{
    "0608/astro-ph0608353_arXiv.txt": {
        "abstract": "We determine what aspects of the density field surrounding galaxies most affect their properties.  For Sloan Digital Sky Survey galaxies, we measure the group environment, meaning the host group luminosity and the distance from the group center (hereafter, ``groupocentric distance''). For comparison, we measure the surrounding density field on scales ranging from 100 $h^{-1}$ kpc to 10 $h^{-1}$ Mpc.  We use the relationship between color and group environment to test the null hypothesis that only the group environment matters, searching for a residual dependence of properties on the surrounding density. Generally, red galaxies are slightly more clustered on small scales ($\\sim$ 100--300 $h^{-1}$ kpc) than the null hypothesis predicts, possibly indicating that substructure within groups has some importance. At large scales ($> 1$ $h^{-1}$ Mpc), the actual projected correlation functions of galaxies are biased at less than the 5\\% level with respect to the null hypothesis predictions.  We exclude strongly the converse null hypothesis, that only the surrounding density (on any scale) matters.  These results generally encourage the use of the halo model description of galaxy bias, which models the galaxy distribution as a function of host halo mass alone.  We compare these results to proposed galaxy formation scenarios within the Cold Dark Matter cosmological model. ",
        "introduction": "\\label{intro}  Galaxy properties are a strong function of their environment. In particular, galaxies in dense environments are older, redder, more concentrated, higher in surface brightness, and more luminous than galaxies in voids (\\citealt{hubble36a, oemler74a, davis76a, postman84a, dressler80a, santiago92a, hermit96a, zehavi02a, norberg02a, hogg03b, blanton05b, kauffmann04a, weinmann05a, maller05a, martinez06a}).  Clearly, part of understanding galaxy formation must involve understanding why regions with different initial conditions (a different initial cosmological density field) result in such different galaxy populations.  Quantitative understanding of this variation of clustering with galaxy type has improved in the last twenty years or so with the advent of large extragalactic samples. Among the papers written on this subject, what the authors mean by ``galaxy type'' has varied considerably. For some ``type'' has meant a measure of its optical morphological properties --- is it an elliptical, a dwarf elliptical, a lenticular, a spiral, or an irregular? Others have preferred measures of galaxy structure such as size or concentration, or measures of the galaxy star-formation history, such as color or emission line flux.  A number of recent papers (\\citealt{kauffmann04a, blanton05b, quintero06a}) have shown that star-formation related properties such as color and emission line flux are directly affected by environment, while structural properties such as surface brightness and concentration (which are more closely related to classical morphology) are not.  In addition, \\citet{kauffmann04a} have shown that environment does not affect measures of very recent star-formation such as emission lines (the last 10 Myr) over and above its affect on longer time scale measures such as optical color (the last 1 Gyr).  Whether other classical morphological measurements, such as spiral arm properties, are related to environment directly, remains to be seen. In this paper, we will simplify based on these results to only consider the relationship between environment and galaxy colors.  Astronomers have not only differed by what they mean by ``galaxy type,'' they also have differed in how they measure ``environment.'' Usually their choices have been motivated by practicality --- after all, the ultimate measure of environment, the local mass density field, is currently observationally inaccessible. In some cases, astronomers have used distance from the center of a cluster or surface density within that cluster as a measure of environment (\\citealt{oemler74a, dressler80a, postman84a, gomez03a, quintero06a}). In other cases, we have used field samples and measured the number of nearby galaxies relative to the mean density, using various measures (\\citealt{hashimoto98a, hashimoto99a, balogh04a, hogg04a, kauffmann04a, blanton05b}). Sometimes, we do not measure the density around individual objects, but instead measure the average environments of classes of galaxies. For example, we can measure as a function of galaxy properties the mean environments around galaxies (\\citealt{hogg03b,blanton05b}) or their correlation functions (\\citealt{davis76a, santiago92a, hermit96a, norberg02a, zehavi04a}). In this paper, we will consider all of these approaches to understand what description of the density field is most directly related to galaxy colors.  All of these analyses have found the same qualitative results. The ``later type'' galaxies --- diffuse, low surface brightness, blue spiral or irregular galaxies --- populate the low density regions. The ``earlier type'' galaxies --- concentrated, high surface brightness, red elliptical or S0 galaxies --- populate the high density regions. However, all of the measures of environment are correlated at least weakly to each other, so one expects them all to show the same qualitative trends.  If we can determine what aspects of environment are {\\it most directly} related to galaxy properties, it will be an important clue in determining the way in which environment affects galaxy formation. Making this determination is the goal of this paper.  We choose here two classes of environmental measurements: ``group environment'' measurements and ``surrounding density measurements'' on varying scales.  Here, we repeatedly use the term ``group environment'' to refer to two parameters regarding the galaxy environment: the luminosity of the host group or cluster (group luminosity) and the radius from the center of that group (groupocentric distance). By ``surrounding density'' we mean the density with respect to the mean in a cylinder around each galaxy, where the cylinder is aligned in the redshift direction to integrate over nonlinear redshift space distortions. The radius of the cylinder in the transverse direction determines the scale. We explain these density measurements in more detail below.  These results yield insights into how the spatial distribution of galaxies is related to that of the dark matter.  In particular, a recently developed way of describing the relationship between galaxies and mass is by using the ``halo occupation distribution'' (HOD), which quantifies the distribution of galaxies as a function of host dark matter halo mass (see \\citealt{berlind02a} and references therein). The HOD description is an extremely convenient way of parameterizing the physics that relates galaxies and mass, and of marginalizing over the possible relationships when trying to use galaxy clustering to constrain cosmology (\\citealt{abazajian05a}).  It typically assumes that the distribution of galaxy luminosities and types depends only on the mass of the host halo, and not on the larger scale density field. We test this assumption here by associating galaxy groups in the observations with dark matter halos, and asking whether indeed the masses of these groups are the only quantities that are relevant to galaxy properties.  In Section \\ref{sdss} we describe the Sloan Digital Sky Survey (SDSS; \\citealt{york00a}) and the \\citet{berlind06a} group catalog resulting from it.  In Section \\ref{environs}, we find that both group luminosity and groupocentric distance are independently related to galaxy properties. We then consider whether the mean density around galaxies is related to galaxy colors over and above the dependence expected to result just from the dependence on group environment. We only find a large residual dependence on very small scales ($\\sim 300$ $h^{-1}$ kpc). Meanwhile, even in regions with very different densities on large scales, the galaxy population is similar as long as the group environment is similar. In Section \\ref{correlation} we demonstrate the same results in terms of the correlation functions of red and blue galaxies.  In Section \\ref{others} we compare our results to similar studies of others, finding agreement. In Section \\ref{theory}, we compare our results to similar investigations in the theoretical realm. In Section \\ref{conclusions} we discuss the implications of our results for theories of galaxy formation and for interpreting large-scale structure statistics. ",
        "conclusions": "\\label{conclusions}  We have described the relationship between galaxy colors and their group environments, in terms of the luminosity of their host group and their distance from the center of that group (groupocentric distance). Furthermore, we have searched for any residual relationship between galaxy color and measures of the surrounding environment on various scales, finding only a slight dependence at large scales ($> 1$ $h^{-1}$ Mpc) and a stronger dependence at the smallest scales ($<300$ $h^{-1}$ kpc). Measured at any scale, the variation of galaxy colors with the surrounding density field does not explain the variation of color with group environment: the group environment always yields extra information.  Since other properties, such as concentration and surface brightness, do not appear to correlate with the density field independently of color (\\citealt{blanton05b,kauffmann04a, quintero06a}) these results likely extend to other galaxy properties as well. In fact, we have checked this proposition explicitly for concentration and found no dependence of concentration on the larger scale density field once the color and group environment is fixed.  When interpreting these results, keep in mind that we have considered for the most part only galaxies with $M_r - 5\\log_{10} h < -19$ (though in Figure \\ref{bias} we also look at slightly lower luminosity galaxies), which limits us to dark matter halos of $\\sim 5\\times 10^{11}$ $h^{-1}$ $M_\\odot$ or greater. Whether the assembly time or large-scale density affects galaxies in smaller halos is an open question.  These results have important consequences for the study of galaxy formation: \\begin{enumerate} \\item That the blue fraction has no residual relationship with the large scale density field (such as 1 $h^{-1}$ Mpc and above) demonstrates that galaxy formation is closely tied to physics within each halo. The location of that halo in the large-scale density field appears not to be important, nor (if CDM is correct) the assembly time of the host halos (at fixed halo mass). \\item The residual dependence of color on small scale density may indicate the importance of surviving substructure within the group (though it is hard to disentangle possible projection effects causing this dependence). If so, it may be a signature of the processing of blue galaxies into red galaxies in moderately sized groups prior to infall into large groups. \\end{enumerate}  These results are also important for studies of large-scale structure that depend on the use of the halo occupation distribution formalism to model small-scale clustering. A typical simplification (not a necessary one) of those models is that how galaxies occupy halos depends only on their mass, not their larger scale environment. By showing that the {\\it relative} distribution of different types of galaxies is not affected by the larger scale environment (while it manifestly is affected by the group luminosity) we lend credence to this assumption. Furthermore, theoretically speaking, the large-scale density field is related to the assembly time of the halos, and we might expect assembly times to be related to the ages of the galaxies. Either the halos containing our $L_\\ast$ galaxy sample do not have much relationship between their assembly time and large-scale clustering, or the process of galaxy formation is not much affected by the assembly time of the halo at a given halo mass. Either possibility is encouraging to those trying to use the halo model to interpret the medium-scale correlation functions of similar samples in terms of cosmological models."
    },
    "0608/astro-ph0608165_arXiv.txt": {
        "abstract": "The $w-w'$ plane, defined by the equation of state parameter for the dark energy and its derivative with respect to the logarithm of the scale factor, is useful to the study of classifying the dynamical dark energy models. In this note, we examine the evolving behavior of the two-field quintom models with $w$ crossing the $w=-1$ barrier in the $w-w'$ plane. We find that these models can be divided into two categories, type A quintom in which $w$ changes from $>-1$ to $<-1$ and type B quintom in which $w$ changes from $<-1$ to $>-1$ as the universe expands. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608215_arXiv.txt": {
        "abstract": "We report on the spectroscopic monitoring of GCIRS16SW, an Ofpe/WN9 star and LBV candidate in the central parsec of the Galaxy. SINFONI observations show strong daily spectroscopic changes in the K band. Radial velocities are derived from the He~{\\sc i} $2.112 \\mu m$ line complex and vary regularly with a period of 19.45 days, indicating that the star is most likely an eclipsing binary. Under various assumptions, we are able to derive a mass of $\\sim$ 50 \\msun\\ for each component. ",
        "introduction": "\\label{intro}  The central cluster constitutes one of the largest concentrations of massive stars in the Galaxy \\citep{genzel03}. Nearly 100 OB and Wolf-Rayet stars are confined in a compact region of radius $\\sim$ 0.5 parsec centered on the super-massive black hole associated with the radio source SgrA* \\citep{paumard06}. Among this population of young massive stars, six are thought to be Luminous Blue Variables (LBV): IRS16NE, IRS16C, IRS16NW, IRS16SW, IRS33E and IRS34W \\citep{paumard04,trippe06}. LBVs are evolved massive stars experiencing strong variability in both photometry and spectroscopy due to their proximity to the Humphreys-Davidson limit \\citep{hd94}, a region of the HR diagram where the luminosity of the stars reaches the Eddington luminosity so that instabilities develop in their atmospheres, leading to strong mass ejection and drastic changes in the stellar properties (\\teff, radius). The six stars mentioned above are only LBV ``candidates'' (LBVc) since they have not been observed to experience the strong outbursts and photometric changes typical of bona fide LBVs such as $\\eta$ Car \\citep{dh97}. However, their luminosities and spectra are very similar to stars known to be ``quiescent'' LBV, i.e. stars having experienced an LBV event in the past and being now in a more stable phase. In addition, one of them - IRS34W - has shown photometric variability on timescales of months to years which was interpreted as the formation of dust from material previously ejected by an LBV outburst \\citep{trippe06}.  Among these six stars, IRS16SW deserves special attention. This star was claimed to be a massive eclipsing binary by \\citet{ott99} since its K band magnitude displays regular variations with a periodicity of 9.72 days. However, the absence of a second eclipse in the light-curve lead \\citet{depoy04} to the conclusion that the binary scenario was not correct, and that IRS16SW was instead a pulsating massive star, a class of star predicted by theory but not observed so far.  Here, we present results of the spectroscopic monitoring of IRS16SW revealing periodic variations of radial velocities which are interpreted as the signature of a massive spectroscopic and eclipsing binary .   \\begin{figure*} \\epsscale{0.9} \\begin{center} \\includegraphics[width=8.5cm,height=6.5cm]{f1a.eps} \\includegraphics[width=8.5cm,height=6.5cm]{f1b.eps} \\end{center} \\caption{Montage of spectra of IRS16SW around the region of He~{\\sc i} 2.112 $\\mu m$ (left) and \\brg\\ (right) taken between October 2$^{nd}$ and October 12$^{th}$ 2005. The changes in both the line shape and position are clearly seen. \\label{var_lines}} \\end{figure*} ",
        "conclusions": "\\label{disc}   \\subsection{Binary versus pulsating variable} \\label{disc_bin}  \\citet{depoy04} argue that IRS16SW was a pulsating massive star based on the absence of second minimum in the light-curve and the difficulty to fit this light-curve in the binary scenario (but see Sect.\\ \\ref{orb_sol}). They compare the observed variation in K magnitude to the predictions of the dynamical models of \\citet{dg00} and conclude that there is a reasonable qualitative agreement. However, there are some quantitative discrepancies. First, the amplitude of the variation is much larger than predicted: although \\citet{dg00} do not compute K band photometry, one can estimate the variation in this band to be at most 0.2 mag (inspection of their Table 2 reveals that the amplitude of photometric variations decreases with wavelength and is $\\lesssim 0.2$ mag in the I band), while we observe 0.55 mag.  Second, the period we derive $-$ 19.45 days $-$ is larger than expected in the pulsating scenario \\citep[see Table 1 of][]{dg00}.  Concerning spectroscopic changes, although in principle one can not completely rule out the possibility that they are due to motions of the atmosphere and fluctuations of the physical parameters (\\teff, radius) due to pulsations \\citep{dg00}, the timescales are again not consistent: $\\sim$ 1 day for pulsations compared to 19.45 days observed. Besides, so far there are no theoretical predictions of spectroscopic changes caused by pulsations to which we could compare our observed spectra. The binary nature of IRS16SW is thus strongly favored.  The absence of secondary eclipse in the light-curve is explained by the similar K band magnitude of the two components. This is another indication that both stars are very similar. \\citet{depoy04} report the presence of a variation in $H-K$ on a period of 9.725 days, $H-K$ being bluer when the system is brighter. A similar trend was observed in the optical photometry of the massive contact binary V606 Centauri \\citep{lorenz99}. We interpret this as a sign of heating in the contact zone, making the spectral energy distribution in this region bluer. Massive binaries are indeed known to produce X-rays through colliding winds, and we might expect the same kind of interaction and heating around the contact region. Again, since this region is seen twice during an orbital revolution, an observed period half the true orbital period is naturally derived from the $H-K$ curve.    \\subsection{Stellar evolution and the LBV phenomenon} \\label{disc_evol}   Whether or not all massive stars go through the LBV phase is still under debate. \\citet{langer94} and \\citet{pasquali97} argue that this is the case, while other observational \\citep{crowther95} and theoretical \\citep{mm05} studies indicate that the most massive stars ($M \\gtrsim 60 \\msun$) may skip this phase. This is an important issue since although short, the LBV phase is crucial in the mass loss history, and consequently in the subsequent evolution, of massive stars.  Recent studies by \\citet{smith06} even claim that most of the mass of hot stars is lost during the LBV phase.  Here, we provide an accurate measurement of the present mass of a candidate LBV, confirming that a star with an initial mass larger than 50 \\msun\\ (and likely of the order 60-70 \\msun) may become a LBV. This is a strong constraint for evolutionary models.  Of course, one could argue that the star's evolution was affected by binarity. However, inspection of the spectral morphology and physical properties \\citep[see][]{paco97} of IRS16SW and the other LBVc in the Galactic Center shows similarities. Our monitoring of IRS16SW also includes the other LBVc in the Galactic Center. Except for IRS16NE, none of these stars showed any spectroscopic variation, ruling out the possibility that they are close binaries. IRS16NE showed some RV fluctuations \\citep[see also][]{tanner06}, but so far we do not have enough data point to sample an hypothetical RV curve. Hence, the similarity between the spectrum of IRS16SW and those of the other single LBV candidates leads us to the conclusion that binarity has not (yet?) significantly affected the evolution of IRS16SW. Since the mass of the two components is similar, one can speculate that IRS16SW is composed of two stars initially equally massive that have so far evolved in parallel in a detached system, without influencing each other. They may have just entered the LBV phase during which contact was achieved due to their respective expansion. This event probably happened very recently (the LBV phase lasting $\\sim 10^{5} yr$) so that the general properties of both components have not yet been affected by mass transfer and binary evolution. Such a scenario is consistent with both stars displaying similar spectra (but see Sect.\\ \\ref{spec_disentang} for caution words).   Assuming that IRS16SW was not affected by binary evolution, the properties of the LBV candidate in IRS16SW can thus be used to constrain evolutionary models of single massive stars.  In their recent models including rotation, \\citet{mm05} stressed that the LBV phase is not systematically reached above 45 \\msun. Here, we have an example for which it is the case (under the assumption that IRS16SW will turn $-$ or has already turned in the past $-$ into a genuine LBV)."
    },
    "0608/astro-ph0608509_arXiv.txt": {
        "abstract": "The gravitational collapse of a non-rotating, black-hole-forming massive star is studied by $\\nu$-radiation-hydrodynamical simulations for two different sets of realistic equation of state of dense matter. We show that the event will produce as many neutrinos as the ordinary supernova, but with distinctive characteristics in luminosities and spectra that will be an unmistakable indication of black hole formation. More importantly, the neutrino signals are quite sensitive to the difference of equation of state and can be used as a useful probe into the properties of dense matter. The event will be unique in that they will be shining only by neutrinos (and, possibly, gravitational waves) but not by photons, and hence they should be an important target of neutrino astronomy. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608023_arXiv.txt": {
        "abstract": "{ We present a detailed study of the single radio pulses of PSR B0656+14, a pulsar also known to be a strong pulsed source of high-energy emission. } { Despite the extensive studies at high-energy wavelengths, there is little or no published work on its single-pulse behaviour in the radio band. In this report we rectify this omission. } { Radio observations using the Westerbork Synthesis Radio Telescope at 1380 MHz and the Arecibo Observatory at 327 and 1525 MHz are used to investigate the single-pulse behavior of PSR B0656+14. A comparison is made with the phenomena of giant pulses and giant micropulses. } { We have found that the shape of the pulse profile of PSR B0656+14 requires\\rm an unusually long timescale to achieve stability (over 25,000 pulses at 327 MHz). This instability is caused by very bright and narrow pulses with widths and luminosities comparable to those observed for the RRATs. Many pulses are bright enough to qualify as ``giant pulses'', but are broader than those usually meant by this term.  At 327 MHz the brightest pulse was about 116 times brighter than the average pulse.  Although the most powerful pulses peak near the centre of the profile, occasional sudden strong pulses are also found on the extreme leading edge of the profile. One of them has a peak flux of about 2000 times the average flux at that pulse longitude. No ``break'' in the pulse-energy distributions is observed, but nevertheless there is evidence of two separate populations of pulses: bright pulses have a narrow ``spiky'' appearance consisting of short quasi-periodic bursts of emission with microstructure, in contrast to the underlying weaker broad pulses. Furthermore, the spiky pulses tend to appear in clusters which arise and dissipate over about 10 periods. We demonstrate that the spiky emission builds a narrow and peaked profile, whereas the weak emission produces a broad hump, which is largely responsible for the shoulders in the total emission profiles at both high and low frequencies.  } {} ",
        "introduction": "\\object{PSR B0656+14}, at a spin-down age of 111,000 yrs, is one of three nearby pulsars in the middle-age range in which pulsed high-energy emission has been detected. These are commonly known as ``The Three Musketeers'' (\\citealt{bt97}), the other two being Geminga and PSR B1055--52. Middle-aged pulsars (roughly defined to be those whose spin-down ages range from 50,000 yrs to 300,000 yrs) are of interest to neutron star theorists because they allow the detection of high-energy thermal radiation from the star's surface, and hence can provide tests for models of surface cooling and atmosphere composition. In more energetic young pulsars (the Crab) the pulsed non-thermal emission is dominant, and completely masks the thermal emission at all wavelengths. In cool, older pulsars blackbody surface temperatures are expected to have fallen below the detection threshold, with the likely exception of their heated polar caps (the polar cap of one such older pulsar, B0943+10, having recently been detected in soft X-rays (\\citealt{zsp05}).  However pulsars such as B0656+14 are also of interest because they provide tests between competing emission models for the highest energy components. Their hard X-ray emission can variously be interpreted as a product of outergaps located in the outer magnetosphere (\\citealt{cz99}), or as originating closer to the surface magnetic poles (\\citealt{hm98}). Fundamental to these discussions are the determination of the geometric location of the various high-energy peaks, and their relation to the radio emission peak, which is generally taken to define the polar cap of the pulsar's dipole axis. In the case of PSR B0656+14, a major clue is that observations from optical to hard X-ray all show one or both of the two peaks always at the same two phase longitudes (\\citealt{dcm+05,kmmh03,pzs02,ssl+05,zps96}), while the radio peak falls almost, but not quite, halfway between these peaks.  PSR B0656+14 was included in a recent extensive survey of subpulse modulation in pulsars in the northern sky (\\citealt{wes06}). In the single pulses analysed for this purpose the unusual nature of this pulsar's emission was very evident, especially the ``spiky'' nature of the subpulses. But what was most striking was that the pulsar occasionally exhibited exceptionally powerful and longitudinally narrow subpulses reminiscent of ``giant'' pulses, hitherto reported for only a handful of pulsars, including the Crab pulsar (\\citealt{sr68}) and mostly young or millisecond pulsars (e.g. \\citealt{spb+04}). More recently, giant pulses have been discovered in two old pulsars (\\citealt{ke04}). If their presence could be confirmed in PSR B0656+14, this would demonstrate that the phenomenon might also be found in pulsars of intermediate age and hence at any stage of a pulsar's lifetime.  We therefore set out to explore the full nature of PSR B0656+14's pulse behaviour in the radio band. PSR B0656+14 has a period (0.385 sec) not greatly below the average for all pulsars and at radio frequencies exhibits an apparently unremarkable single-peak integrated profile. Yet at any wavelength the integrated profile of a pulsar conceals as much information as it yields, and this pulsar has proved to be no exception. It is remarkable that, despite the extensive studies at high-energy wavelengths, there is little or no published work on its single-pulse behaviour in the radio band. In this report we rectify this omission, and, in addition to confirming the presence of extremely bright pulses across the profile, demonstrate that the pulsar's radio emission is far from that typical of older better-known pulsars.  While this work was being prepared the discovery of RRATs (Rotating RAdio Transients) was announced. These are sources which emit single powerful pulses, separated by long intervals, and were identified as isolated pulsars with periods of between 0.4 and 7 seconds. However, their relation to the known pulsar population was unclear.  The intermittent giant pulses we have detected in PSR B0656+14 led us to argue in a separate paper that this pulsar, were it not so near, could itself have appeared as an RRAT (\\citealt{wsr+06}). This paper is complementary to that work, showing that this pulsar's bright pulses are narrow and ``spiky'', occur at a wide range of central longitudes and can be differentiated from a weaker but steadier underlying emission. A third paper is planned in which we will present polarization data and attempt to link the radio emission -- in a single geometric structure -- to the high-energy peaks.    The details of the radio observations are described in the next section. In Sect. \\ref{SctProfile} the pulse profile is discussed, in Sect. \\ref{SctSinglePulse} the properties of the single pulses and in Sect. \\ref{SctSeparation5} we will show that the emission can be decomposed into ``spiky'' and ``weak'' emission. In Sect. \\ref{SctDiscussion} our results are discussed and summarized. ",
        "conclusions": "Discussion and conclusions}  \\subsection{Comparison with the known population of pulsars}  We have shown that PSR B0656+14 intermittently emits pulses that are extremely bright compared to normal pulsars and with pulse energies well above 10 {\\Eav} these pulses formally qualify as giant pulses. Nevertheless these pulses differ from giant pulses and giant micropulses in important ways:  \\begin{itemize} \\item The pulse-energy distribution of PSR B0656+14 does not show a power-law tail (this is possibly also true for PSR B1133+16; \\citealt{kkg+03}) \\item The widths of the bright pulses of PSR B0656+14 are much greater than those of the extremely narrow giant pulses (e.g. \\citealt{spb+04}), although similar to that of the giant micropulses. \\item The brightest pulses of PSR B0656+14 are an order of magnitude brighter (both in integrated energy and in peak flux) than the giant micropulses of the Vela pulsar (\\citealt{jvkb01}) and PSR B1706--44 (\\citealt{jr02}). \\item The bright pulses of PSR B0656+14 are not strongly constrained in pulse longitude. In most other pulsars, both fast (e.g. B1937+21, B1821--24) and slow (e.g.  B0031--07) the pulse longitudes where giant pulses appear are highly constrained, whether towards one edge or centrally (e.g \\citealt{cst+96,kt00,rj01,ke04}). Also the giant micropulses detected for the Vela pulsar and the young pulsar B1706--44 are highly confined (in both cases to the leading edge of the pulse profile). \\item The claimed empirical correlation between the appearance of giant pulses and high light-cylinder magnetic field strengths clearly fails for PSR B0656+14 since its value (770 Gauss) is well below those of most pulsars exhibiting giant pulses (around $10^5$ Gauss; \\citealt{cst+96}). However, PSR B0031--07 and a number of other slow pulsars also easily fail this test, so the correlation may only be valid for millisecond pulsars. \\item The bright pulses of B0656+14 do not coincide with any of the observed X-ray peaks, in contrast to those of the well-known pulsars exhibiting giant pulses. \\end{itemize}  Despite the differences between the spiky emission of PSR B0656+14 and the giant micropulses listed above, there are also interesting parallels between PSR B0656+14 and Vela. The pulse profile of PSR B0656+14 has a shoulder on the trailing edge reminiscent of a similar feature discovered in Vela (\\citealt{jvkb01}) and B1937+21 (\\citealt{cst+96}). In both, the shoulder is not associated with the giant pulses (\\citealt{kt00}), and that also appears to be the case in PSR B0656+14. Furthermore both pulsars have low-frequency intensity modulations and neither pulsar shows a coincidence between the locations of the radio and high-energy emission.   Many of the exceptional properties of PSR B0656+14 have led us to point out (\\citealt{wsr+06}) that this pulsar, were it not so near, could have been discovered as an RRAT. Chief among these are the infrequent, but luminous pulses with very high peak fluxes and widths comparable with those of the RRATs. Furthermore, like RRATs, PSR B0656+14 has a light-cylinder magnetic field strength which is not exceptional.  Surprisingly, the middle-aged PSR B0656+14 also shares many properties with recycled millisecond pulsars. For instance both have wide profiles that have a more or less constant width or component separation over a very wide frequency range (e.g. \\citealt{kll+99})\\footnote{However, it should be noted that invariant profile widths is not a property confined to millisecond pulsars. \\citep{mr02a} have shown that a whole class of bright slow pulsars have component widths and component separations independent of frequency.}. Furthermore the emission of millisecond pulsars is usually highly polarized with a flat linear polarization profile (e.g. \\citealt{xkj+98}). PSR B0656+14 is 92\\% linear polarized at 400 MHz and also the position angle is remarkably flat (\\citealt{gl98}). An interesting comparison might be drawn between the B0656+14 and a number of millisecond pulsars which are also known to have slowly-varying unstable profiles (e.g. \\citealt{rk03,kxc+99,bs97}), possibly resulting from spiky emission in one or more components.   \\subsection{The spiky emission}  We show in this paper that the emission of PSR B0656+14 can be characterized by spiky emission that is not confined to a narrow pulse-longitude range, but generally concentrated toward the centre of its profile.  This spiky emission has a low occurrence rate within each pulse and is often quasi-periodic in structure. In addition, pulses with spiky emission have a tendency to cluster in successive pulses. By contrast, there is also a weak emission, which appears with a high occurrence rate over the full width of the pulse and varies little from pulse to pulse. Furthermore, the profiles of the spiky and weak emission have very different shapes and frequency evolution. Taken together, these properties all argue for two distinct types of emission.   Separate properties of weak and strong emission can be supported by the fits of the pulse-energy distribution. However, in Fig. \\ref{lred_fit} one can see that all pulses with a peak-flux weaker than the average are below the noise level. Therefore more sensitive observations (with at least an order of magnitude increase in $S/N$) are required to see single weak pulses and to find out if there exists a break in the pulse-energy distribution.  Such an observation would show if the spiky and weak emission are associated with the extreme ends of a single smooth energy distribution, or if they are two physically distinct distributions.   Pulse profiles of many pulsars can accurately be decomposed into a small number of Gaussian emission components (\\citealt{kwj+94}).  This is not possible for the isosceles shape of the profile of the spiky emission at high frequencies. This, together with the width of the central component that is invariant with frequency, may suggest that it reflects a permanent geometrical feature. It might be related to an intensity-enhancing (e.g. caustic) effect that boosts the otherwise weak emission.  Due to the low occurence rate of the spiky emission, the profile requires an unusually long timescale to achieve stability (over 25,000 pulses at 327 MHz). Despite the high variability in shape of the spiky emission profile on short timescales, we found that the radio power output remains remarkably stable and this may indicate that processes ``repackage'' the power within the magnetosphere, rather than generate it.     The emission of PSR B0656+14 shows a quasi-periodicity of $20P_1$, which is intensity modulated, rather than phase modulated. Similar fluctuations without drift are often associated with core emission within a conal structure (e.g. PSR B1237+25; \\citealt{sr05}). To the best of our knowledge, there have been no theoretical attempts to account for the periodicity of this emission. Possibly the modulation should be considered in the framework of the charge flow balance within the entire magnetosphere. Although we conventionally suppose in such a pulsar that we only have a view of one pole, it has recently been plausibly suggested (\\citealt{dfs+05}) that in selected stars (all of them pulsed X-ray sources, as is PSR B0656+14) it is possible to clearly detect radiation from particles downflowing towards the pole that is hidden from our view.  The widely-separated high-energy peaks located well away from the phase of the polar cap point to activity in the outer magnetosphere, possibly associated with an outer gap. Also in the case of PSR B0656+14 we may be dealing with emission which is stimulated, either directly on indirectly, by an inflow - as well as an outflow - of charged particles (an analogy with rainfall is not inappropriate), with short-lived showers of particles intermittently injected into the polar regions in a quasi-periodic manner. This idea will be explored in a subsequent paper, in which we will also present polarization data and attempt to link the radio emission to the high-energy peaks."
    },
    "0608/astro-ph0608090_arXiv.txt": {
        "abstract": "We consider dark matter annihilation into Standard Model particles and show that the least detectable final states, namely neutrinos, define an upper bound on the total cross section.  Calculating the cosmic diffuse neutrino signal, and comparing it to the measured terrestrial atmospheric neutrino background, we derive a strong and general bound.  This can be evaded if the annihilation products are dominantly new and truly invisible particles. Our bound is much stronger than the unitarity bound at the most interesting masses, shows that dark matter halos cannot be significantly modified by annihilations, and can be improved by a factor of 10--100 with existing neutrino experiments. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608573_arXiv.txt": {
        "abstract": "We present the first far-red spectra of the L subdwarf 2MASS J16262034+3925190 and the late M/early L subdwarf 2MASS J16403197+1231068.  We confirm the ultracool subdwarf nature of these objects.  2M1626+3925 shows strong K~I absorption, like an L4 dwarf,  and is very similar to, but hotter than, the late L subdwarf 2MASS J05325346+8246465.  It is unambiguously an L subdwarf.  2M1640+1231 is very similar to SSSPM J1444-2019, which has been classified as sdM9 or early sdL.   In contrast to the hotter M subdwarfs, L subdwarfs are  characterized by not only enhanced hydrides but also strong TiO.  Progress in a classification system requires identification of more subdwarfs to map out their variations. ",
        "introduction": "Although most stars in the solar neighborhood are members of the Galactic disk, there is a trace population of old, low-metallicity stars.   The spectra of near-solar metallicity low-mass stars can be classified in the M dwarf scheme \\citep{khm} and are dominated in the red by molecular absorption, most notably TiO.   The spectra of metal-poor low-mass stars are dramatically different, and can be classified in the scheme of \\citet{g97}[hereafter G97], who divided the stars into M dwarfs (M), subdwarfs (sdM) and extreme subdwarfs (esdM), on the basis of the TiO and CaH features.  The sdM have reduced but significant TiO absorption compared to M dwarfs, while extreme M subdwarfs have weak or even absent TiO.  At and below the hydrogen-burning limit, solar metallicity dwarf spectra begin to look qualitatively different, and are described by the L dwarf classification scheme \\citep{k99}.  TiO, and later VO, weaken and disappear as dust grains form; the neutral alkali (Na, K, Rb, Cs) lines are strong; and the hydrides strengthen (see Kirkpatrick 2006 for a complete review and references.)  Although the G97 system covers the LHS stars \\citep{l79,faintlhs2}, new surveys have begun to identify even cooler subdwarfs  and have extended the sdM and esdM sequences \\citep{esdm7,sdm8,esdm65, sdm95,esdm8} to later spectral types.   Four subdwarfs have merited suggestions that they should be considered L subdwarfs by their discoverers.   \\citet{latesdl} found 2MASS J05325346+8246465\\footnote{Hereafter 2M0532+8246; we follow this convention for abbreviation for other object designations.} as part of a T dwarf search ---  it is dramatically different in both the red and near-IR from any L dwarf and is most likely a metal-poor brown dwarf.  \\citet{earlysdl}  discovered LSR1610-0040 as part of a photographic plate-based proper motion survey --- it resembles an sdM but has anomalously strong Rb~I and Cs~I lines, leading to the suggestions that it should be considered an early L dwarf.  However, both \\citet{c1610} and \\citet{b1610} have concluded it is peculiar mid-M subdwarf with unusual abundances.  Most recently, \\citet{sss1444} report \"the formal classification of SSSPM J1444-2019 is sdM9, although it appears likely that it is, in fact, an L-type subdwarf.\"   \\citet{latesdl2} reports 2MASS J16262034+3925190 as an L subdwarf that is not as cool as 2M0532+8246  We present red spectra of three cool subdwarfs:  LSR1610-0040, 2M1626+3925, and 2MASS J16403197+1231068.  Like 2M1626+3925, 2M1640+1231\\citep{bother} was discovered in a 2MASS-based T dwarf search.  It  is classified as \"sdM8?\" on the basis of a near-infrared spectrum, but existing \"sdM9\" type objects have been classified in the red.  We present the data in Section~2 and discuss issues in the classification and interpretation of these objects in Section~3. ",
        "conclusions": "2M1640+1231 can be understood as an subdwarf analog to an M8/M9 dwarf; the current optical classification suggests a type even later than the sdM8 IR classification -- either sdM9 or early sdL.  It is very similar to SSSPM J1444-2019.  2M1626+3925 can be understood as an L subdwarf; it is similar to, but 'earlier' (hotter) than,  2M0532+8246.   LSR1610-0040 does not appear to be in between these two objects; however, the features pointed out by \\citet{earlysdl} are confirmed.   Both \\citet{b1610} and \\citet{c1610} have extensive discussions of this peculiar object and its classification.  In contrast to the situation in the range sdM0-sdM7, the latest M and L subdwarfs show {\\it enhanced} TiO absorption relative to late-M and L dwarf spectra.  Physically, this is due to dust being relatively less important than in solar-metallicity dwarfs, as suggested by \\citet{latesdl}.  There are two important points that emerge from the difficulty in dealing with the spectra of these stars. First,  if the \"L subdwarf\" term is to suggest features in common with the L dwarf sequence, both 2M1640+1231 and LSR1610-0040, as well as SSSPM J1013-1356 and SSSPM J1444-2019, should be considered M subdwarfs.  However, as noted by \\citet{earlysdl} and \\citet{sss1444}, we have run out of single digit sdM types, especially if SSSPM J1013-1356 is an sdM9.5, pointing to an early L designation. The G97 system has been successful in distinguishing the dwarf, subdwarf and extreme subdwarfs bins, and in its numerical types it distinguishes within these bins. Types were assigned so that CaH absorption was similar for all M, sdM and esdM; for the hotter stars recent work confirms that the CaH is primarily sensitive to temperature \\citep{ww06}, but since the latest types have CaH stronger than in M dwarfs this system necessarily breaks down.  LHS 377 is {\\it by definition} an sdM7 in the G97 system, but it was the only object later than sdM5 known at that time!  At least in the red spectrum here, an sdM9 classification of 2M1640+1231 -- the same as SSSPM J1444-2019 -- suggests that existing numerical types are not wildly inappropriate.  It is therefore possible that the community may prefer to describe such objects as SSSPM J1444-2019 and 2M1640+1231 as sdM9, and therefore compress the existing late sdMs that have extended the sdM sequence into the sdM6-sdM8 subclasses. The community could alternatively allow the sdM spectral classes to extend beyond sdM9; e.g., sdM10, sdM11, etc., if it is desirable to avoid sdL types. The behavior of TiO and CaH suggests that in subdwarfs that the transition from M to L is rather different at low metallicity, so a one-to-one matching may not be a good description of the spectra anyway: 2M1626+3925 is dominated by the same CaH and TiO features as sdMs in the limited range 6500-7300\\AA.  Furthermore, the 2MASS L subdwarfs are {\\it blue} in  $J-K_s$, yet L dwarfs are much {\\it redder} in $J-K_s$ than M dwarfs.     The change in the overall appearance once enough subdwarfs have been discovered should govern the M/L transition, and is likely to be better linked to physical changes.    It may well be, that happened with T dwarfs, that near-infrared classification is more useful:  see \\citet{esdm8} for just such a suggestion. The best solution is to wait until more late sdM/early sdL's have been discovered and the relative merits of red and near-IR typing can be assessed.  Second, the fact that the three objects presented in this paper seem to be an extension of the sdM sequence, rather than the esdM sequence, yet do not clearly fit into a sequence strongly suggests that we need more objects simply to characterize the appearance of the spectra near the M/L transition and that an extra parameter may be needed to describe them.  Such factors as unresolved binaries, changes in $[\\alpha/Fe]$ abundances,  and higher sensitivity to $[m/H]$ might all play a role.  The fact that an object,   LSR1610-0040, that is so difficult to fit into the existing framework \\citep{earlysdl,c1610,b1610}  has already been found should serve as a cautionary warning.  \\citet{sss1449}'s suggestion that the latest sdM and sdL's are only $[m/H]=-0.5$ is also consistent with the idea that the coolest dwarfs are very sensitive to metallicity; on the other hand, it seems to us, following G97, that if $[m/H]=-0.5$ caused subdwarf classification then there would be large numbers of the them already discovered in the disk, rather than only a handful of objects with halo kinematics.  Rather, we believe these objects are likely to have $[m/H] <-0.5$ (i.e., $[m/H] \\approx -1$).  Detailed studies of the objects discussed in this paper, as well as such objects as  the blue M7 LHS 1317 \\citep{faintlhs2}, are needed to answer these questions.  In any case, we suggest that all numerical spectral subtypes of sdM7 and later, be considered tentative until a large sample with both optical and near-IR spectra can be characterized."
    },
    "0608/astro-ph0608059_arXiv.txt": {
        "abstract": "We consider particles with low free or proper eccentricity that are orbiting near planets on eccentric orbits. Via collisionless particle integration we numerically find the location of the boundary of the chaotic zone in the planet's corotation region. We find that the distance in semi-major axis between the planet and boundary depends on the planet mass to the 2/7 power and is independent of the planet eccentricity, at least for planet eccentricities below 0.3. Our integrations reveal a similarity between the dynamics of particles at zero eccentricity near a planet in a circular orbit and with zero free eccentricity particles near an eccentric planet. The 2/7 law has been previously explained by estimating the semi-major at which the first order mean motion resonances are large enough to overlap. Orbital dynamics near an eccentric planet could differ due to first order corotation resonances that have strength proportional to the planet's eccentricity. However, we find the corotation resonance width at low free eccentricity is small.  Also the first order resonance width at zero free eccentricity is the same as that for a zero eccentricity particle near a planet in a circular orbit. This accounts for insensitivity of the chaotic zone width to planet eccentricity. Particles at zero free eccentricity near an eccentric planet have similar dynamics to those at  zero eccentricity near a planet in a circular orbit. ",
        "introduction": "Chaotic diffusion associated with the overlap of resonances has been shown to be responsible for instabilities in the solar system (e.g., see \\citealt{holman93,lecar01,lecar02,tsiganis05}). For the restricted 3-body problem \\citet{wisdom80} first showed that the width of the chaotic zone near a planet could be explained by calculating the location at which the first order mean motion resonances are large enough to overlap. The zone width has been measured numerical and predicted theoretically \\citep{wisdom80,duncan89,murray97} for a planet in a circular orbit, though some work has extended the stability analysis for bodies in orbits near circular and eccentric binary stars \\citep{holman99,mudryk06}. The stability of bodies at low eccentricity residing in multiple planet extrasolar systems have also been investigated numerically (e.g., \\citealt{rasio96,lepage04,barnes04}).  Recently \\citet{quillen06b} suggested that the edge of Fomalhaut's eccentric ring could be due to truncation by a 0.1 eccentricity Neptune mass planet. The nearby star Fomalhaut hosts a ring of circumstellar material \\citep{aumann85,gillett85} residing between 120 and 160 AU from the star \\citep{holland98,dent00,holland03}. {\\it Spitzer Space Telescope} infrared observations of Fomalhaut reveal a strong brightness asymmetry in the ring \\citep{stapelfeldt04,marsh05}. Recent {\\it Hubble Space Telescope} ({\\it HST}) observations show that this ring has both a steep and eccentric inner edge \\citep{kalas05}. The sharp disk edge suggested that the dust particles are in orbits with low free or proper eccentricity, thus the ring has eccentricity equal to the forced eccentricity caused by secular perturbations from the proposed planet. Such a configuration  is possible if inelastic collisions in the disk damp the eccentricities of particles, resulting in a particle distribution that moves along nearly closed streamlines or closed and non-self-intersecting orbits. Fomalhaut's ring has a intermediate collision timescale of $10^3$ orbits, estimated from its normal disk opacity $\\tau \\sim 1.6\\mu$m at 24$\\mu$m \\citep{marsh05}.  While previous theoretical and numerical studies have considered orbit stability near  the corotation region for planets on circular orbits, little work has been done considering the stability of orbits near a planet on an eccentric orbit. The dynamical problem of an object orbiting a planet in a circular orbit has a conserved quantity, the Jacobi integral, that is not conserved when the planet is on an eccentric orbit. Surfaces of section have been used to illustrate the types of orbits (tori or area filling) for the restricted three-body system \\citep{wisdom85,winter97}. However, when the planet is eccentric there is no extra integral of motion making it difficult to create surfaces of section. We are motivated here to consider the role of the planet's eccentricity in setting the boundary of non-stochastic orbits in the corotation region.  We focus here on particle orbits that have nearly zero free eccentricity and so have eccentricity set by the forced eccentricity due to secular perturbations from the planet. ",
        "conclusions": "In this paper we have investigated the dynamics of low free eccentricity collisionless massless particles in the plane near a planet  on a eccentric orbit. By determining the semi-major axis at which the particle lifetime increases, we measure the width of the chaotic zone near the planet. For eccentricity $e_p < 0.3$ we find that the chaotic zone width is independent of the planet's eccentricity and matches that predicted by the 2/7 law. The eccentricity dispersion in the disk edge and the lifetime of particles within the chaotic zone is also nearly independent of planet eccentricity.  These results suggest there is a similarity in the dynamics of particles at zero free eccentricity and those at zero eccentricity near a planet in a circular orbit.  To account for our numerical results we have explored  the dynamics of a  Hamiltonian system that takes into account first order secular perturbations and the two terms that comprise each first order mean motion resonance.  With a coordinate transformation we have rewritten the Hamiltonian in terms of an action variable that depends on the free eccentricity rather than the eccentricity.  At low free eccentricity we find that the new Hamiltonian resembles the Hamiltonian of a low eccentricity particle near a planet in a circular orbit. This accounts for the lack of sensitivity of the particle dynamics on planet eccentricity.  For orbits with eccentricity equal to the forced eccentricity there is a  region in the plane with longer lived orbits  (compared to the planet orbital period) in close proximity to eccentric planets. Three dimensional simulations that incorporate collisions are needed to see if these orbits tend to be populated by long lived particles, as proposed for Fomalhaut's eccentric ring \\citep{quillen06b}.  \\vskip 0.1truein We thank the referee, Dr. Tsiganis, for numerous comments that have significantly improved this manuscript. Support for this work was in part provided by National Science Foundation grants AST-0406823 and PHY-0552695 the National Aeronautics and Space Administration under Grant No.~NNG04GM12G issued through the Origins of Solar Systems Program, and HST-AR-10972 to the Space Telescope Science Institute."
    },
    "0608/astro-ph0608603_arXiv.txt": {
        "abstract": "We discuss color selection of rare objects in a wide-field, multiband survey spanning from the optical to the mid-infrared.  Simple color criteria simultaneously identify and distinguish two of the most sought after astrophysical sources:  the coolest brown dwarfs and the most distant quasars.  We present spectroscopically-confirmed examples of each class identified in the IRAC Shallow Survey of the Bo\\\"otes field of the NOAO Deep Wide-Field Survey.  \\tdwarf\\ is a T4.5 brown dwarf at a distance of approximately 42~pc, and \\qso\\ is a radio-loud quasar at redshift $z = 6.12$.  Our selection criteria identify a total of four candidates over 8 square degrees of the Bo\\\"otes field.  The other two candidates are both confirmed $5.5 < z < 6$ quasars, previously reported by Cool et al. (2006).  We discuss the implications of these discoveries and conclude that there are excellent prospects for extending such searches to cooler brown dwarfs and higher redshift quasars. ",
        "introduction": "Wide-area surveys are one of the most powerful tools for observational astronomy, and have led to discoveries ranging from Earth-crossing asteroids to the most distant quasars.  Historically, when technology allows a wavelength regime to be newly probed, either in terms of sensitivity or area, one of the first tasks is a large, shallow survey to see what astrophysical phenomena lurk in the uncovered territory. In recent years, major advances from this line of research include the discovery of the coolest Galactic stars by the Two Micron All Sky Survey (2MASS), ultraluminous infrared galaxies by the {\\it Infrared Astronomical Satellite}, the most distant quasars by the Sloan Digital Sky Survey (SDSS), and the power spectrum of the cosmic microwave background, first by the {\\it Cosmic Background Explorer} and later refined by the {\\it Wilkinson Microwave Anisotropy Probe}.  Such fundamental scientific discoveries have been a major incentive and reward for NASA's Explorer program and other large projects.  The pace of scientific discovery relies on such programs continuing.  The mid-infrared regime has been made newly accessible by the launch of the {\\it Spitzer Space Telescope} \\markcite{Werner:04}(Werner {et~al.} 2004).  At its least competitive, shortest waveband, $3.6 \\mu$m, {\\it Spitzer} is still more than five orders of magnitude more efficient than the largest ground-based observatories for areal surveys.  For the longest wavebands, ground-based observations are simply not possible.  Even compared to previous space-based missions, {\\it Spitzer} offers several orders of magnitude increase in mapping efficiency.  Thus inspired, we have undertaken a shallow, wide-area 3.6 to $8.0~\\mu$m survey with {\\it Spitzer}, summarized in \\S2.  We discuss two of the rare, interesting astronomical sources which are ideally suited to selection by combining deep optical data with shallow mid-infrared data:  the coolest Galactic brown dwarfs and the most distant quasars. The former, of course, are not actually rare in the cosmos; their faint optical magnitudes merely delayed their discovery until recent years and continue to make them ``rare'' in terms of known, spectroscopically-confirmed examples.  Section~3 discusses the selection criteria used to identify such sources and \\S4 describes our spectroscopic observations which confirmed both a cool brown dwarf (spectral class T4.5) and a high-redshift ($z = 6.12$) quasar.  The implications for these discoveries are described in \\S4, and \\S5 summarizes the results and discusses future prospects.  Throughout we adopt a ($\\Omega_m, \\Omega_\\Lambda$) = ($0.3, 0.7$) flat cosmology and $H_0 = 70 \\kmsMpc$. Unless otherwise stated, all magnitudes are quoted in the Vega system. ",
        "conclusions": "Initial spectroscopic follow-up of candidates was obtained with the Multi-Aperture Red Spectrometer \\markcite{Barden:01}(MARS; Barden {et~al.} 2001) on the Mayall 4m telescope at Kitt Peak.  MARS is an optical spectrograph which uses a high resistivity, p-channel Lawrence Berkeley National Laboratory CCD with little fringing and very high throughput at long wavelengths ($\\simlt 10,500$ \\AA).  On the nights of UT 2006 March 24 $-$ 26, we obtained spectra of red sources in the Bo\\\"otes field using the 1\\farcs7 wide long slit, OG550 order-sorting filter, and the VG8050 grism. Across much of the optical window, the instrument configuration provides resolution $R \\approx 1100$ spectra, as measured from sky lines filling the slit.  \\tdwarf\\ was observed for 1.5~hr on UT 2006 March 24, split into three dithered 1800~s exposures.  \\qso\\ was observed for 1~hr on UT 2006 March 25, split into three 1200~s exposures.  The data were processed following standard optical, slit spectroscopy procedures. The nights were not photometric, but relative flux calibration of the spectra was achieved with observations of the spectrophotometric standards Feige~34 and PG~0823+546 \\markcite{Massey:90}(Massey \\& Gronwall 1990) obtained during the same observing run.  The extracted, calibrated MARS spectra are presented in Figs.~\\ref{fig.tdwarfopt} and \\ref{fig.qso}.  The bright star 3\\farcs3 east of \\qso\\ made extraction of the fainter target challenging, resulting in systematic fluctuations of the background at the 1~$\\mu$Jy level.  Near-infrared spectroscopy of \\tdwarf\\ was obtained with the cryogenic, cross-dispersed  Near-Infrared Echelle Spectrograph \\markcite{Mclean:98}(NIRSPEC; McLean {et~al.} 1998) on the Keck~II 10m telescope atop Mauna Kea. We first obtained $J$- and $H$-band spectroscopy on UT 2006 April 05. An AB nod sequence with a total on-source integration time of 200~s per grating setting was used.  For both grating settings, the G2~V star GSPC P300-E from \\markcite{Colina:97}Colina \\& Bohlin (1997) was used for both telluric correction and flux calibration.  An additional $J$-band spectrum was acquired on UT 2006 May 11.  On this night a 300 s integration was taken both on-source and off-source, and the F0 star BD+66 1089 was acquired for telluric correction and flux calibration.  Fig.~\\ref{fig.tdwarf} presents the combined near-infrared spectrum.     \\subsection{\\tdwarf:  Mid-T Brown Dwarf}   \\begin{figure}[!t] \\begin{center} \\plotfiddle{f3.eps}{3.1in}{-90}{45}{45}{-180}{260} \\end{center}  \\caption{Optical spectrum of \\tdwarf, optically classified as a T5$\\pm$2 brown dwarf, obtained with the MARS spectrograph on KPNO~4m telescope. The relative flux calibration was determined from observations of standard stars from the same observing runs with the same instrumental configurations.  The spectrophotometric scale was estimated from the imaging.  The dotted spectrum shows 2MASS~J055919.14$-$140448.8, classified as a T5 brown dwarf at optical wavelengths (Burgasser et al. 2003a).}  \\label{fig.tdwarfopt} \\end{figure}   \\begin{figure}[!t] \\begin{center} \\plotfiddle{f4.eps}{3.1in}{-90}{45}{45}{-180}{260} \\end{center}  \\caption{Near-infrared spectrum of \\tdwarf\\ obtained with the NIRSPEC spectrograph on the Keck~II telescope, classified as a Galactic T4.5 brown dwarf from these data.  The relative flux calibration was determined from observations of standard stars from the same observing runs with the same instrumental configurations.  The spectrophotometric scale was estimated from the imaging.  The dotted line shows the infrared spectrum of 2MASS~J0559$-$1404, classified as a T4.5 brown dwarf in the near-infrared (McLean et al. 2003).}  \\label{fig.tdwarf} \\end{figure}   The spectrum of \\tdwarf\\ shows the classic signatures of a T dwarf. The optical spectrum in Fig.~\\ref{fig.tdwarfopt} shows a sharp rise to the longest wavelengths, indicative of a cool temperature and strong absorption by the pressure-broadened wings of \\ion{K}{1} (and to some extent \\ion{Na}{1}).  Even more telling are the $J$- and $H$-band spectra in Fig.~\\ref{fig.tdwarf} that show strong CH$_4$ and H$_2$O absorption, the former of which is the hallmark of spectral class T.  As this object has both optical and near-infrared spectra, we can classify on both the optical and near-infrared classification schemes. The optical typing of T dwarfs is somewhat crude because the $\\leq 1 \\mu$m spectra show less variation than at longer wavelengths.  Nonetheless, \\markcite{Burgasser:03}Burgasser {et~al.} (2003a) have established standards for classes T2, T5, T6, and T8.  In the $6000 - 10000$ \\AA\\ range the best diagnostic is the 9300 \\AA\\ band of H$_2$O.  Unfortunately our MARS spectrum has not been telluric corrected so the depth of this water feature will be influenced by both the earth's atmosphere as well as the atmosphere of the brown dwarf itself.  This feature in \\tdwarf\\ is not as deep as in the spectrum of a T8, so the true spectral type must be earlier than that.  Comparisons with the T2, T5, and T6 standards obtained with Keck \\markcite{Burgasser:03}(Burgasser {et~al.} 2003a) show that the overall slope most resembles that of the T5.  Given the coarseness of classification in this wavelength regime, we can assign only a crude optical spectral type of T5$\\pm$2.  In the near-infrared the situation is much improved.  In this wavelength regime there is a full set of standards for each spectral subtype from T0 to T8 \\markcite{Burgasser:06}(Burgasser {et~al.} 2006).  Using Keck NIRSPEC spectra from \\markcite{Mclean:03}McLean {et~al.} (2003) of the \\markcite{Burgasser:06}Burgasser {et~al.} (2006) standards, we find that the individual $J$-band spectra best match a type intermediate between T4 and T5. A similar fit to the $H$-band data alone gives the identical result. These results point to a solid near-infrared spectral type of T4.5.  Shown in Figs.~\\ref{fig.tdwarfopt} - \\ref{fig.tdwarf} are comparisons of the spectra of \\tdwarf\\ and 2MASS~J0559$-$1404, which is the optical T5 standard and typed as T4.5 on the \\markcite{Burgasser:06}Burgasser {et~al.} (2006) near-infrared scheme. (That is, 2MASS~J0559$-$1404 has the same type as \\tdwarf\\ in both wavelength regimes.)  Note the similarities between the two spectra.  2MASS~J0559$-$1404 has a well measured trigonometric parallax of 97.7$\\pm$1.3 mas \\markcite{Dahn:02}(Dahn {et~al.} 2002) and an absolute magnitude of $M_J = 13.75 \\pm 0.04$, which allows us to estimate a distance to \\tdwarf\\ of 42 pc, assuming both T dwarfs are single.  However, 2MASS~J0559$-$1404 is the most overluminous object in the early-/mid-T ``hump'' on the Hertzsprung-Russell diagram \\markcite{Vrba:04, Golimowski:04}(see Vrba {et~al.} 2004; Golimoski {et~al.} 2004), leading some researchers to believe that it might be a close, equal-magnitude double despite all current evidence to the contrary \\markcite{Burgasser:03b, Gelino:05, Liu:06}(Burgasser {et~al.} 2003b; Gelino \\& Kulkarni 2005; Liu {et~al.} 2006).  Correcting for this possibility, we find that \\tdwarf\\ might be as close as 30 pc.  No other optically classified T dwarfs are known of spectral type T5; the only other T dwarf with a measured trigonometric parallax and near-infrared type of T4.5 is SDSS~J020742.48+000056.2.  The parallax measurement of 34.85$\\pm$9.87 mas for SDSS~J0207+0000 \\markcite{Vrba:04}(Vrba {et~al.} 2004) implies $M_J = 14.51 \\pm 0.64$, which is very uncertain but lends some weak support to the closer distance estimate for \\tdwarf.   The first images to detect the brown dwarf were the NDWFS $I$-band observations obtained on UT 2000 April 28, 3.7~yr prior to the IRAC imaging.  Comparing ten nearby sources detected in both the $I$-band and $3.6 \\mu$m observations, \\tdwarf\\ has a detected proper motion of $0\\farcs1 \\pm 0\\farcs03\\, {\\rm yr}^{-1}$, in a southerly direction. This is comparable in amplitude to the expected reflex solar motion for a source at $\\approx 40$~pc.  Interestingly, the star 5\\farcs7 east of the brown dwarf shows a higher proper motion, $\\mu = 0\\farcs3 \\pm 0\\farcs03\\, {\\rm yr}^{-1}$ in the NW direction.  The colors of this $R = 22.3$ star (ISS~J142951.3+333010) are relatively blue, $B_W - R = 0.6, R-I = 0.6$, suggesting a relatively hot white dwarf at a distance of several hundred pc, moving at several hundred km sec$^{-1}$.  The $4.5 \\mu$m flux of the brown dwarf is 2.7~mag brighter than the survey limit (\\ie $V/V_{\\rm max} = 0.024$), whereas the $I$ magnitude is only 0.9~mag above the limit ($V/V_{\\rm max} = 0.3$).  This suggests that the $I < 23$ requirement imposed to provide robust morphological selection of unresolved sources is a significant limiting factor.  We estimate that our selection criteria restrict our sensitivity to brown dwarfs of spectral type T3 to T6.  The former limit comes from the IRAC color criterion \\markcite{Patten:06}(Patten {et~al.} 2006).  The latter limit comes from available data (J.D.~Kirkpatrick \\etal in prep.) suggesting that $I$-band flux drops dramatically for spectral types cooler than T6.  From \\markcite{Vrba:04}Vrba {et~al.} (2004) and \\markcite{Golimowski:04}Golimoski {et~al.} (2004), the range T3 to T6 corresponds very roughly to $T_{\\rm eff} = 1500$ to $1100$~K.  Using a model which forms brown dwarfs at a constant rate over 10~Gyr with power law mass functions of index 0.4 to 1.3 \\markcite{Reid:02}(Reid, Gizis, \\& Hawley 2002) and the theoretical models of \\markcite{Burrows:03}Burrows, Sudarsky, \\& Lunine (2003) which give luminosities and $T_{\\rm eff}$ as a function of brown dwarf mass and age, we expect $3 - 5$ brown dwarfs in the IRAC shallow survey to meet our selection criteria.  Intriguingly, there should be a similar number of dwarfs with $T_{\\rm eff} < 750$~K above the [4.5] flux limit, although our $I < 23$ requirement would exclude them from the present sample.  Given our desire to understand more fully the physical nature of the L/T transition, the newly discovered T4.5 brown dwarf can serve as another probe of the overluminosity of the early-/mid-T hump.  Its magnitudes of $J = 16.88$ and $K_s = 16.99$ make it a difficult but not impossible target for a dedicated near-infrared parallax program such as the on-going one at the US Naval Observatory in Flagstaff \\markcite{Vrba:04}(Vrba {et~al.} 2004).  More importantly, \\tdwarf\\ is the first example of a field T dwarf selected by mid-infrared photometry supplemented by other ground-based optical and near-infrared data.  This implies that a very similar selection technique to be employed by the {\\it Wide-Field Infrared Survey Explorer} \\markcite{Eisenhardt:03}({\\it WISE}; Eisenhardt \\& Wright 2003), planned for launch in 2009, is sound and will be capable of discovering other T dwarfs, and hopefully cooler Y dwarfs \\markcite{Kirkpatrick:03}(Kirkpatrick 2003).  {\\it WISE} will sample hundreds of times more volume than the IRAC Shallow Survey in bands similar to [3.6] and [4.5], and should reveal whether there are brown dwarfs closer to the Sun than Proxima Centauri.     \\subsection{\\qso:  $z = 6.12$ Quasar}  \\begin{figure}[!t] \\begin{center} \\plotfiddle{f5.eps}{3.1in}{-90}{45}{45}{-180}{260} \\end{center}  \\caption{Spectrum of \\qso, a quasar at $z = 6.12$, obtained with the MARS spectrograph on the Kitt Peak 4m Mayall telescope.   The relative flux calibration was determined from observations of standard stars from the same observing run with the same instrumental configuration. As the nights were not photometric, the spectrophotometric scale has been estimated from the imaging.}  \\label{fig.qso} \\end{figure}   The spectrum of \\qso\\ (Fig.~\\ref{fig.qso}) clearly shows the strong Ly$\\alpha$ emission and strong Ly$\\alpha$ decrement of a $z \\geq 6$ quasar.  At a redshift of $z = 6.12$, \\qso\\ is emitting when the universe was 0.89~Gyr old, or only 7\\%\\ of its current age.  This is the tenth $z \\geq 6$ quasar identified to date, with the prior nine identified by the Sloan Digital Sky Survey \\markcite{Fan:06}(SDSS; Fan {et~al.} 2006).  \\qso\\ was identified independently by \\markcite{McGreer:06}McGreer {et~al.} (2006) using different selection criteria. Two characteristics separate \\qso\\ from the other nine $z \\geq 6$ quasars known.  First, while the other nine were identified from 6550 deg$^2$ of the wide-area, shallow SDSS optical survey with $J$-band follow-up, \\qso\\ was identified in a more sensitive, multi-wavelength survey of only 8 deg$^2$.  Consequently, this is the least luminous quasar known at $z \\approx 6$.  Secondly, the Faint Images of the Radio Sky at Twenty-cm survey \\markcite{Becker:95}(FIRST; Becker, White, \\& Helfand 1995) identifies a source with an integrated flux of 1.03~mJy within 1 arcsec of the quasar coordinates. \\qso\\ is thus the only $z \\geq 6$ mJy-radio source currently known.  The evolution of the fraction of quasars which are radio loud, and, in fact, the definition and very existence of such a dichotomy, has been the subject of substantial literature.  Some researchers prefer a definition based on the radio-optical ratio $R_{\\rm ro}$ of the specific fluxes at rest-frame 6~cm (5~GHz) and 4400 \\AA\\ \\markcite{Kellerman:89}(Kellerman {et~al.} 1989). The other common definition divides the populations at some rest-frame radio luminosity; \\eg \\markcite{Gregg:96}Gregg {et~al.} (1996) uses a cutoff value for the 1.4 GHz specific luminosity, $L_{\\rm 1.4~GHz} = 10^{32.5}\\, h_{50}^{-2}\\, {\\rm ergs}\\, {\\rm s}^{-1}\\, {\\rm Hz}^{-1}$ to separate radio loud and radio quiet sources.\\footnote{An Einstein-de~Sitter cosmology is assumed.} The latter definition is immune to obscuration from dust, and, as argued by \\markcite{Peacock:86}Peacock, Miller, \\& Longair (1986) and \\markcite{Miller:90}Miller, Peacock, \\& Mead (1990), is the more physically meaningful definition.  Based on radio observations of all $z > 4$ quasars known as of mid-1999 and using the radio luminosity definition, \\markcite{Stern:00a}Stern {et~al.} (2000) found that approximately 12\\%\\ of quasars are radio loud, with no evidence of this fraction depending on either redshift (for $2 \\simlt z \\simlt 5$) or optical luminosity (for $-25 \\simgt M_B \\simgt -28$).  For a typical radio spectral index $\\alpha = -0.5$ and an Einstein-de~Sitter cosmology for comparison with previous literature, \\qso\\ has a radio luminosity of $L_{\\rm 1.4~GHz} = 1.33 \\times 10^{33}\\, h_{50}^{-2}\\, {\\rm ergs}\\, {\\rm s}^{-1}\\, {\\rm Hz}^{-1}$, classifying it as radio loud.  \\markcite{McGreer:06}McGreer {et~al.} (2006) show that this source is still classified as radio loud based on a radio-optical ratio definition. \\qso\\ is thus the most distant radio-loud quasar known.   \\qso\\ is only slightly fainter in luminosity than the $z \\geq 6$ SDSS quasars, so we consider all ten $z > 6$ quasars as a single sample, deferring issues of the likelihood of our having found such a source in our drastically smaller survey (discussed next).  The implication is that the radio loud fraction remains near 10\\%\\ out to $z \\approx 6.5$. Conventional wisdom and morphological studies suggest that luminous, radio loud AGN are preferentially identified with early-type galaxies \\markcite{McLure:99}(\\eg McLure {et~al.} 1999).  Theory can explain the trend, since early-type galaxies are likely the products of major mergers and two coalescing supermassive black holes appear necessary to create black holes of sufficient spin to generate highly collimated jets and powerful radio sources \\markcite{Wilson:95}(\\eg Wilson \\& Colbert 1995).  Assuming the radio loud -- luminous host galaxy relation remains robust at high redshift, the apparent discovery that $\\approx 10\\%$ of quasars are radio loud out to the highest redshifts probed has interesting implications for the formation epoch of massive galaxies. In hierarchical models of galaxy formation, late-type (less massive) systems form first and mergers are required to form the early-type (more massive) systems.  Eventually, therefore, one expects the radio-loud fraction of AGNs to fall precipitously with redshift. Our results show this epoch lies beyond $z \\approx 6$, providing further evidence for an early formation epoch for massive galaxies.  The stellar masses of $i$-dropout galaxies in the Great Observatories Origins Deep Survey \\markcite{Giavalisco:04a}(Giavalisco {et~al.} 2004) leads to a similar conclusion from a very different data set and line of argument \\markcite{Yan:05, Yan:06, Eyles:06}(Yan {et~al.} 2005, 2006; Eyles {et~al.} 2006).    How likely was the discovery of this distant quasar in an 8 deg$^2$ field? Interpolating the $K_s$ and $3.6~\\mu$m photometry for \\qso\\ implies $m_{\\rm AB}[(1+z) 4400 {\\rm \\AA}] \\approx 19.6$, or $M_{\\rm AB}(4400) = -27.2$.  For a typical quasar optical spectral index, the conversion between AB-system $M_{\\rm AB}(4400)$ and Vega-system $M_B$ is $M_B = M_{\\rm AB}(4400) + 0.12$ \\markcite{Stern:00a}(\\eg Stern {et~al.} 2000), implying $M_B = -27.1$ for \\qso.  Our $J$-band photometry implies a continuum flux density of $\\approx 18~\\mu$Jy at $1~\\mu$m, or a rest-frame UV luminosity of $M(1450) = -26.0$, making this source fainter than any of the $z \\approx 6$ quasars identified by the SDSS \\markcite{Fan:06}(Fan {et~al.} 2006).  \\markcite{McGreer:06}McGreer {et~al.} (2006) found a slightly brighter absolute magnitude, $M(1450) = -26.4$, likely due to their alternate methodology whereby a quasar template fit to the IRAC data was used to derive the rest-frame UV luminosities.  We estimate the number of high-redshift quasars expected from our selection criteria using the \\markcite{Fan:04}Fan {et~al.} (2004) high-redshift quasar luminosity function, derived from the SDSS.  We approximate high-redshift quasar spectra as step functions, with zero flux below redshifted Ly$\\alpha$ and a flat SED (in $f_\\nu$) redward of Ly$\\alpha$, and we approximate the NDWFS $I$-band filter as a tophat function.  Our selection criteria restrict our sensitivity to quasars at $5.5 \\simlt z \\simlt 6.5$. The lower redshift limit comes from the $R - I$ color requirement, determined from the \\markcite{Richards:06}Richards {et~al.} (2006) model discussed in \\S3; indeed, the $z = 5.39$ quasar identifed by \\markcite{Cool:06}Cool {et~al.} (2006) is too blue in $R - I$ to meet our selection criteria (Fig.~\\ref{fig.colcol}).  The upper redshift limit corresponds to Ly$\\alpha$ shifting out of the $I$-band filter.  The \\markcite{Fan:04}Fan {et~al.} (2004) luminosity function predicts 3.3 quasars at $5.5 < z < 6.5$ with $I < 23$ in our 8 deg$^2$ survey.  This prediction exactly matches the current results, though, notably, the faintest of the high-redshift Bo\\\"otes quasars has $I = 22.0$, suggesting that more quasars remain to be discovered with $22 < I < 23$ and that the faint end slope of the high-redshift quasar luminosity function is steeper than currently assumed.  Of the 3.3 quasars predicted at $5.5 < z < 6.5$, only 0.3 are expected to be at $z > 6$, or, for 12\\%\\ of quasars being radio-loud \\markcite{Stern:00a}(Stern {et~al.} 2000), we only had a 4\\%\\ chance of identifying a $z > 6$ radio-loud quasar in this survey.  While it is premature to make strong claims from this small sample, our results imply possible rapid evolution in the faint end of the quasar luminosity function and in the radio loud fraction at high redshift."
    },
    "0608/astro-ph0608329_arXiv.txt": {
        "abstract": "Errors in the kinematic distances, under the assumption of circular gas orbits, were estimated by performing synthetic observations of a model disk galaxy. It was found that the error is $< 0.5 \\kpc$ for most of the disk when the measured rotation curve was used, but larger if the real rotation curve is applied. In both cases, the error is significantly larger at the positions of the spiral arms. The error structure is such that, when kinematic distances are used to develope a picture of the large scale density distribution, the most significant features of the numerical model are significantly distorted or absent, while spurious structure appears. By considering the full velocity field in the calculation of the kinematic distances, most of the original density structures can be recovered. ",
        "introduction": "Since the classic work by \\citet{oor58}, there have been many attempts to use the kinematic properties of the diffuse gas to determine the large scale spiral structure of the Milky Way. Very early in the study of the Galaxy, it was determined that the orbits of the disk components of the Galaxy were not very different from circular, with an orbital frequency that decreases monotonically as a function of galactocentric radius. These facts allow the use of the kinematic distance method as a first approximation to map the gaseous component of the galactic disk. Two of the main strengths of this method (that it can be used for a very large fraction of the Galaxy and that it can be applied to the gaseous component of disk, which is notoriously difficult to obtain a distance to) make it particularly useful for this goal. Nevertheless, it was soon realized that the deviations from circular orbits, however small in absolute value, might have a strong impact on how we see the Galaxy.  One of the first difficulties of the kinematic distance method appeared in the determination of the rotation curve, namely, in the fact that the circular rotation velocity measured for positive galactic longitudes (northern Galaxy) did not match the one measured for negative longitudes (southern Galaxy). The simplest way to reconcile these observed rotation laws is to take their average, assuming that the differences generated by non-axisymmetric structure will cancel out. \\citet{ker62} showed that this approximation leads to large north-south asymmetries that, given their heliocentric nature, seemed unlikely. It became clear that the complex kinematic structure revealed in the diffuse gas surveys, like the presence of gas at forbidden velocities, or the oscillations in the rotation curve, was itself a consequence of the spiral structure that was being sought. Given the importance of the rotation curve for the understanding of the galactic dynamics, in addition of the determination of kinematic distances, a different approach to measuring the rotation curve was needed.  A frequently used method to obtain the galactic rotation curve involves measuring the full velocity field of discreet sources that might share the motion of the diffuse gas (young objects like \\ion{H}{2} regions, for example), and then averaging the so obtained azimuthal velocities. This approach has lead to models of the rotation curve \\citep{bra93,mac05} that might more closely trace the real mass distribution of the Galaxy, but introduce new sources of error when used to obtain kinematic distances. Nevertheless, generally accepted models of the spiral structure \\citep{geo76,tay93} have been obtained using this assumption. \\citep[For a nice review of the early work, see][]{ker69}  Another approach involves the modelling of the non-circular motions of the gas, instead of forcing the assumption of circular orbits. In a recent paper, \\citet{fos06} used an analytic approach for the velocity field of the outer Galaxy. In this work, a numerical model of the galactic disk, with full MHD, is used to further explore the effects of non-circular motions in the image one would obtain of the Galaxy when we rely on the kinematic method for distances. An observer is imagined inside the numerical model, which is assumed similar to the Milky Way, and the analysis that this observer would perform is reproduced. In \\S \\ref{simulation_sec}, a brief description of the numerical simulation is presented; in \\S \\ref{observations_sec} the selection of the observer's position is described, and how the measurement of the rotation curve was emulated; \\S \\ref{errors_sec} presents an analysis of the errors in the kinematic distances and how they affect the image the observer generates of his/her home galaxy; finally, \\S \\ref{discusion_sec} summarizes the results. ",
        "conclusions": "\\label{discusion_sec}  The effect of the circular orbits assumption on our idea of the large scale structure of the Galaxy was explored. Since these errors might be quite large at the position of the spiral arms, the study of the spiral structure of the Galaxy and objects associated with it is particularly affected. By simulating the way an imaginary observer inside the model galaxy might try to infer the structure of the gaseous disk, it was found that the circular orbits assumption destroys the spiral structure and creates spurious features in the measured distribution.  The method of kinematic distances is a powerful one since it allows measurement of distances to diffuse sources and it is easily applicable to a large fraction of the galactic disk. Even if the measured rotation curve includes deviations that do not reflect the true large scale mass distribution, Figure \\ref{lon_err_fig}(a) shows that the errors in the distance are, in fact not very large for most of the galactic disk; in fact, the distance errors that arise from using the true rotation curve are larger. In both cases, however, the errors are quite large at the positions of the spiral arms. If we want to use this distance method for objects associated with the spiral structure, we need to consider non-circular motions (as has been succesfully shown by \\citealp{fos06} for a set of \\ion{H}{2} regions and SNR).  One possibility to achieve this is to try to determine the full velocity field of the galactic disk. But direct measure of distances to the diffuse gas component is quite difficult (therefore the strength of the kinematic distance method). So, we need to use discreet objects and assume that they share their velocity with the diffuse component (\\citealp[ for example]{bra93}; \\citealp[see also discussion in][]{min73}). Yet another difficulty arises when tangential velocities and distances are required beyond the solar neighborhood.  Another approach at determining the full velocity field is to model it. Recently, \\citet{fos06} used an analytical model of the density and velocity fields of the diffuse gas, with parameters for the model fitted to \\ion{H}{1} observations of the outer Galaxy. Despite the fact that their density and velocity models are not consistent in the hydrodynamics sense, and that the model do not include the dynamical effects of magnetic fields, they were able to add features of the Galaxy that are currently difficult to incorporate to numerical models, like the disk's warp or the rolling motions associated with the spiral arms. Further numerical studies should allow the development of a more realistic analytical model.  Instead of an analytic model, a numerical model was used in the present work to obtain density and velocity fields. Since the focus is in large scale velocity structures, an eulerian code provides a good approach. Also, since the galactic magnetic field has been proved to be an important component of the total ISM pressure \\citep{bou90}, its effect in the gas dynamics is likely to be important; therefore, a full MHD simulation was required. The large scale forcing is also trascendant; since the azimuthal shape of the spiral perturbation appears to have an influence on the gaseous response \\citep{fra02}, the usual sinusoidal perturbation was deemed too simplistic and a self-consistent model for the perturbing arms was chosen. At the present time, the galactic warp and the vertical motions associated with the spiral arms \\citep{gom04a,gom04b} could not be considered at the necessary resolution.  In this work, it has been shown that it is possible to recover most of the gaseous structure of the galactic disk using kinematic distances, as long as the full velocity field is considered. Nevertheless, applying these results to the Milky Way is a whole new issue, since obtaining the full velocity field is not trivial. For the procedure used here, how close the numerical simulation is to the real Galaxy remains the weak point of this approach. The computation cost of a realistic enough simulation is still too high to allow a parameter fitting analysis. So, the remaining question is if the velocity field that results from the simulation yields a determination of the distance to a given object, or only an estimation of the distance error. The answer to that question is left to the reader's criterion."
    },
    "0608/astro-ph0608172.txt": {
        "abstract": "{} % {We study the uniformity of the distribution of compact flat-spectrum AGN on the sky and the evolution of their relativistic jets with cosmic epoch.} % {A complete sample of compact extragalactic radio sources at 15\\,GHz was recently compiled to conduct the MOJAVE program. The MOJAVE sample comprises 133 radio-loud flat-spectrum AGN with compact relativistic outflows detected at parsec scales.} % {Analysis of the population of flat-spectrum quasars of the sample reveals that the pc-scale jets of quasars have intrinsic luminosities in the range between $\\sim10^{24}\\,{\\rm W\\,Hz^{-1}}$ and $\\sim10^{27}\\,{\\rm W\\,Hz^{-1}}$ and Lorentz factors distributed between $3\\la\\gamma \\la30$. We find that the apparent speed (or Lorentz factor) of jets evolves with redshift, increasing from $z\\sim0$ to $z\\sim1$ and then falling at higher redshifts ($z\\sim2.5$) by a factor of 2.5. The evolution of apparent speeds does not affect significantly the evolution of the beamed luminosity function of quasars, which is most likely to be dependent on the evolution of radio luminosity. Furthermore, the beamed radio luminosity function suggests that the intrinsic luminosity function of quasars has a double power-law form: it is flat at low luminosities and steep at high luminosities. There is a positive evolution of quasars at low redshifts ($z<0.5$) and strong negative evolution at redshifts $>1.7$ with space density decline up to $z\\sim2.5$. This implies that the powerful jets were more populous at redshifts between $0.5$ and $1.7$. We show that the evolution of compact quasars is luminosity dependent and it depends strongly on the speed of the jet suggesting that there are two distinct populations of quasars with slow and fast jets which evolve differently with redshift.} % {} ",
        "introduction": "The advantage of studying the cosmological evolution of radio sources free of dust obscuration was realized by \\cite{longair66}. He demonstrated that the cosmic evolution of low and high luminosity radio sources sampled at low radio frequency (178\\,MHz) are different with later ones having much stronger evolution. The morphological difference between FRI and FRII radio sources \\citep{fanaroff74} and their division into low and high luminosity classes below and above $L_{\\rm 178\\,MHz}\\approx10^{25}$\\,W\\,Hz$^{-1}$\\,sr$^{-1}$ was used by \\cite{wall80} to show that the population of FRI radio sources does not evolve strongly with cosmic time whilst the evolution of FRII sources is much stronger.  The orientation-based unification scheme \\citep{scheuer87,barthel89} suggests that a single parent population of radio sources appears as radio galaxies and quasars when viewed at different angles to the line of sight, with quasars beeing seen from $45^{\\circ}$ to the line of sight and radio galaxies from larger angles. It was successfully used to describe a diverse appearance of observed characteristics of radio sources and their evolution \\citep{urry95,jackson99,grimes04} at different radio frequencies. The orientation of the jet and its beamed emission are two important keys to identify that the BL Lacs and quasars are beamed parent objects of FRI and FRII radio sources.  At low frequencies (178\\,MHz), these sources appear as extended (lobe-dominated) steep-spectrum FRI and FRII radio sources. \\cite{willott01} used the combined sample of 3CRR, 7C and 6CE surveys (see references in Willott et al. 2001) to model the radio luminosity function of FRI/FRII radio sources. They showed that a dual-population model of luminosity function fits the data well, but it requires a differential positive density evolution between $z\\sim0$ and $z\\sim2$. No evidence of a negative evolution (redshift `cut-off') is found at high redshifts.  The surveys at frequencies $\\ge1.4$\\,GHz include both lobe- and core-dominated sources in which the jet beaming effect appears to be dominant. These radio sources enhanced by Doppler beaming of the jet are identified with flat-spectrum quasars and BL Lacs. At these frequencies, the cosmological evolution of flat-spectrum sources is more complex, involving the evolution of both non-beamed and beamed emission. In addition to the luminosity/density evolution it also may depend on the evolution of relativistic effects with cosmic time. The evolution of flat-spectrum quasars was studied in several works \\citep{dunlop90,jackson99,cirasuolo05,wall05}. \\cite{jackson99} showed that the evolution and beaming of powerful radio sources may be described by a dual-population unification scheme for FRI and FRII sources which, depending on the viewing angle, appear as lobe-dominated steep-spectrum and core-dominated flat-spectrum sources. Using data from a combined sample of radio sources at 2.7\\,GHz \\cite{dunlop90} showed that both pure luminosity and luminosity/density evolution fit the observed redshift and source-count data. They presented clear evidence of a redshift cut-off for both steep- and flat-spectrum quasars at high redshifts, $z\\ga2$. A high redshift cut-off is also evident in the Parkes quarter-Jansky flat-spectrum quasars \\citep{wall05} and in the sample of 352 faint flat-/steep-spectrum quasars \\citep{cirasuolo05} -- with radio fluxes $\\ge$1~mJy at 1.4\\,GHz -- selected from the radio (Faint Images of the Radio Sky at Twenty centimeters) and optical (2dF QSO Redshift Survey) surveys. They found an indication of a significant negative evolution for faint quasars at $z>1.8$ consistent with negative evolution of bright flat-spectrum radio quasars at 2.7\\,GHz, and inconsistent with no-evolution of steep-spectrum bright radio quasars at 151\\,MHz \\citep{willott98}. The reduction of the comoving space density of quasars at high redshifts is also found in X-ray and optically selected samples (e.g., Schmidt et al. 1991; Fan et al. 2001 and references therein; Hasinger 2004; Silverman et al. 2005) indicating that the redshift cut-off of optically-selected quasars is not a selection effect due to obscuration by dust.  It is not clear whether the redshift cut-off of flat-spectrum sources is present in high radio frequency samples selected on the basis of beamed emission and whether the cosmological evolution of beamed sources depends on the properties of the jets such as the intrinsic luminosity and Lorentz factor of the jet (the analytical relationships between the latter parameters were studied by \\cite{lister97} for simulated flux-limited samples of relativistic jets). %The analytical relationships %between the latter parameters were studied by \\cite{lister97} for a %simulated flux-limited samples of relativistic jets. The predicted %flux density, redshift and apparent speed distributions were compared %to the Caltech-Jodrell Bank sample of bright flat-spectrum AGN. %They concluded that the observational data of the sample are %insufficient to distinguish between correlated and independent models %of the luminosity--Lorentz factor relation. To examine these problems we use the (first) complete sample of 133 bright AGN at 15\\,GHz which has been compiled recently by \\cite{lister05} as an extension of the Very Long Baseline Array (VLBA) 2\\,cm survey \\citep{kellermann98,zensus02} and other programs.  This sample of flat-spectrum AGN was designed for a long-term observational program (named the MOJAVE survey) to study the structure and evolution of compact jets. The radio emission of AGN originates in the compact jets at scales from a few parsecs to tens of parsecs. Although the sample is small, includes relatively bright sources and is restricted to $z\\la2.7$, it provides a wealth of information about radio emission, structure and kinematics of parsec-scale jets. The completeness of the sample makes it possible to investigate the cosmological evolution of flat-spectrum AGN in the context of evolution of the jet parameters. In Sect. 2, the complete sample of flat-spectrum AGN is presented. In Sect. 3, we investigate the homogeneity of the MOJAVE sources on the sky. The source counts of AGN, evolution of apparent speeds of jets, and luminosity function and its evolution with cosmic epoch are examined in Sect. 4. Section 5 involves the summary of results. In this paper we use a flat cosmology ($\\Omega_{\\rm m} + \\Omega_{\\rm \\Lambda}= 1$) with non-zero lambda, $\\Omega_{\\rm m} = 0.3$ and $\\rm H_{\\rm 0} = 70\\, km\\,s^{-1}\\,Mpc^{-1}$.  %The observed radio %luminosity of BL Lacs and quasars at 15\\,GHz is highly boosted and %radiated predominantly by the approaching jet (references XXX). %Earlier studies \\citep{urry84,urry91,lister03} provided analytical %predictions for the %beamed luminosity function of the jets assuming   %The intrinsic radio luminosity %function is altered by the Doppler beaming of the jet which is a %function of the speed and viewng angle of the jet. Earlier studies %\\citep{urry84,urry91,lister03} provided analytical predictions for the %beamed luminosity function of the jets assuming   % The %beamed radio luminosity is a function of intrinsic (non-boosted) %luminosity and Doppler factor of the jet which depends on the speed %and viewng angle of the jet.   %__________________________________________________________________ ",
        "conclusions": "The principal results of statistical analysis of bright flat-spectrum compact AGN from the flux-limited MOJAVE sample are: (a) the distribution of 133 AGN is uniform on the sky; (b) the slope of the source counts of radio sources is aligned with the slopes of source counts of independent high radio frequency surveys suggesting that the MOJAVE is a \\emph{complete} flux-limited sample; (c) Lorentz factors of the jets of quasars lie in the range from $\\sim3$ to $\\sim30$ and Doppler factors of the jets are distributed between $\\delta_{\\rm min}>1$ and $\\delta_{\\rm max}\\la60$ with the realistic upper limit on $\\delta_{\\rm max}\\approx30$ for the MOJAVE sample; (d) the mean apparent transverse speed of jets of quasars (or the mean Lorentz factor) increases by a factor of two (from 6.5 to 14) with increasing redshift up to $z\\sim1$ and then gradually decreases at high redshifts by a factor of 2.5; (e) the intrinsic luminosities of compact jets are distributed in the range from $\\sim10^{24}\\,{\\rm W\\,Hz^{-1}}$ to $\\sim10^{27}\\,{\\rm W\\,Hz^{-1}}$ and their intrinsic radio luminosity function is flat ($B\\sim0$) at low luminosities and becomes steeper ($B\\sim-2.5$) at high luminosities; (f) the weak positive evolution of quasars is evident at low redshifts, $z<0.5$, and a statistically significant redshift cut-off is present at $z \\geq 1.7$; (g) the cosmological evolution is different (1) for high- ($L_{15}>10^{27.9}\\,{\\rm W\\,Hz^{-1}}$) and low-luminosity quasars ($L_{15}<10^{27.9}\\,{\\rm W\\,Hz^{-1}}$) suggesting that their evolution depends strongly on luminosity, and (2) for fast and slow jets, which indicates the existence of two distinct populations of quasars one with slow jets ($\\beta_{\\mathrm a}<6$) and high intrinsic luminosities which does not show a significant evolution with redshift and one with $\\beta_{\\mathrm a}>6$ and low intrinsic luminosities having a pronounced redshift cut-off at $z>1.7$.  The differential evolution of the population of fast and slow jets at $z>1.7$ may explain inconsistent results obtained for a redshift cut-off of quasars sampled at high and low radio frequencies \\citep{dunlop90,willott01} with former one providing no evidence for any decline in the comoving space density. If slow jets reside in high luminosity FRII radio sources and fast jets produce faint FRII structures, which are too weak to be detected at 151\\,MHz, then the low-frequency sample will include the population of quasars with low jet speeds which do not show significant evolution up to $z\\sim2.5$ (Fig.~\\ref{fig:vvm-speed}). The population of fast jets, however, will be included in the high frequency samples because of highly beamed emission of the jet. This naturally explains the redshift cut-off of quasars sampled at 1.4\\,GHz, 2.7\\,GHz \\citep{dunlop90,cirasuolo05,wall05} and significant redshift cut-off at $z>1.7$ found for our sample (Figs.~\\ref{fig:vvm-b-l},\\ref{fig:vvm-speed}).  %% %\\begin{figure} %   \\centering %%   \\resizebox{\\hsize}{!}{\\includegraphics{z-ro.eps}} %   \\includegraphics[angle=-90,width=8.2cm]{Va-D.ps} % %   \\caption{Apparent transverse speed versus Doppler factor of the jet %   for different values of the viewing angle ($\\theta=2^{\\circ}; %   3^{\\circ}; 5^{\\circ}; 10^{\\circ}; 30^{\\circ}; 90^{\\circ}$) of the %   jet and its speed ($\\gamma=2; 5; 10;20; 30; 90$).} % %\\label{fig8} %\\end{figure} %%  %The presence of the positive evolution of compact quasars at %low-redshifts indicates that the activity induced by the jet was much %common at redshifts between $0.5<z<1.7$ with a statistically %significant ($>99.5$\\,\\%, see Table 3) redshift cut-off at $z > %1.7$. This result confirm the redshift cut-off found at slightly %higher redshifts, $z>2$ \\citep{dunlop90}, for flat- and %steep-spectrum quasars sampled at 2.7\\, GHz.  At low-redshifts, the %positive evolution of compact quasars is consistent with slightly %stronger evolution of quasars at 2.7\\,GHz.   %At low-redshifts, $z<1.8$, the strong positive evolution is seen for %bright quasars sampled at 15\\,GHz, 2.7\\,GHz \\citep{dunlop90} and %151\\,MHz \\citep{willott98,willott01}, while the faint quasars %sampled at 1.4\\,GHz do not show evidence for a strong decrement in the %space density. This is a strong evidence that the space density of %quasars peaks at smaller redshifts for fainter quasars (similar to the %evolutionary behaviour of X-ray selected quasars found by %\\cite{ueda03} indicating that the population of faint radio quasars %forms later than the high-luminosity population. The space density %decline at redshifts beyond $z>3$ is also confirmed for optically- and %X-ray-selected quasars (see for example, Schmidt et al. 1991; Fan et %al. 2001 and references therein; Hasinger 2004) %\\cite{hasinger04}).  %{\\tt Discuss the V/Vm and number density evolution in terms of %  $L_{\\rm int}$ and $\\gamma$ of the jet.}  %{\\tt Generate all figures and results for 94 quasars (not 93)}  %\\emph{151/178 MHz surveys}. At these frequencies all sources have %extended radio lobe structures from kpc to Mpc scales. Schmidt (1969) %and Longair \\& Scheuer (1970) studied the evolution of steep-spectrum %quasars and radio galaxies from the complete 3CRR sample (Laing et %al. 1983) compiled at 178 MHz.  They found a positive evolution for %quasars with the $\\langle V/V_{\\rm max} \\rangle = %0.68\\pm0.042$. Willott et al. (2001) used the combined sample of 3CRR, %7C and 6CE surveys (see references in Willott et al. 2001) to model %the radio luminosity function (RLF) of FRI/FRII radio sources. They %showed that a dual-population model of RLF fits the data well, but it %requires a differential positive density evolution between $z\\sim0$ %and $z\\sim2$. No evidences of a negative evolution is found at high %redshifts.  %The complete samples compiled at different frequencies sample different %populations of AGN which show a positive evolution with cosmic %epoch. It naturally explains that physical processes occurring on %hundreds of kpc and on pc-scales are interconnected. The high-$z$ %cut-off is found to be significant only at high frequencies, 2.7 and 15 %GHz. %{\\tt  More discussion?}  %%   \\renewcommand{\\labelenumi}{(\\roman{enumi})} %   \\begin{enumerate} \\item The distribution of 133 AGN is uniform on %   the sky.  \\item The flat-spectrum quasars are distributed %   non-uniformly in the space: the $V_{\\rm e}/V_{\\rm a}$ distribution %   is biased towards large redshifts with $\\langle V_{\\rm e}/V_{\\rm a} %   \\rangle = 0.59\\pm0.03$, which is indicative of a positive %   evolution with cosmic epoch.  \\item The banded $V/V_{\\rm max}$ test %   shows that (a) the strong positive evolution of flat-spectrum %   quasars is significant at low-redshifts, $z<0.5$, and a %   statistically significant redshift cut-off is present at $z \\geq %   1.7$ with negative evolution up to high redshifts, $z\\sim2.5$;  %      No significant evolution is found between $0.5<z<1.7$. %      (b) the low-/high- luminosity quasars evolve differently with %      redshift which is indicative that the cosmic evolution depends %      strongly on luminosity. %   \\end{enumerate}    %{\\textbf {[To discuss? weather the high- and low-luminosity AGN %are two distinct populations of AGN]}}"
    },
    "0608/hep-th0608220_arXiv.txt": {
        "abstract": "Quantum fluctuation of unstable modes about gravitational instantons causes the instability of flat space at finite temperature, leading to the spontaneous process of nucleating quantum black holes. The energy-density of quantum black holes, depending on the initial temperature, gives the cosmological term, which naturally accounts for the inflationary phase of Early Universe. The reheating phase is attributed to the Hawking radiation and annihilation of these quantum black holes. Then, the radiation energy-density dominates over the energy-density of quantum black holes, the Universe started the Standard cosmology phase. In this phase the energy-density of quantum black holes depends on the reheating temperature. It asymptotically approaches to the cosmological constant in matter domination phase, consistently with current observations. ",
        "introduction": "In recent years, as observational data concerning on cosmology are rapidly accurate, the theoretical understanding of our Universe has been greatly profounded ever before. The inflation \\cite{linde} in early Universe and the acceleration of present Universe \\cite{wein} are two most important issues in modern cosmology and fundamental physics. Both issues are closely related to the cosmological term in Einstein equation. The observational value of the cosmological constant \\begin{equation} \\Lambda^{\\rm obs} = 8\\pi G \\epsilon_\\Lambda \\sim 10^{-118}m_p^2, \\label{lambdaobs} \\end{equation} where the Newton constant $G=1/m_p^2$, the Planck mass $m_p\\sim 10^{19}$GeV and a related characteristic energy-density \\begin{equation} \\epsilon_\\Lambda \\sim 10^{-120}m_p^4. \\label{edobs} \\end{equation} This energy-density is of $120$ order smaller than the Planck energy-density, $\\epsilon_{\\rm Planck} = m_p^4$.  Much effort in understanding these issues has been made for many decades and there are many interesting ideas and innovative theoretical developments based on either simple models \\cite{q} or complex theories \\cite{s}. In these approaches, a scalar field slowly-rolling downwards an effective potential, which mimics a possible vacuum-energy variation, plays a crucial r\\^ole in driving inflation; whereas this scalar field with a mass of order the current Hubble scale, can possibly account for acceleration.  In this article, we attempt to study these two issues within a framework based on the instability of flat space at finite temperature, due to an unstable quantum fluctuation about gravitational instantons. We shall use the signature $(- + + +)$. ",
        "conclusions": ""
    },
    "0608/astro-ph0608421_arXiv.txt": {
        "abstract": "We employ the one-dimensional Fourier representation (1DFR) to analyze the 3-year \\wmap de-biased internal linear combination (DILC) map and its possible contamination by galactic foregrounds. The 1DFR is a representation of the spherical harmonic coefficients for each $\\l$--mode using an inverse Fourier transform into one-dimensional curves. Based on the {\\it a priori} assumption that the CMB signal should be statistically independent of, and consequently have no significant correlation with, any foregrounds, we cross-correlate the 1DFR curves of $2 \\le \\l \\le 10$ modes, which are claimed by the \\wmap team to be free of contamination and suitable for whole sky analysis. We find that 8 out of the 9 modes are negatively cross correlated with the foreground maps, an event which has a probability of only $9/512 \\simeq 0.0176$ for uncorrelated signals. Furthermore, the local extrema of the 1DFR curves between the DILC and those of the foregrounds for $\\l=2$ and 6 are correlated with significance level below 0.04. We also discuss the minimum variance optimization method and use the properties of the measured cross-correlation to estimate the possible level of contamination present in the DILC map. ",
        "introduction": "The cosmic microwave background (CMB) radiation contains a wealth of information about our Universe. Not only does the angular power spectrum of the CMB temperature fluctuations allow us to determine cosmological parameters that shape our understanding of the Universe, but the issue of primordial Gaussianity and statistical isotropy of the CMB also taps the most fundamental principle in cosmology, and any robust detection of a  deviation would have far-reaching and indeed revolutionary implications.  Interpreting the term ``Gaussianity'' in the most relevant way for cosmology, i.e. as meaning that the fluctuations form a statistically homogeneous and isotropic Gaussian random field \\cite{bbks,be}, the issue of non-Gaussianity in the CMB data was first raised by \\cite{cobeng} in the COBE data \\cite{cobe}. After the release of the 1-year Wilkinson Microwave Anisotropy Probe (\\wmap) data \\cite{wmapresults,wmapfg,wmapng,wmapcl}, departures from Gaussianity have been detected with various methods \\cite{wmaptacng,gaztanagz,coleskuiper,park,eriksenmf,santanderng,romanng,hansen,mukherjee,larson,phaserandomwalk,mcewen,edingburgh,mnn}. Although Gaussianity (as we define it) requires statistical isotropy, a special focus and results on the breaking of large-scale statistical isotropy have been reported in \\cite{toh,dtzh,eriksenasym,copi,schwarz,evil,bernui,abramolow}. Among them, the most notable anomaly involves the alignment of the quadrupole and octupole alignment in the general direction of Virgo. In the 3-year \\wmap data \\cite{wmap3ytem,wmap3ycos} the previously detected departures from Gaussianity and statistical isotropy still persist \\cite{wmap3yrng,cruz,mcewenwmap3,copi3y}. The question is whether they are attributed to foreground residuals \\cite{abramo}, systematic errors \\cite{dipolestraylight}, local structures \\cite{void} or even new physics \\cite{jaffebianchi,gordon}.  However, it is important to eliminate the least exciting possibilities before getting carried away by the most exotic ones. In particular, one has to be cautious about foreground residuals as they are surely present in the derived CMB map. The elimination of foregrounds and the extraction of CMB signal is based on the concept of minimum variance optimization with the {\\it a priori} assumption that CMB and the foregrounds are statistically independent: linearly combining all available maps and minimizing the variance of the combined map to reduce the foreground residuals as much as possible. The 3-year de-biased internal linear combination (DILC) map is based on such a concept by the \\wmap science team. This approach can certainly help to achieve the optimal estimate of the angular power spectrum for the CMB. It cannot, however, guarantee the optimal morphology (spatial distribution). Since most of the anomalies, with the exception of the reported low quadrupole power \\cite{efstathiouquadrupole,dtzh,efstathiou}, are related to the pattern of fluctuations rather than simply their amplitude, one needs to check the derived CMB map very carefully before any scientific conclusion is reached. In this light, several authors have investigated possible foreground contaminations \\cite{foreground,faraday,template,2n} and the effect of different subtraction methods \\cite{dt,park3y}.  The {\\it a priori} assumption that the CMB and the foregrounds are statistically independent and thus have no significant cross-correlation is not only the backbone of the minimum variance optimization method, it is also the most basic principle (preceding Gaussianity of the CMB), furnishing the most fundamental statistical test for foreground cleaning. Note that statistical independence involves an ensemble of universes, there will, therefore, inevitably be accidental cross-correlations caused by chance alignments, particularly on large angular scales due to the cosmic variance effect \\cite{cosmicvariance}.  Efforts on the investigation on large-scale foreground contamination by direct cross-correlation were made in the harmonic domain. Based on the close connection between the phases of the spherical harmonic modes, and morphology \\cite{morph}, \\cite{ndv03,ndv04,autox} examine foreground residuals by cross-correlation of phases between the \\wmap internal linear combination maps and the derived foreground maps.   In this paper, we use a new representation of the spherical harmonic modes to analyze the cross correlation between the DILC and foreground maps. Instead of using the the two-dimensional spheres, we use an inverse Fourier transform on the spherical harmonic modes for each $\\l$ producing one-dimensional curves for each harmonic scale. We therefore call such representation one-dimensional Fourier representation (1DFR). Chiang and Naselsky \\cite{autox} first devised the 1DFR to illustrate connection between phase coupling and morphology, particularly relating to local extrema higher than 3$\\sigma$. Chiang, Naselsky and Coles \\cite{wmap3yrng} use the same representation to demonstrate the anomalies in the distribution of global extrema of the 1DFR curves of the DILC map for the $\\l \\le 10$ modes.  This paper is arranged as follows. In Section 2 we summarize the Gaussian random hypothesis of the CMB. In Section 3 we introduce the 1DFR and its properties and advantages. We then use simple cross-correlation and extremum correlation to examine foreground contamination in the DILC map in Section 4. We discuss the minimum variance optimization method and use cross-correlation coefficients to estimate the level of foreground contamination in Section 5. The conclusions are presented in Section 6. ",
        "conclusions": "In this paper we have proceeded from the basic assumption that the CMB should have no or little correlation (within cosmic variance limit) with the foregrounds, to examine foreground contamination in the \\wmap DILC map. We use the newly-developed representation 1DFR of the spherical harmonic coefficients for cross and peak correlation. The 1DFR has an advantage in that the $m=0$ modes do not alter the morphology, and when using the standard definition of cross correlation of the 1D curves, they are intrinsically excluded in the calculation.  We have tested the DILC map for $2 \\le \\l \\le 10$, although any diagnostics of correlation can be applied to all $\\l$ modes for both the \\wmap and the upcoming \\planck data. We find that eight out of nine modes are negatively-correlated, for which the probability is only 0.0176 if they are really independent. We also find the local extrema of the 1DFR curves are correlated, particularly for $\\l=2$ and 6 modes, with significance level less than $0.04$. Both analyses have indicated that the DILC quadrupole is contaminated with foregrounds (with cross correlation coefficient $X< -0.8$ and extremum statistics at significance level $\\alpha<0.04$). We then used the cross correlation coefficients to explore the level of the foreground contamination.  The axiom that the CMB has no correlation with the foregrounds provides the most fundamental criterion that should be applied as a foreground cleaning check before any Gaussianity tests are performed and cosmologists get carried away by suggestions of new physics. We suggest that for the upcoming \\planck mission, one should not ignore the kind of information we have presented here, before any scientific conclusion is reached.  One final remark for our results is that the DILC map is derived from the concept of minimum variance optimization under the {\\it a priori} assumption that CMB and the foregrounds are statistically independent. That we have only one universe, hence there is inevitably accidental correlation shall introduce some error in the estimation of the weighting coefficients, and consequently introduce residuals in the resultant map. Even if the correlation between the CMB and the foregrounds happens to be zero, the variance of the foreground residuals is minimum, but not necessarily zero. Although we can extract the angular power spectrum as close as possible to that of the CMB using such methods, the signal obtained in general still contains residuals.    \\ack PC would like to thank the Niels Bohr International Academy for their support during the Summer Institute when this paper was written. We acknowledge the use of the NASA Legacy Archive for extracting the \\wmap data. We also acknowledge the use of \\healpix \\footnote{\\tt http://www.eso.org/science/healpix/} package \\cite{healpix} to produce $\\alm$ from the \\wmap data and the use of \\glesp \\footnote{\\tt http://www.glesp.nbi.dk} package \\cite{glesp}.  \\newcommand{\\combib}[3]{\\bibitem{#3}}       %  \\newcommand{\\autetal}[2]{{#1\\ #2 \\etal}}    % \\newcommand{\\aut}[2]{{#1\\ #2}}              % \\newcommand{\\saut}[2]{{#1\\ #2},}             % \\newcommand{\\laut}[2]{{and #1\\ #2},}         %  \\newcommand{\\refsd}[6]{{\\it #1}, #5 {\\it #2} {\\bf #3} #4 [#6]} % \\newcommand{\\unrefs}[6]{{\\it #1}, #5 {\\it #2} #3 [#6]}         %   \\newcommand{\\book}[6]{#5, #1, #2}                     %   \\newcommand{\\proceeding}[6]{#5 {\\it Proc. of #4} (#2) [#6]}   %    \\def\\apj{\\it Astrophys. J} \\def\\apjl{\\it Astrophys. J. Lett.} \\def\\apjs{\\it Astrophys. J. Supp.} \\def\\ijmpd{\\it Int. J. Mod. Phys. D} \\def\\mn{\\it Mon. Not. R. Astron. Soc.} \\def\\nature{\\it Nature} \\def\\aa{\\it Astro. \\& Astrophys.} \\def\\prl{\\it Phys.\\ Rev.\\ Lett.} \\def\\prd{\\it Phys.\\ Rev.\\ D} \\def\\pr{\\it Phys.\\ Rep.}   \\def\\cup{Cambridge University Press, Cambridge, UK} \\def\\princetonpress{Princeton University Press} \\def\\worldpress{World Scientific, Singapore} \\def\\oxfordpress{Oxford University Press}"
    },
    "0608/astro-ph0608617_arXiv.txt": {
        "abstract": "% We study the influence of the frictional heating on the wind line-profiles of SMC stars. For this purpose we use our NLTE wind code to obtain consistent occupation numbers of studied levels and our radiative transfer code to solve the radiative transfer equation in moving media. We compare predicted wind line profiles calculated with and without frictional heating for a~low-luminosity SMC star with a weak wind and discuss the relevance of frictional heating as a solution of a \"weak-wind problem\" for this star. ",
        "introduction": "For lower luminosity hot stars the predicted mass-loss rates \\mdot\\ are significantly higher than those inferred from observations \\citep[``weak-wind problem'', e.g.][see also Fig.~\\ref{dmdt}]{martin}. Wind structure of these stars may be modified by the effects connected with multicomponent nature of wind flow (like frictional heating). Here we discuss influence of frictional heating on wind line profiles.  \\begin{figure}[h!] {\\centering% \\resizebox{0.36\\hsize}{!}{\\includegraphics{krticka2_fig1.eps}}\\\\} \\caption{Comparison of predicted \\mdot\\ and \\mdot\\ derived from observations.} \\label{dmdt} \\end{figure} ",
        "conclusions": ""
    },
    "0608/astro-ph0608567_arXiv.txt": {
        "abstract": "Forthcoming advances in  direct gravitational wave detection from kilohertz to nanohertz frequencies have unique capabilities to detect signatures from or set meaningful constraints on a wide range of new cosmological phenomena and new fundamental physics. A brief survey is presented of the post-inflationary  gravitational radiation backgrounds predicted in  cosmologies that include intense new classical sources such as first-order phase transitions, late-ending inflation, and dynamically active mesoscopic extra dimensions.  LISA will provide the most sensitive direct probes of such phenomena  near  TeV energies or Terascale. LISA will also deeply probe the broadband background, and possibly bursts, from  loops of cosmic superstrings predicted to form in current models of brane inflation. ",
        "introduction": "In recent years, cosmology has converged on a ``concordance'' model that continues to meet with remarkable success wherever new observational tests are introduced. The direct data span a large dynamic range of structural scales, including power spectrum probes over about a factor of $10^4$ in linear comoving scale.  The microwave background spectrum  and the products of cosmological nucleosynthesis confirm basic model assumptions and microphysics back to the first fraction of a second of the Big Bang.  An important part of the standard scenario is inflation.  The generic inflationary  predictions of  near global flatness and nearly scale-invariant, Gaussian fluctuations are well confirmed by the data. We thus seem to have some direct information, on very large comoving scales, dating from some very early time, prior to the Terascale (or TeV energies), when the Hubble length was less than about a millimeter and the Hubble time less than a few picoseconds.  At the same time, advances in string theory have opened up a wide range of new possibilities for the initial conditions and the behavior of the universe at early times.  Seven extra dimensions seem mandatory for consistency; they must be  compactified, or gravity must be localized by some other means;  depending on the order of events, these possibilities may require extensions to the usual framework of inflation based on field theory in 3+1 dimensions. Initial conditions, and even apparently fundamental parameters such as the shapes of effective potentials in the field theory Lagrangian,   may be strongly shaped by anthropic selection from a string-theory ``landscape'', changing  the character of arguments about what is natural for inflationary scenarios.  Branes of various dimensionalities  introduce new behavior not present  in a system of fields in 3D space, including  new possibilities for inflation itself, for the generation of fluctuations, and for new kinds of macroscopic defects such as cosmic superstrings.  Such a large range of new possibilities can be introduced because the observational tests of concordance cosmology cannot directly probe a very large swathe of early cosmic history. Viewed in terms of logarithmic scales in energy space or cosmic scale factor, the unexplored region of our past light cone greatly  exceeds the explored region. For example, information from particle content only survives from the decoupling of the relevant species; for most dark matter candidates, this is well below a GeV in temperature. The baryon content is a presumed relic of the pre-TeV period, which sets the only firm lower limit on the end of inflation. And while it is likely true that the largest-scale perturbations date from very early,  many (possibly as many as about 27)  orders of magnitude of expansion before the end of  inflation, the direct data on these cover only  a relatively brief period and a few orders of magnitude of comoving scale. We have almost no direct data on what happened during about 30 orders of magnitude of cosmic expansion; for almost all of this period, we do not even have solid theoretical arguments to believe that we understand the  dimensionality of space.  Advances in detection of high frequency gravitational waves are starting to provide a way to explore  observationally this very large range of scales, epochs and phenomena not previously accessible. A new generation of ground-based detectors operating in the  sub-kilohertz band\\cite{LIGOURL,VIRGOURL}, such as LIGO, GEO, and VIRGO, are now collecting data at their design sensitivity, and will be significantly upgraded in the years ahead.  The space interferometer\\cite{LISAURL}  LISA, operating in the millihertz band, will have its core technology proven in a test flight in 2009 and will itself fly in about ten years. Pulsar surveys are providing larger samples of clean pulsars for timing-based detection at frequencies of inverse years to decades\\cite{Kaspi:1994hp,Thorsett:1996dr,McHugh:1996hd,Lommen:2001ax,Lommen:2002je}. Together, these techniques are maturing to the point where we will have cosmologically sensitive information over a range of frequencies from  kilohertz to nanohertz. Since gravitational waves  penetrate all of cosmic history at sub-Planckian densities, and even propagate some distance into extra dimensions, these techniques will access directly many phenomena that are currently hidden from view.  This brief survey represents an introduction and  update; more detail on many of the ideas presented here appear in earlier reviews and surveys\\cite{Hogan:1998dv,Maggiore:1999vm,Cutler:2002me,Hughes:2002yy,Buonanno:2003th,Chongchitnan:2006pe}. Also, since the emphasis is on classically-generated waves that might be directly detected in the near future (and not on cosmic background techniques), inflationary quantum perturbations are little discussed, even though they dominate calculations in the cosmological literature;  direct detection of these waves might occur in the more distant future\\cite{Efstathiou:2006ak,Boyle:2005se,Smith:2005mm,Crowder:2005nr}. ",
        "conclusions": ""
    },
    "0608/astro-ph0608084_arXiv.txt": {
        "abstract": "In this paper we report the discovery of a high energy component of the X-ray spectra of U Gem , which can be observed while the source is in outburst. We used Chandra and XMM-Newton observations to compare the quiescence and outburst X-ray spectra of the source. The additional component may be the result of the reflection of X-rays emitted from an optically thin plasma close to the white dwarf, from the optically thick boundary layer during the outburst. Another possible explanation is that some magnetically channeled accretion may occur onto the equatorial belt of the primary causing shocks similar to the ones in the intermediate polars as it was suggested by \\citep{w2002}. We have also found a timing structure at about 73 mHz ($\\sim$13.7 s.) in the RXTE observation, resembling dwarf novae oscillations (DNOs). ",
        "introduction": "Dwarf novae are highly variable mass-exchanging binary star systems containing a white dwarf and a late-type K or M star. They are characterized by occasional outburst episodes which typically take place on a timescale of 100 days and last for about 15 days. Outbursts are believed to be triggered by a thermal instability in the accretion disk that drives the disk from a low-temperature, low mass accretion rate, quiescent state to a hot, high-$\\dot{\\rm{m}}$ outburst state.  About every 120 days prototypical dwarf nova U Gem exhibits outbursts, where the V magnitude of the system rises from $\\sim14$ to $\\sim9$ \\citep{s1984}. During outbursts U Gem is a bright EUV source. This emission is interpreted as optically thick radiation from a $\\sim140.000 \\rm{K}$ boundary layer \\citep{l1996}. Besides, there is observational evidence that the observed energy from the boundary layer and that from the disk is comparable \\citep{l1996} which is in accordance with the simple mass accretion scenario \\citep{p1977}. In quiescence the UV spectra of U Gem is dominated by the white dwarf emission as observed with IUE \\citep{k1991}, Hopkins Ultraviolet Telescope (HUT) \\citep{l1993} and Hubble Space Telescope (HST) \\citep{l1994}. Studies of the white dwarf in U Gem indicate evidence for an equatorial accretion belt \\shortcite{c1997}, \\shortcite{l1993} which is cooling during the quiescence interval.  X-ray emission (0.2 - 10 keV) has been observed from U Gem both in quiescence and outburst \\shortcite{s1978}, \\shortcite{c1984}, \\shortcite{s1996}. Unlike other well known dwarf novae systems, like SS Cygni, hard X-ray flux of U Gem does not decrease in outbursts \\citep{m2000}. The fact is that the increase in the hard X-ray (2-15 keV) flux is less than that of the soft X-ray flux (0.1-4 keV) \\shortcite{c1984}, \\shortcite{m2000}. Chandra observations of U Gem in quiescence indicate that the hard X-ray emission arises from a gas with a small scale height ($<10^{7}cm$) close to the white dwarf \\citep{s2002}, (hereafter S02).  Here we present detailed hard X-ray emission properties of U Gem in outburst phase for the first time. We find an extra non-thermal X-ray component arising during outburst which can be intrepreted as a sign of a transient magnetosphere occuring in the boundary layer of U Gem as a result of the increase of the mass accretion rate \\citep{w2002}.  In section 2, we present observations and our data analysis. In sections 3 and 4, we show observational differences between quiescence and outburst phases, respectively. In section 5, we discuss these results. ",
        "conclusions": "In this work we have investigated U Gem's X-ray behaviour during outburst. We found that unlike quiescence, one cannot fit the X-ray spectrum of U Gem with a simple mekal based multi-temperature spectral model. Instead, an additional non-thermal component, most likely a power-law, is needed. We should also note that when fitting the excess hard X-ray emission with a power-law, deviations from the continuum can be seen (see Figures \\ref{subtracted} and \\ref{outburst_spectra}). We think that these deviations arise because of the broadening and increase of the fluxes of the X-ray emission lines when compared to the quiescence phase. In order to be able to obtain better fits to the outburst data one should also be able to model the broadening of the lines physically, which is currently unavailable.  Interpretation of the extra X-ray component can be made as follows: Since this component was observed only during outburst phase, there must be a mechanism which is temporarily available in outbursts. During quiescence, the mass accretion rate is not enough to create an optically thick boundary layer, but in outbursts there is observational evidence for such a boundary layer \\citep{l1996}. Comparing the X-ray spectra obtained in quiescence and outburst we see that the optically thin plasma observed in quiescence, can also be seen in outburst with approximately the same temperature, but there exists an additional spectral component. Reflection of this X-ray emitting plasma from the optically thick boundary layer might occur in outburst giving rise to the extra X-ray component. Such reflection models are applied to AGNs, furthermore a similar detection for reflection during an outburst was reported by \\shortcite{d1997} for SS Cygni. According to the reflection scenario, besides this extra component, one would expect to see $K_{\\alpha}$ fluorescent lines \\shortcite{b1995} and \\shortcite{d1995}, due to the reflection of hard X-rays from the relatively cold EUV emitting boundary layer. Unfortunately in the high resolution Chandra spectra of U Gem during outburst those lines are either very weak or absent especially when compared to the resonance lines of Fe appearent near 1.86 \\AA. Moreover such reflection components are most effectively seen as a bump around 20 keV \\citep{d1997}.  Another plausible explanation for this hard X-ray component arises with the help of Low Inertia Magnetic Accretor model (LIMA) \\citep{w2002}. According to the model, during outbursts the spin up of an accretion belt of the weakly magnetic white dwarf will enhance the intrinsic magnetic field and with this way, magnetically channelled accretion can occur even on low field CV primaries. With such an accretion onto the equatorial belt of the primary, accretion curtains and shocks in the same manner as in the standard intermediate polar structure are possible \\citep{w2002}. This model is useful for explaining the dwarf nova oscillations (DNOs) observed in outbursts. In the model these DNOs are interpreted as the pulsations at the rotation period of this transient magnetosphere.  As we have noted above, an equatorial accretion belt has been deduced from UV observations \\citep{c1997}, \\citep{l1993}. So if accretion curtains and shocks are likely to occur in U Gem during outbursts then a thermal bremsstrahlung or a power-law like emission may be observed from the cooling post-shock gas showing us an extra X-ray component. From the Chandra observations in quiescence (S02) and HST observations during outburst \\citep{s1997} it has been suggested that magnetic accretion is possible for U Gem.  We find a significant structure in one power spectrum that is similar to quasi periodic oscillations (QPO). This feature appears around 73 mHz ($\\sim 13.7$ s.). Transient DNOs have been observed for U Gem, during outbursts, in the 25 s. range, by EXOSAT \\citep{m1988} and EUVE \\citep{l1996}. These oscillations may be arising from the transient magnetosphere which is predicted by the LIMA model. We, therefore, suggest that the extra X-ray spectral component as well as transient QPO we detected originates from a transient magnetosphere.  Another issue is the appearent changes in the emission line profiles, whose some of the properties are summarized in Table \\ref{emis_line}. Increase of the fluxes can be understood by assuming an high mass transfer to the inner disc during outuburst. This will increase the probability of collisions between particles, which will of course increase collisional ionization. However, explanation of the broadening of the lines can not be made as easy as increased fluxes, one assumption we can make is that the emission lines are now coming from a region rotating with Keplerian velocities, which influences the idea that the matter, at the outer parts of the accretion disc, moving towards to the white dwarf during the outburst. Although this idea is mainly in agreement with the standard outburst theories \\citep{w1995}, it is not obvious from only one observation, because there are still some lines which are not as broadened as expected to form in a region rotating with a Keplerian velocity.  In order to understand this phenomenon further broadband observations of U Gem and as well as other CVs, are needed to investigate the overall variations, particularly in the boundary layer, between the quiescence and outburst phases."
    },
    "0608/astro-ph0608037_arXiv.txt": {
        "abstract": "We calculate the scattering of X-rays by interstellar dust, for a dust model that reproduces the observed wavelength-dependent extinction and polarization of starlight. On interstellar sightlines that produce appreciable starlight polarization, we predict that the dust-scattered X-ray halo around point sources will have measurable azimuthal asymmetry due to scattering by partially-aligned nonspherical grains. We calculate the expected halo asymmetry. X-ray halo asymmetry provides a new test of interstellar dust models. ",
        "introduction": "}  Interstellar dust grains scatter X-rays through small angles, as was first pointed out by Overbeck (1965), Slysh (1969), and Hayakawa (1970). Because of this scattering, an image of an X-ray point source includes a ``halo'' of X-rays that have been scattered by dust grains near the line of sight. First observed by Catura (1983) using the {\\it Einstein} observatory, scattered X-ray halos have since been measured by a number of telescopes, including {\\it Einstein} (e.g., Mauche \\& Gorenstein 1986), {\\it ROSAT} (e.g., Predehl \\& Schmitt 1995), {\\it Chandra} (e.g., Smith, Edgar, \\& Shafer 2002), and {\\it XMM-Newton} (e.g., Costantini, Freyberg, \\& Predehl 2005).  The observed polarization of starlight (Hall 1949; Hiltner 1949) requires that interstellar grains be both appreciably nonspherical and partially aligned, with starlight propagating through the dusty interstellar medium becoming linearly polarized as the result of preferential attenuation of one of the linear polarization modes (``linear dichroism''). In addition to producing polarization of starlight, aligned dust grains produce polarized thermal emission at wavelengths from the far infrared to the microwave.  Observers seeking to measure the polarization of the cosmic microwave background must subtract this polarized ``galactic foreground'' from observations of the microwave sky. Realistic models of nonspherical interstellar grains are therefore of interest for many reasons.  There continue to be many uncertainties concerning the composition and geometry of interstellar grains (for a recent review, see Draine 2003a), and X-ray absorption and scattering can be used to test grain models. We consider a specific grain model consisting of spherical carbonaceous grains and oblate spheroidal silicate grains, with size distributions and size-dependent degree of alignment adjusted to reproduce the observed wavelength dependence of both interstellar extinction and interstellar polarization.  The scattering ``halo'' produced by a nonspherical grain will not be azimuthally symmetric, and therefore the population of aligned interstellar grains can be expected to produce asymmetric X-ray scattering halos. Here we develop a method for calculating the differential scattering cross section for X-rays incident on grains with arbitrary geometry, and apply it to calculate X-ray scattering halos for a realistic model of interstellar dust.  We find that the X-ray scattering halo produced by this model of aligned intersellar grains has appreciable and observable asymmetry. We propose statistics $R_\\ell^{(I)}$ to measure the asymmetry of observed halos. For two models of partially-aligned interstellar grains, we calculate the expected values of $R_2^{(I)}$ and $R_4^{(I)}$, and discuss their observability. The predicted values of $R_2^{(I)}$ should be measurable on sightlines where there is appreciable polarization of starlight. The ratio $R_4^{(I)}/R_2^{(I)}$ is sensitive to the shape of the scattering grains.  The paper is organized as follows: In \\S\\ref{sec:x-ray scattering} we discuss anomalous diffraction theory, as applied to interstellar grains, and our implementation of it. In \\S\\ref{sec:spinning precessing grains} we discuss averaging over the grain rotation expected for partially-aligned suprathermally-rotating grains. In \\S\\ref{sec:azimuthal asymmetry} we introduce statistics $R_\\ell^{(I)}$ to quantify azimuthal asymmetries in scattering halos. We obtain a realistic grain model in \\S\\ref{sec:grain model}, and in \\S\\ref{sec:anisotropic scattering by aligned dust} we use this model to predict $R_2^{(I)}$ and $R_4^{(I)}$ for sightlines where the magnetic field direction is uniform and perpendicular to the line-of-sight. We discuss using halo asymmetry measurements as a test of grain models in \\S\\ref{sec:discussion}. Our results are summarized in \\S\\ref{sec:summary}.  A reader concerned only with the predicted observability of this phenomenon may choose to proceed directly to \\S\\ref{sec:discussion}. ",
        "conclusions": "Discussion} \\subsection{Observational Considerations}  We will assume that there is no uncertainty regarding the position of the point source on the sky; this position is taken to be the origin of a polar coordinate system $(\\theta,\\phi_\\obs)$ with the direction of $\\phi_\\obs=0$ taken to be whatever coordinate system is convenient to the observer. We further assume that the image $I(\\theta,\\phi_\\obs)$ has had the instrumental point spread function (p.s.f.) subtracted; this is especially important if the p.s.f.\\ is not azimuthally symmetric. Real images may contain extraneous background or foreground sources that need to be recognized and removed; we assume that standard methods are used to interpolate the scattered intensity $I(\\theta,\\phi_\\obs)$ in such regions.  In the following discussion it is assumed that $I(\\theta,\\phi_\\obs)$ is the observed scattered intensity.  An annulus $[\\theta_1, \\theta_2]$ is chosen that contains a strong scattered signal, with $\\theta_1$ chosen to avoid the central core that will be dominated by the instrumental p.s.f., and $\\theta_2$ chosen to ideally extend beyond the radius where the scattered halo intensity begins to significantly decline. If $I(\\theta,\\phi_\\obs)$ is the observed intensity (corrected for nonuniform backgrounds), a new coordinate system $\\phi\\equiv\\phi_\\obs-\\Delta$ is defined, with the rotation angle $\\Delta$ determined by the condition \\beqa \\label{eq:condition for phi=0} \\Delta &=& \\frac{1}{2}\\arctan(A/B) ~~~, \\\\ A &\\equiv& \\int_{\\theta_1}^{\\theta_2} d\\theta\\sin\\theta \\int_0^{2\\pi}d\\phi_\\obs I(\\theta,\\phi_\\obs)\\sin(2\\phi_\\obs) ~~~, \\\\ B &\\equiv& \\int_{\\theta_1}^{\\theta_2} d\\theta\\sin\\theta \\int_0^{2\\pi}d\\phi_\\obs I(\\theta,\\phi_\\obs)\\cos(2\\phi_\\obs) ~~~. \\\\ \\eeqa In this coordinate system $(\\theta,\\phi)$, the anisotropy statistics $R_\\ell^{(I)}(\\theta_1,\\theta_2)$ can be calculated from eq.\\ (\\ref{eq:R_ell^(I)}). The resulting $R_2^{(I)}(\\theta_1,\\theta_2)$ and $R_4^{(I)}(\\theta_1,\\theta_2)$ can be compared to the theoretical predictions in Figure \\ref{fig:R_2,R_4: 0.5,1,2keV; mix; sheet; b/a=1.4,2} (for dust at a single distance) or Figure \\ref{fig:R_2,R_4: E=0.5,1,2keV; mix; uniform; b/a=1.4,2} (for uniformly-distributed dust).  We have calculated the anisotropy statistics $R_\\ell^{(I)}$ for plane-parallel distributions of dust: dust in a sheet at a single distance and dust uniformly distributed between observer and source. In real situations, gradients in the dust density distribution perpendicular to the line-of-sight (on linear scales of order $\\theta_{\\rm halo}\\approx 1\\arcmin\\times1\\kpc\\approx0.3\\pc$) can contribute to the $R_\\ell^{(I)}$. How can halo anisotropies due to aligned dust grains be separated from anisotropies due to dust density gradients? \\begin{enumerate} \\item The halo anisotropy due to dust grain alignment will have $R_1^{(I)}=R_3^{(I)}=0$. Dust density gradients, on the other hand, will generally contribute to all of the $R_\\ell^{(I)}$, and in fact would be expected to make the largest contribution to $R_1^{(I)}$. Therefore the magnitude of $R_1^{(I)}$ and $R_3^{(I)}$ can be used to estimate the contribution of dust density gradients to $R_2^{(I)}$. \\item The halo anisotropy due to aligned dust grains will be aligned with the direction of starlight polarization: polarized starlight has ${\\bf E}$ parallel to the short axis of the aligned grains.  Therefore the scattered halo will have its major axis parallel to the direction of starlight polarization: the angle $\\Delta$ should coincide with the direction of starlight polarization. \\item If the anisotropic signal $R_2^{(I)}$ is from aligned grains, the angle $\\Delta$ will be insensitive to the choice of annulus $[\\theta_1,\\theta_2]$. If $\\Delta$ is found to vary significantly from one annulus to another, the observed anisotropy will have an appreciable component from some source other than aligned dust grains. -- nonuniform dust, or other X-ray sources in the field. \\end{enumerate}  Ideally, the X-ray scattering halo would be observed for a source where the optical starlight polarization is known either for the source itself or a nearby stellar companion.  In some cases, the X-ray source may not be bright enough at optical wavelengths to permit polarization measurements.  However, the polarization of starlight is known to display large-scale coherence over the sky (see, e.g., Mathewson \\& Ford 1970); therefore if starlight polarization has been observed for a star that is nearby on the sky, with a similar amount of reddening, it is reasonable to presume that the dust in front of the candidate X-ray source would produce a similar degree of optical polarization.  \\begin{figure}[h] \\includegraphics[angle=270,width=16cm]{f18.eps} \\caption{\\footnotesize \\label{fig:polmap} Line segments show ratio direction of starlight polarization, with length proportional to $p(V)/E(B-V)$, where $p(V)$ is the polarization, and $E(B-V)$ is the reddening. Data from Heiles (2000). In this region of the sky, the dust grain alignment is evidently quite uniform, and the polarization per unit $E(B-V)$ is $\\sim2/3$ of the empirical maximum. } \\end{figure}  Directions with large observed starlight polarization tend to lie close to the galactic plane (so that significant amounts of reddening will be present) and in directions perpendicular to the general direction of the local magnetic field. An example of a region with large observed starlight polarization per unit $E(B-V)$ is $l = 130 \\pm 10\\deg$.  Figure \\ref{fig:polmap} shows $p(V)/E(B-V)$ for stars in a section of the Galactic plane near $l\\approx 130\\deg$. The stars shown are limited to $0.30 \\leq E(B-V) \\leq 0.60$~mag (i.e., $0.9 \\ltsim A_V \\ltsim 1.9$~mag). The large-scale coherence of the polarization direction is apparent, and it is also evident that the ratio of polarization to reddening is relatively uniform.  This indicates that the magnetic field is relatively well-ordered, and that the efficiency of grain alignment is also relatively uniform.  An X-ray source in this general region would presumably be observed through the same aligned dust as the stars for which the polarization has already been measured.  \\subsection{Predictions}  We have considered two models for interstellar grains, both with partially aligned silicate oblate spheroids, but for two different axial ratios.  The aligned grains that produce polarization of starlight also produce anisotropic X-ray scattering. While we have explicitly calculated the anisotropic X-ray scattering only for the optimal case where the magnetic field $\\bB$ is unidirectional and perpendicular to the line-of-sight, we expect that for other magnetic field geometries, the anisotropy $R_2^{(I)}$ will decrease in proportion to $\\pmax/E(B-V)$, because for optically-thin conditions both $R_2^{(I)}$ and $\\pmax/E(B-V)$ respond similarly to changes in the magnetic field direction.\\footnote{ E.g., both are unchanged if $\\bB\\rightarrow-\\bB$; both vanish if $\\bB\\parallel\\bzhat$; and both vanish if 50\\% of the dust has symmetry axis $\\bahat_1\\parallel\\bxhat$, and 50\\% has $\\bahat_1\\parallel\\byhat$. } We can therefore predict the ratio of the fractional anisotropy of the X-ray halo to the ratio of the optical polarization to total reddening.  We suppose that it is possible to measure the X-ray halo for $\\theta_1<\\theta<\\theta_2$, with $\\theta_1\\ltsim 200\\arcsec$, and $\\theta_2 \\gtsim 300\\arcsec$. If the dust is distributed approximately uniformly between observer and source, from Figure \\ref{fig:R_2,R_4: E=0.5,1,2keV; mix; uniform; b/a=1.4,2} we see that for X-ray energy $E = 1\\keV$, \\beq \\label{eq:result for R2/(pmax/E(B-V))} R_2^{(I)}(\\theta_1,\\theta_2) \\gtsim 0.09 \\times \\left[ \\frac{\\pmax}{0.09E(B-V)/{\\rm mag}} \\right] ~~~, \\eeq provided the annulus has $\\theta_1\\ltsim 300\\arcsec$, and $\\theta_2\\gtsim 300\\arcsec$. The two cases $b/a=\\sqrt{2}$ and 2 give similar values for the quadrupole anisotropy $R_2^{(I)}$ -- increased grain oblateness is offset by reduced degree of partial alignment. As discussed above, the anisotropy $R_2^{(I)}$ is an increasing function of energy over the range 0.5--2~keV because the relative contribution of the nonspherical silicate grains to the scattering increases with energy.  It is not critical that the dust be distributed uniformly for the result (\\ref{eq:result for R2/(pmax/E(B-V))}) to be applicable, only that a negligible amount of the starlight polarization be produced by dust that is very close to the source, as the anisotropic X-ray halo produced by this dust may be largely lost at small angles unless it is possible to use small values of $\\theta_1$.  Aligned grains are also predicted to produce an octupole anisotropy in the X-ray halo, measured by $R_4^{(I)}$.  Unfortunately, this anisotropy can be washed out by moderate rotation of the magnetic field direction along the line-of-sight. However, if the starlight polarization has $\\pmax\\approx 0.09 E(B-V)/{\\rm mag}$, then we can suppose that the magnetic field in the dusty regions is approximately perpendicular to the line-of-sight and unidirectional (otherwise the net starlight polarization would be reduced).  In this case, we can use the statistic $R_4^{(I)}$ as a test of grain models.  The model with minimal axial ratio $b/a\\approx1.4$ for the spheroids predicts $R_4^{(I)}\\approx 0.03$ at 1~keV (see Figure \\ref{fig:R_2,R_4: E=0.5,1,2keV; mix; uniform; b/a=1.4,2}). If larger values of $R_4^{(I)}$ are observed, it will suggest more extreme grain shapes (for example, $R_4^{(I)}\\approx 0.10$ would be consistent with $b/a\\approx 2$ for the silicate spheroids, but inconsistent with $b/a\\approx\\sqrt{2}$). Therefore, measurement of $R_4^{(I)}$ has the potential to discriminate between grain models with different grain shape.   Summary} The principal results of this paper are as follows: \\begin{enumerate} \\item We show that anomalous diffraction theory can be used for accurate calculation of X-ray scattering from dust grains with arbitrary shape.  Fast Fourier transforms may be employed for efficient calculation of the scattering halo. \\item Differential scattering cross sections are calculated for oblate spheroids, and are shown to have substantial anisotropies if the spheroid axial ratio differs appreciably from unity. \\item Models of spherical carbonaceous grains and partially-aligned silicate spheroids are found that reproduce the observed interstellar extinction and polarization of starlight as a function of wavelength. Confirming a result found previously by Kim \\& Martin (1995), we find that for silicate axial ratio $b/a\\approx\\sqrt{2}$ the model can reproduce the largest observed starlight polarization, provided that the silicate spheroids with $\\aeff\\gtsim0.1\\micron$ are almost perfectly aligned (with their short axis parallel to the magnetic field). For larger axial ratios, the grain alignment need not be so complete. \\item X-ray scattering halos are calculated for aligned interstellar dust grains.  For realistic size distributions and fractional alignments, the scattered halo shows substantial anisotropy. \\item We propose statistics $R_\\ell^{(I)}(\\theta_1,\\theta_2)$ to measure the halo anistropy in an annulus $\\theta_1<\\theta<\\theta_2$. We predict the values of $R_2^{(I)}/[\\pmax/E(B-V)]$. We find that $R_2^{(I)}$ is large enough to be measured on sightlines to X-ray point sources where the grains are aligned so as to produce starlight polarization with $\\pmax/E(B-V)\\gtsim 0.05/{\\rm mag}$. \\item We show that the octupole anisotropy $R_4^{(I)}$ is sensitive to the assumed grain shape.  On favorable sightlines, $R_4^{(I)}$ can be used to constrain the geometry of interstellar grains. \\end{enumerate}"
    },
    "0608/astro-ph0608171_arXiv.txt": {
        "abstract": "% We review recent high-resolution CO observations of distant QSOs obtained at the Very Large Array.  The aim of these observations is to resolve the molecular gas distribution in these extreme objects both spatially and in velocity space. They provide unique information regarding the small-scale distribution, the extent, and the brightness temperatures of the molecular gas in these early systems. E.g., the structure and dynamics of the molecular gas may reveal whether or not mergers can be the cause of the ongoing starburst activity. The observations also allow for a first estimate of the dynamical gas mass. Currently, only the VLA is able to obtain resolutions in CO of up to 0.15$''$ which is needed to resolve typical galactic structures of sizes $\\sim$1 kpc. We present new high-resolution VLA imaging of high--$z$ QSOs (BRI 1335--0417, APM 08279+5255 and J1148+5251).  These observations pave the road to future ALMA observations where resolutions of order 0.1$''$ will be obtained routinely. ",
        "introduction": "Over the last few years, the study of high redshift QSOs has been revolutionized in three ways. First, wide field surveys have revealed 100's of high--$z$ QSOs, right back to the epoch of cosmic reionization ($z > 6$; e.g., Fan et al. 2006). Second, it has been shown that most (all?) low redshift spheroidal galaxies have central super-massive black holes (SMBH), and that the black hole mass correlates with bulge velocity dispersion.  This M$_{BH}$--$\\sigma_{\\rm v}$ correlation suggests coeval formation of galaxies and SMBH, thereby making SMBHs a fundamental aspect of the galaxy formation process (Gebhardt et al. 2000).  And third, mm surveys of high redshift QSOs find that 30$\\%$ of the sources are `hyper-luminous infrared galaxies' ($L_{FIR} = 10^{13}$ L$_\\odot$), corresponding to thermal emission from warm dust, and that this fraction is {\\sl independent of redshift out to $z = 6.4$} (Beelen et al. 2004, Omont et al. 2004). If the dust is heated by star formation, the implied star formation rates are extreme ($> 10^3$ M$_\\odot$ year$^{-1}$), consistent with the formation of a large elliptical galaxy on a dynamical timescale of 10$^8$ years.  \\\\ ",
        "conclusions": "The recent observations presented here have shown that resolved molecular gas {\\em maps} of objects out to the highest redshifts can already be achieved with today's instrument (i.e. the VLA). Although these observations are time intensive, they have dramatically improved our knowledge of the properties of the molecular gas phase at high redshift. E.g., they have helped to constrain the source sizes and kinematics, which in turn constrain the dynamical masses (e.g. in the case of J1148+5251, and B1335-0417). Also, they have provided clear evidence for co-eval starburst and AGN activity (e.g., J2322+1944) and revealed likely interactions with objects that are barely visible in the optical (e.g.  B1202-0725). These results show the enormous potential for future studies of the molecular gas content and dynamical masses in a statistically significant sample of high--redshift galaxies using ALMA, where resolutions of $<0.1''$ will be achieved routinely."
    },
    "0608/astro-ph0608492_arXiv.txt": {
        "abstract": "We would like to discuss prospects for neutrino observations of the core-collapse supernova progenitor during neutrino-cooled stage. We will present new theoretical results on thermal neutrino and antineutrino spectra produced deep inside the pre-supernova core. Three competing processes: pair-, photo and plasma-neutrino production, are taken into account. The results will be used to estimate signal in existing and future neutrino detectors. Chance for supernova prediction is estimated, with possible aid to core-collapse neutrino and gravitational wave detectors in the form of early warning. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608347_arXiv.txt": {
        "abstract": "The evidence for a shortage of exosolar planets with semimajor axes $-1.1 \\leq \\log (a/AU) \\leq -0.2$ is investigated. It is shown that this valley results from a gap in the radial distribution of planets, orbiting stars with masses $M_* \\geq 1.2 M_{\\odot}$ (the high-mass sample, HMS). No underabundance is found for planets orbiting stars with smaller masses. The observational data also indicate that within the HMS population it is preferentially the more massive planets with $M \\sin(i) \\geq 0.8 M_J$ that are missing. Monte-Carlo simulations of planet formation and migration are presented that reproduce the observed shortage of planets in the observed radius regime. A dependence on the disk depletion timescale $\\tau_{\\rm dep}$ is found. The gap is more pronounced for $\\tau_{\\rm dep} = 10^6 - 10^7$ yrs than for $\\tau_{\\rm dep} = 3 \\times 10^6 - 3 \\times 10^7$ yrs. This might explain the observed trend with stellar mass if disks around stars with masses $M_* \\geq 1.2 M_{\\odot}$ have shorter depletion timescales than those around less massive stars. Possible reasons for such a dependence are a decrease of disk size and an increase of stellar EUV flux with stellar mass. ",
        "introduction": "Within the past decade, nearly 200 extrasolar planets have been detected. Their orbital properties and masses and the relationships between planetary properties and those of the central stars are providing interesting insight into the complexity of planetary system formation and offer important constraints for theoretical models. Recently, extrasolar planets have been found around low-mass stars (M dwarfs) by high-precision radial-velocity (Butler et al.~2004, Bonfils et al.~2005) and microlensing surveys (Beaulieu et al.~2006, Gould et al.~2006) with indications that Jupiter-mass planets are rare around M dwarfs (Laughlin et al.~2004, Ida \\& Lin 2005).  Theoretical models have been developed to explain the statistical properties of exoplanets. Ida \\& Lin (2005) presented Monte-Carlo simulations that predict mass and semimajor axis distributions of extrasolar planets around stars with various masses. Their prescription of planet formation (Ida \\& Lin 2004a) is based on the core-accretion scenario in which it is assumed that Jupiter-mass gas-giant planets form through 1) grain condensation into kilometer-sized planetesimals, 2) runaway planetesimal coagulation (Wetherill \\& Stewart 1989; Kokubo \\& Ida 1996), 3) oligarchic growth of protoplanetary embryos (Kokubo \\& Ida 1998, 2000), and finally 4) gas accretion onto solid cores (Mizuno 1980; Bodenheimer \\& Pollack 1986; Pollack et al.~1996; Ikoma et al.~2000). The models show that the characteristic {\\it mass} distribution of planets around M dwarfs should be quite different from that around FGK dwarfs, which agrees well with the results of microlensing surveys (Beaulieu et al.~2006; Gould et al.~2006). The core accretion model also predicts a strong dependence of the likelihood for finding planets on the stellar metallicity (Ida \\& Lin 2004b; Kornet et al.~2005) which is in agreement with observations (Santos et al.~2000; Fischer \\& Valenti 2005).  Like the mass distribution, the {\\it semimajor axis} distribution of exoplanets needs to be explained by any consistent theoretical model of planet formation. One of the most puzzling questions that arose already when the first Jovian planets around main sequence stars were discovered (Mayor \\& Queloz 1995) is their wide range of orbital radii, with distances ranging from $2 \\times 10^{-2}$ AU to the observational limit of $\\sim 5$ AU. According to the core-accretion scenario, preferred formation locations of Jovian planets are the outer regions of protoplanetary disks where the temperatures are low enough for ice to survive (e.g. Ida \\& Lin 2004a). However, hot Jupiters have been observed very close to the central star ($\\la 0.1$AU). It is unlikely that these objects formed in these inner regions, although Jovian planets can form inside the ice boundary in relatively massive disks (Ida \\& Lin 2004a). Orbital migration due to gravitational interaction between the planet and the gaseous disk (Goldreich \\& Tremaine 1980; Lin \\& Papaloizou 1986a,b; Lin et al.~1996) is considered as the most promising mechanism to generate hot Jupiters. It starts with relatively fast type I migration (e.g. Ward 1986; Tanaka et al. 2002). As soon as the planet has become massive enough to open a gap in the disk it enters the phase of type II migration where the infall rate is connected with the viscous evolution of the disk. Numerical simulations have explored the complex migration process in great details (see e.g. de Val-Borro et al.~2006 for a recent comparison of different numerical simulations).  The Monte Carlo models of Ida \\& Lin (2004b) predict that the (differential) distribution for extrasolar planets with radial velocities larger than 10 m/s is an increasing linear function in log $a$ from $a \\sim 0.1$ AU up to $\\sim 1$ AU and decreases with $a$ for $a \\ga 1$AU, independent of disk metallicity, if the disk depletion timescale is comparable to the disk diffusion timescale at the radius where Jovian planets are formed (The diffusion timescale determines the migration speed). Very similar results were found by Armitage et al.~(2002) and Trilling et al.~(2002). Some planets might survive in the outer regions. Near the inner disk edge ($\\la 0.1$AU), those planets that migrate inwards from the outer regions may pile up (Lin et al.~1996). As a result, a gap (valley) would be created at intermediate semimajor axes (equivalently, intermediate periods) in the distribution. Udry et al.~(2003) indeed point out a shortage of planets in the 10-100 day period range in the observed data (see also Zucker \\& Mazeh 2002) which they interpret as a signature for a transition region between two categories of planets that suffered different migration scenarios.  In this paper we analyse again the valley in the period and separation distribution of extrasolar planets taking into account the newest set of observational data. We point out that in the observational data, the valley is much more pronounced around F dwarfs than GK dwarfs which may reflect a dependence on disk size and/or on the EUV flux from the central star. Section 2 presents the observational evidence for a gap around F dwarfs. Section 3 discusses a series of Monte-Carlo simulations that lead to a dependence of the planet semimajor axis distribution on stellar mass. A discussion of the results follows in section 4. ",
        "conclusions": "We should start the discussion with a word of caution. Despite the fact that almost 200 exoplanets are known, the statistics for more advanced correlations like the one discussed here is still poor. In addition, for small numbers spurious correlations between different variables can always be found. They would appear to have a high statistical significance which is however caused by the biased selection of the data. At the moment we cannot rule out that the planet shortage in the separation range of  0.08 AU $\\leq a \\leq 0.6$ AU, first discussed by Udry et al. (2003), is such a case.  However, {\\it if} this valley in period and separation distribution exists, we argue that it is due to a lack of planets with masses M $\\sin$(i) $>$ 0.8 M$_J$ orbiting around stars with masses M$_* \\geq$ 1.2 M$_{\\odot}$ which provides interesting new insight on the origin of massive planets orbiting close to their parent stars and the dependence on stellar properties. We point out that smaller size protoplanetary disks around more massive stars can account for a deeper valley around these stars. If MRI turbulence is responsible for the disk viscosity, $\\alpha$ values for the viscosity would not significantly depend on stellar mass. Adopting a constant $\\alpha$, smaller disk sizes around more massive stars would lead to shorter disk depletion timescales, which prevents Jovian planets formed in the outer regions from migrating considerably. Strong EUV radiation from massive (pre-main) stars may also inhibit the migration across the regions at a few AU. Another possibility is a larger radius of the ice boundary for massive stars that may also inhibit smoothing-out the valley. The valley would then be most pronounced for planets around massive stars.  To confirm the existence of the HMS period valley and test the scenario presented here, more detections of extrasolar planets around more massive stars (BA dwarfs) are needed. If our model is correct, the valley in the period/separation distribution of Jovian planets would become even more pronounced for even more massive stars. Because several planets around KG giants, which were BA dwarfs in their main sequence phase, have been discovered, plenty of planets should also exist around BA dwarfs. Spectroscopic observations are not easy for BA dwarfs because of a smaller number of absorption lines and rapid rotation.  However, in spite of the difficulty, relatively massive planets should be detectable. Transit surveys may not have serious problems to search for planets around BA dwarfs, however, detection is limited to short-period planets. Detailed observations of differences in disk sizes between T Tauri stars and Herbig Ae/Be stars are also needed to test our model.   {\\bf Acknowledgments:} A. Burkert acknowledges benefits from the activities of the RTN network \"PLANETS\" supported by the European Commission under the agreement No HPRN-CT-2002-0308. A. Burkert would like to thank the colleagues at the astronomy department (Lick Observatory) of the University of California, Santa Cruz for their hospitality."
    },
    "0608/astro-ph0608636_arXiv.txt": {
        "abstract": "We present a Fourier analysis of the clustering of galaxies in the combined Main galaxy and Luminous Red Galaxy (LRG) Sloan Digital Sky Survey (SDSS) Data Release 5 (DR5) sample. The aim of our analysis is to consider how well we can measure the cosmological matter density using the signature of the horizon at matter-radiation equality embedded in the large-scale power spectrum. The new data constrains the power spectrum on scales $100$--$600\\mpcoh$ with significantly higher precision than previous analyses of just the SDSS Main galaxies, due to our larger sample and the inclusion of the LRGs. This improvement means that we can now reveal a discrepancy between the shape of the measured power and linear CDM models on scales $0.01<k<0.15\\hompc$, with linear model fits favouring a lower matter density ($\\Omega_M=0.22\\pm0.04$) on scales $0.01<k<0.06\\hompc$ and a higher matter density ($\\Omega_M=0.32\\pm0.01$) when smaller scales are included, assuming a flat $\\Lambda$CDM model with $h=0.73$ and $n_s=0.96$. The lower matter density favoured by fitting our SDSS data for $0.01<k<0.06\\hompc$ is a better match to the best-fit WMAP 3-year cosmological model, and to results from the positions of the baryon oscillations observed in the SDSS DR5 power spectrum. This discrepancy could be explained by scale-dependent bias and, by analysing subsamples of galaxies, we find that the ratio of small-scale to large-scale power increases with galaxy luminosity, so all of the SDSS galaxies cannot trace the same power spectrum shape over $0.01<k<0.2\\hompc$. However, the data are insufficient to clearly show a luminosity-dependent change in the largest scale at which a significant increase in clustering is observed, although they do not rule out such an effect. Significant scale-dependent galaxy bias on large-scales, which changes with the $r$-band luminosity of the galaxies, could potentially explain differences in our $\\Omega_M$ estimates and differences previously observed between 2dFGRS and SDSS power spectra and the resulting parameter constraints. ",
        "introduction": "The evolution of perturbations in the early universe imprints characteristic scales that depend on the average matter density \\citep{silk68,peebles70,sunyaev70,bond84,bond87,holtzman89}. Fundamentally, the growth of fluctuations is intimately linked to the Jeans scale; perturbations smaller than the Jeans scale do not collapse due to pressure support, while larger perturbations are free to grow through gravity. In the radiation dominated era, the dark matter has negligible density compared to the photon-baryon fluid, and the perturbations in this fluid are stabilised by the high radiation pressure at a time when the sound speed was of order $c/\\sqrt{3}$. Consequently, the Jeans scale is of order the horizon scale until matter-radiation equality, after which it reduces to zero when the matter dominates. We therefore see that the horizon scale at matter-radiation equality will be imprinted in the distribution of fluctuations -- this scale marks a turn-over in the growth rate of fluctuations.  In a model with only collisionless dark matter, all lengths scale with the horizon scale at matter radiation equality, which is a multiple of $(\\Omega_Mh^2)^{-1}$, where $h=H_0/100\\,{\\rm km\\,s^{-1}Mpc^{-1}}$. Consequently, CDM transfer function fitting formulae (such as Eq. G3 of \\citealt{bbks}) were traditionally created as a function of $q\\equiv k / (\\Omega_Mh^2 {\\rm Mpc})$. When analysing galaxy redshift surveys, the comoving distance--redshift relation introduces another factor of $h$ into the measurements, so the data constrain the transfer function in $ k / \\hompc$. Consequently, fits between model and data constrain the degenerate parameter combination $\\Omega_Mh$.  Such cosmological constraints are important for breaking degeneracies between cosmological parameters that exist when fitting just CMB data \\citep{eisenstein99,efstathiou99,percival02}. With increasingly precise temperature and polarisation CMB constraints \\citep{hinshaw06,page06,spergel06}, the additional cosmological role of galaxy surveys to provide cross-checks is becoming increasingly important. In this paper, we carefully test the methodology behind using the shape of the galaxy power spectrum to provide cosmological constraints.  Extracting the cosmological information encoded in the galaxy power spectrum has motivated many previous surveys. Early studies, with of order $10^4$ galaxies, include the CfA \\citep{vogeley92,park94}, LCRS \\citep{shectman96} and PSCz \\citep{saunders00} surveys. These surveys were able to show that the shape of the power spectrum required a relatively small matter density, at odds with the simple Einstein-de Sitter model (e.g. \\citealt{efstathiou90,tadros99}). The same general shape of the power spectrum has been recovered by deprojecting the APM \\citep{maddox90,maddox96} survey, the parent catalogue of the 2dFGRS \\citep{efstathiou01,PB03}, and by deprojecting the photometric component of the SDSS catalogue \\citep{scranton02,connolly02,tegmark02,dodelson02,szalay03}. More recently, photometric redshifts have been exploited to improve the deprojection of galaxy distances and the accuracy with which the power spectrum can be measured from photometric data. Two studies have recently been completed considering Luminous Red Galaxies (LRGs) within the SDSS survey \\citep{padmanabhan06,blake06}, showing a marked improvement in accuracy on previous work.  However, the most precise measurement of the 3-dimensional galaxy power spectrum is still recovered from redshift surveys. Over the past 5 years there has been a revolution in the number of galaxy redshifts measured, and the size of the volume probed by these surveys. Leading the way have been two large galaxy redshift surveys, the 2dF Galaxy Redshift Survey (2dFGRS; \\citealt{colless01,colless03}) and Sloan Digital Sky Survey (SDSS; \\citealt{york00}). In this paper we analyse the relative clustering strengths of galaxy samples as a function of scale. The overall clustering amplitude is also potentially interesting in constraining cosmological models, but takes more work to decode. This normalisation is known to be a strong function of both galaxy colour and luminosity \\citep{park94,norberg01,norberg02,zehavi02}, so an understanding of galaxy bias is required before cosmological constraints can be derived from this statistic. However, the effect of galaxy bias on the shape of the power spectrum is less certain, and is the subject of this paper. On small scales, where non-linear corrections to the matter power spectrum are important, the shapes of galaxy power spectra are known to depend on galaxy colour and luminosity (e.g. \\citealt{cole05}). However, on large-scales, the effects on the shape are more uncertain.  On larger scales where the matter in the universe is still expected to be predominantly in the linear regime $k\\ls0.15\\hompc$ \\citep{smith03}, discrepancies currently exist between the shapes of the power spectra recovered from the 2dFGRS and SDSS. The problem is demonstrated by the measurements of $\\Omega_M$ obtained from such fits. Assuming a Hubble parameter $h=0.72$, the \\citet{tegmark04} SDSS main galaxy analysis favoured $\\Omega_M=0.296\\pm0.032$. Similar values of $\\Omega_M\\simeq0.3$ were found by alternative analyses of the red selected SDSS main galaxy power spectrum \\citep{pope04}, a set of SDSS spectroscopic Luminous Red Galaxies (LRGs) \\citep{eisenstein05}, and SDSS photometrically selected LRGs \\citep{padmanabhan06,blake06}. In combination with the WMAP 1-year data \\citep{spergel03}, the \\citet{eisenstein05} analysis of the LRGs provided the constraint $\\Omega_M=0.273\\pm0.025$ from a combination of the overall shape of the correlation function and the peak caused by baryon acoustic oscillations. In contrast, a lower matter density $\\Omega_M=0.231\\pm0.021$ is favoured by measurements of clustering of blue selected galaxies in the 2dFGRS \\citep{cole05}, who included a simple model for scale-dependent bias, although this had a relatively minor effect on the recovered matter density. Recent results from the 3-year WMAP data have provided an independent constraint on the matter density, finding $\\Omega_M=0.234\\pm0.035$ (mean constraint from Table 2 of \\citealt{spergel06}) from a better resolution of the third peak height. It is clear that these discrepancies between measurements are not linked to a single technique or particular analysis.  When the constraints from the galaxy power spectra are combined with the CMB data the discrepancy is still clear. From Table~5 of \\citet{spergel06}, we see that combining the 3-year WMAP data with additional cosmological constraints from the 2dFGRS power spectrum of \\citet{cole05} gives $\\Omega_M=0.236^{+0.016}_{-0.029}$. However, when the WMAP data is combined with constraints from the SDSS power spectrum of \\citet{tegmark04}, the higher $\\Omega_Mh$ constraint from the large-scale structure data increases the best fit to $\\Omega_M=0.266^{+0.025}_{-0.040}$ (see Table~6 of \\citealt{spergel06}).  The tension between measurements from different large-scale structure experiments and the CMB observations is at the level of approximately $2\\sigma$, and it is therefore possible that it could be explained by cosmic variance. Alternative explanations include a scale-dependent galaxy bias on scales $k\\ls0.15\\hompc$, or a systematic problem with one of the data sets. In this paper we test these possibilities by measuring the redshift-space power spectrum from the latest SDSS sample, Data Release 5 (DR5), which contains approximately twice as many galaxies as previously analysed, and 60\\% more LRGs than used in \\citet{eisenstein05}. By optimally analysing all of the galaxies to calculate the underlying power spectrum, we obtain the most accurate determination of the redshift-space power spectrum ever obtained for any sample of galaxies. First, we will test if the discrepancy between previous 2dFGRS and SDSS power spectra remains. Second, the size of the sample means that we can consider how well the galaxies trace the mass, by testing the simple hypothesis that the shape of the power spectrum matches linear CDM models over $k<0.15\\hompc$, and looking for discrepancies over these scales. Finally, the number of galaxies and volume covered are now sufficient to accurately measure power spectra for subsamples, splitting the catalogue to test for general changes in the shape of the power spectrum as a function of galaxy properties.  In a parallel paper, \\citet{tegmark06} present an analysis of a largely overlapping data set, drawn from SDSS DR4, with a focus on the implications for multi-parameter cosmological model fits. There are a number of differences in the analysis methods. \\citet{tegmark06} use a pseudo-Karhoenen-Loeve method \\citep{vogeley96,tegmark97} to estimate the real space galaxy power spectrum, using finger-of-god compression and linear theory to remove redshift-space distortion effects.  We use the Fourier method of \\citet{PVP}, which extends that of \\citet{FKP}, to estimate the angle-averaged (monopole) redshift-space galaxy power spectrum.  We combine the LRG and main galaxy samples, while \\citet{tegmark06} concentrate on the LRGs, after showing that they have a power spectrum shape consistent with that of the main sample. In addition, the many technical decisions that go into these analyses, regarding completeness corrections, angular masks, K-corrections, and so forth, were made independently for the two papers, and they present different tests for systematic uncertainties.  Despite these many differences of detail, our conclusions agree to the extent that they overlap (this will be discussed in \\citealt{tegmark06}), a reassuring indication of the robustness of the results. ",
        "conclusions": "\\label{sec:discussion}  We have analysed the clustering of galaxies in the largest sample (in terms of number and volume covered) to date. We have shown that the SDSS main galaxy and LRG samples can be naturally combined, because both sets of galaxies, after careful luminosity selection, can be described using the same average large-scale luminosity-dependent bias model. Analysing the combination of data allows us to include cross-correlation between the two data sets as well as internal auto-correlations within either the LRG or main galaxy samples.  Because of the speed and simplicity of the Fourier method used, we have been able to consider a number of tests of the recovered power spectrum, and have eliminated a number of possible systematics from our analysis. The only unusual result was a possible problem with the Galactic extinctions; this requires further detailed analysis and modelling that is outside the scope of our paper. In addition to considering a larger data set than previously analysed, our analysis extends previous work by considering the cosmological model to be tested from the start of the analysis -- using this model to convert from redshift to comoving distance, as well as comparing with the resulting power spectrum.  In a companion paper we have considered the cosmological constraints from the baryon acoustic oscillations observed in the power spectrum, where we found $\\Omega_M=0.250^{+0.042}_{-0.021}$ for flat $\\Lambda$ cosmological models \\citep{percival06b} when combined with the observed baryon acoustic oscillation scale in the CMB \\citep{spergel06}. This paper has instead focused on trying to model the overall shape of the power spectrum, paying particular attention to systematic effects on the resulting matter density constraints. Because of the accuracy with which we can now measure the power spectrum on scales $k<0.05\\hompc$ as a consequence of the large volume covered by the LRG sample, we can test the hypothesis that the shape of the redshift-space galaxy power spectrum matches the matter on scales where the matter clustering is well modelled by linear evolution from the seed perturbations ($0.01<k<0.15\\hompc$). We find a discrepancy between fits to $k<0.06\\hompc$, and $0.06<k<0.15\\hompc$, albeit at only moderate significance. The data on scales $0.01<k<0.06\\hompc$ favour a model with $\\Omega_M=0.22\\pm0.04$, while extending the $k$-range to $0.01<k<0.15\\hompc$ increases the estimated value to $\\Omega_M=0.32\\pm0.01$. Statistically, this discrepancy is only at the 2-3$\\sigma$ level.  It is worth noting that the excess power on scales $k < 0.06\\hompc$ relative to the $\\Omega_M=0.32$ model matches that recently found from analyses of the SDSS photometric-redshift sample \\citep{padmanabhan06,blake06}. This suggests that the discrepancy is therefore not limited to our analysis. The upper $k$ limit for the large-scale fit, $k_{\\rm lim}=0.06\\hompc$, was conservatively chosen based on the observed shape of the power spectrum (Fig.~\\ref{fig:pk_bigplot}). The noise in the data means that there is nothing special in this choice, and we could have split the scales fitted at $0.05<k_{\\rm lim}<0.1\\hompc$, and obtained similar results.  Interestingly, the change in the recovered matter density with the scales fitted matches the bimodality of previous matter density constraints, with the best-fit model on scales $0.01<k<0.06\\hompc$ matching that predicted by the WMAP 3-year data \\citep{spergel06}, the 2dFGRS \\citep{cole05}, and the results from the baryon acoustic oscillations observed in the SDSS DR5 sample. In contrast, the fits to $0.01<k<0.15\\hompc$ match the results from early studies of the SDSS data by \\citet{pope04} and \\citet{tegmark04}, suggesting that we are recovering approximately the same overall shape on these scales as these early analyses. The new data obviously constrain $\\Omega_Mh$ with greater accuracy so, on the scales previously considered, the significance of the offset between 2dFGRS and SDSS analyses has increased. Additionally, the higher value of $\\Omega_M=0.32\\pm0.01$ is now more discrepant with the (mean) WMAP constraint $\\Omega_M=0.234\\pm0.034$ (from Table~2 of \\citealt{spergel06}), and is discrepant with the lower matter density recovered from the positions of the baryon acoustic oscillations in the power spectrum \\citep{percival06b}.  The hypothesis being tested by these model fits is that the contributions to the power spectrum from non-linear structure growth and redshift-space distortions cancel, galaxy bias is scale-independent, and that the matter clusters as expected in a simple CDM model with an unbroken power spectrum. The results provide a contradiction to this hypothesis with a significance of 2-3$\\sigma$. We will now consider each of the elements in turn to try to understand which assumption is breaking down. First, it is worth noting that the effect of non-linear structure growth was discussed in Section~\\ref{sec:model} and shown to be too small to cause the observed power spectrum distortion.  Small-scale redshift-space distortions act in the wrong direction for the observed effect, and it would require a significant break-down of the scale-independent increase in power spectrum amplitude predicted by \\citet{kaiser87} for linear infall observed at large distances to give a scale-independent increase in the power spectrum amplitude and cause the observed change in shape. N-body simulations, and the halo model calculations performed for the 2dFGRS also found that the effect of redshift space distortions on the shape of the power spectrum is small on the scales of interest \\citep{percival01,cole05}. The consistency between the higher matter density favoured fitting to $0.01<k<0.15\\hompc$ in our study, previous studies of the real-space power spectrum \\citep{tegmark04}, and the power spectrum calculated from sets of photometric-redshift LRGs also provide evidence against redshift-space distortions producing the observed effects.  An unwelcome possibility is a systematic problem with one of the data sets, or the analysis method, although the only discrepancy revealed in the tests performed in Section~\\ref{sec:tests} was between the power spectra recovered using different Galactic extinction correction cuts. However, this potential systematic could also explain differences between 2dFGRS and SDSS power spectra if it predominantly affects the SDSS galaxies; analysing the galaxy clustering in the low extinction regions in the SDSS, for which \\citet{yahata06} show that the number density of galaxies does not behave as expected, produces a slight excess of power on scales $k\\simeq0.06\\hompc$ (see Fig.~\\ref{fig:pk_test2}). Additionally, it is perhaps also worth mentioning the heretical possibility that there is a problem with the assumption that the matter clusters as expected in a CDM model with an unbroken post-inflation power law spectrum. However, such a problem would affect the 2dFGRS and SDSS galaxies in the same way, so a further explanation would be required for this difference.  Perhaps the most simple explanation for this inconsistency is that the luminous red galaxies that dominate our combined sample do not trace the linear matter power spectrum as simply as other galaxies, and there is some theoretical work that supports this assertion.  In order to try to obtain evidence for scale-dependent galaxy bias, we have analysed the shape of the SDSS power spectrum, particularly looking for evidence of an increase in small-scale clustering power that depends on galaxy properties. The luminosity-bias correction applied as part of the Fourier method will only correct for the large-scale, scale-independent bias affecting different luminosity galaxies. It is not designed to correct for scale-dependent bias.  We have analysed a number of subsets of our final catalogue in order to look for scale-dependent bias, and find no significant change in shape if we split our sample using the bimodal galaxy colour distribution, or if we split in redshift; we recover similar power spectrum shapes, analysing the data in 5 redshift slices out to $z=0.5$. There is weak evidence for a change in the shape of the power spectrum for $0.01<k<0.15\\hompc$ when splitting the galaxies by luminosity, which varies with the average $r$-band luminosity of the galaxy sample analysed. However, the evidence for this is not conclusive. It is worth noting that we clearly see the effects of colour and luminosity on smaller scales $k>0.2\\hompc$, in line with the results of \\citet{cole05}. If scale-dependent bias also affects large scales then, in principle, this is testable by measurement of the bispectrum on the same data set \\citep{scoccimarro06}: the bispectrum shape and scale dependence respond to bias in a way that differs from, and is therefore not degenerate with, the power spectrum.  A change in power spectrum shape on scales $k<0.15\\hompc$, as a function of luminosity would provide a consistent picture when we compare our recovered power spectrum with previous work. On scales $k\\ls0.06\\hompc$, Fig.~\\ref{fig:pk_cmpr_previous} shows that the shape is consistent with the 2dFGRS power spectrum \\citep{cole05}. If the amplitudes of the power spectra are matched on these large scales, then on smaller scales there is a progression from the low clustering amplitude of the 2dFGRS galaxies, through the main galaxies of the DR3 SDSS sample analysed by \\citet{tegmark04} to the higher clustering strength of the combined main galaxy and LRG sample of this analysis. There is no reason to suggest that this is not a natural progression following the trend observed within the SDSS from galaxies with a low $r$-band luminosity to those with a high $r$-band luminosity. In this interpretation, the ``excess power'' on the largest scales probed by the \\citet{padmanabhan06} and \\citet{blake06} analyses of photometric LRGs is a consequence of comparing to an incorrect reference power spectrum with high $\\Omega_M$.  The true value of $\\Omega_M$ is lower, and the excess power is on small scales, as in Fig.~\\ref{fig:pk_bigplot} here.  The observed clustering was calculated using a weighted average over all galaxy pairs. Consequently, if we analyse the DR5 LRGs and main galaxies separately and match the large-scale clustering amplitudes, we would expect a progression from a lower small-scale amplitude of the main galaxies, through the combined sample, to a higher clustering amplitude of the LRGs. Fig.~\\ref{fig:pk_vs_lum_big} shows that this is indeed what we find.  The combination of scale-dependent bias, redshift-space distortions and non-linear structure growth for a set of galaxies has previously been matched by a simple fitting formulae (the $Q$-model) applied after calculation of the power spectrum \\citep{cole05}. However, it remains to be seen if this simple prescription can adequately describe galaxy bias for LRGs or for a mixed sample of galaxies with different clustering properties; these issues will be considered in \\citet{tegmark06}. The analysis presented in our paper suggests that this correction could become increasingly important for samples as the average $r$-band galaxy luminosity increases. It is clear that, for LRGs, the relation between galaxies and dark matter will need to be carefully modelled in order to extract the maximum possible information from the shape of the observed power spectrum."
    },
    "0608/astro-ph0608400_arXiv.txt": {
        "abstract": "Multicolor ($BVRI$) light curves of the eclipsing novalike variable TT~Tri are presented. The eclipse profiles are analyzed with a parameter-fitting model that assumes four sources of luminosity: a white dwarf primary star, a main-sequence secondary star, a flared accretion disk with a rim, and a bright spot at the intersection of the mass-transfer stream and the disk periphery. Model parameters include the temperatures of the white dwarf ($T_1$) and the secondary star ($T_2$), the radius ($R_\\mathrm{d}$) and temperature ($T_\\mathrm{d}$) of the disk periphery, the inner disk radius ($R_\\mathrm{in}$), the disk power-law radial brightness temperature exponent ($\\alpha$) and thickness parameter ($h_\\mathrm{r}$), and a bright spot temperature enhancement factor ($\\chi_\\mathrm{s}$).  A grid of model model light curves was computed, covering an extensive range of plausible parameter values. The models were then compared with the mean $BVRI$ light curves to determine the optimum values for the fitting parameters and their associated errors. The mass ratio of TT~Tri is poorly constrained in our models, but must lie roughly in the range $0.3\\lessim~q(=M_2/M_1)\\lessim0.9$. Models characterized by mass ratios of $q=0.3$ ($i=76.1^{\\circ}$), $q=0.6$ ($i=72.6^{\\circ}$), and $q=0.9$ ($i=70.4^{\\circ}$) were all capable of providing acceptable fits to the data, although the best fits were achieved for mass ratios near the upper end of the permitted range ($q=0.6$ and $q=0.9$). The values of the remaining fitting parameters were found to be insensitive to the adopted mass ratio and rim thickness. The accretion disk was found to extend to $\\sim50-60$\\% of the distance to the inner Lagrangian point in all models, but came closer to reaching the tidal limit (as expected for steady-state accretion) in the higher mass ratio models. The same behavior was found for the radial temperature profile of the disk, which increased with mass ratio, becoming more consistent with that expected for steady state accretion in the $q=0.6$ and $q=0.9$ models. Models with a truncated inner disk ($R_{\\mathrm{in}}>>R_1$) generally resulted in a higher white dwarf temperature, and a steepening of the disk temperature profile, but were not required to achieve a viable steady-state disk solution. No evidence was found for a luminous bright spot in the system, which is not surprising given the lack of a pre-eclipse ``hump\" in the light curve.  A total of 22 eclipse timings were measured for the system, which yielded an ephemeris for the times of mid-eclipse given by ${\\rm JD}_{\\odot} = 2,453,618.953(3) + 0.1396369(4)~E$. A comparison of the observed brightness and color at mid-eclipse with the photometric properties of the best-fitting model suggests that TT~Tri lies at a distance of $\\sim400-500$~pc. ",
        "introduction": "TT~Tri (GR 286, HS0129+2933) is a little-studied cataclysmic variable star, identified originally as an eclipsing binary by Romano (1978a,b) in a sample of eleven variable stars discovered on Schmidt plates of a field near M33. Except for the refined astrometric measurements of Sharov (1992), the star languished in obscurity until it surfaced in the Hamburg Quasar Survey (Hagen et al. 1995) as HS0129+2933. It has subsequently been suggested to be a member of the SW~Sex subclass of novalike variables by Rodr\\'igues-Gil (2005), who measured an orbital period of 3.35~hr for the system.  The study of eclipsing cataclysmic variables is valuable because it provides the opportunity to study the radiative properties of accretion disks through analysis of the eclipse profiles. In an attempt to further elucidate the properties of this long-neglected system, we have obtained extensive, time-resolved, multi-color CCD photometry of the eclipses of TT~Tri, and subsequently analyzed the resulting eclipse profiles with a parameter-fitting eclipse model. In this paper we present the results of that study. ",
        "conclusions": "We have performed the first multicolor eclipse study of the novalike variable TT~Tri. The best-fitting model parameters are typical of what is expected for a P=3.35~hr novalike variable, and are relatively insensitive to the adopted mass ratio and rim thickness. Our major conclusions are summarized below:  1) The mass ratio of TT~Tri is poorly constrained. As a result, we have considered a wide range of possible values in our analysis: $q=0.3$, $q=0.6$, and $q=0.9$. We found that not only do the higher mass ratio models provide a better fit to the data, they are characterized by disks that are more typical of a system with steady-state accretion (e.g. radii close to their tidal limit with steep radial temperature profiles). Nevertheless, acceptable model fits can be found for $q=0.3$, and thus lower mass ratios cannot be ruled out by the data. Fortunately, the values for the other fitting parameters are very insensitive to the adopted mass ratio.  2) The radius of the accretion disk in TT~Tri extends to $\\sim$50--60\\% of the distance to the inner Lagrangian point, which is somewhat smaller than the maximum size allowed by tidal effects. The disks in the higher mass ratio models fill a larger fraction of their tidal radius, as expected if the disk is in a state of steady accretion. In the $q=0.3$ models, the disk extends to only $\\sim$70\\% of the tidal radius, while for the $q=0.6$ and $q=0.9$ models the disk extends to $\\sim$80\\% and $\\sim$90\\% of the tidal radius, respectively. This result provides indirect evidence that the mass ratio in TT~Tri may be closer to the maximum value allowed for stable accretion, and that the white dwarf may have a relatively low mass ($M_1\\simeq0.3-0.5$M$_{\\odot}$).  3) As found in many other novalike systems, the best fitting models of TT~Tri are not necessarily characterized by disks with steep radial temperature profiles ($\\alpha\\simeq0.75$). For example, SD models with values of $\\alpha$ as low as 0.4 were found to provide acceptable fits to the data, particularly in the cases where $q=0.3$. It is worth recalling that steady-state accretion does not necessarily require $\\alpha=0.75$. To begin with our models fit brightness temperature, not the the effective temperature used in eqn (2). Furthermore, the effects of irradiation of the outer disk from hotter regions near the white dwarf, and the potential cooling effects of winds from the inner accretion disk, both of which would have the effect of reducing the radial disk temperature gradient, have not been included in our models. Despite these caveats, it remains true that the smaller the value of $\\alpha$, the less confident we can be that the disk is in a state of steady accretion.  4) Unlike the case of V~Per, where a thick disk rim was required to achieve temperature profiles approaching $\\alpha=0.75$, the thickness of the rim had little effect in TT~Tri. However, significantly steeper values of $\\alpha$ were found when the inner disk was truncated, as one would expect if the white dwarf is magnetic. A likely explanation for the different behavior between V~Per and TT~Tri is that the orbital inclination is significantly lower in TT~Tri. A lower inclination will make the disk rim less important (e.g. Smak 1994), while making the contribution of the inner disk relatively more important. Despite the fact that the TD models are characterized by significantly larger values of $\\alpha$, as just discussed, steady-state solutions were found for the SD models as well. Thus, in our modeling, we find no compelling evidence for a truncated inner disk in TT~Tri.  5) We estimate that TT~Tri is characterized by a mass accretion rate, $\\dot M\\sim 2\\times10^{17}$~g~s$^{-1}$, which is just above the critical rate for stable accretion. Given the proximity of the accretion rate to the critical value, it is possible that future observations of TT~Tri may reveal dwarf nova eruptions similar to those seen in the Z~Cam stars, where the system fluctuates between state of stable and unstable accretion (e.g. Shafter et al. 2005). The TD models are characterized by somewhat higher accretion rates, making them more typical of novalike systems with stable accretion.  6) Finally, by comparing the observed brightness of TT~Tri, both in and out of eclipse, with properties of the secondary star and accretion disk, we find that the system likely lies at a distance of $\\sim400-500$~pc from Earth. The principle source of error in this estimate comes from uncertainty in the assumed properties of the secondary star."
    },
    "0608/astro-ph0608293_arXiv.txt": {
        "abstract": "On October 28, 2003 an earthward-directed coronal mass ejection (CME) was observed from SOHO/LASCO imagery in conjunction with an X17 solar flare.  The CME, traveling at nearly 2000~km/s, impacted the Earth on October 29, 2003 causing ground-based particle detectors to register a counting rate drop known as a Forbush decrease.  In addition to affecting the rate of cosmic rays, the CME was also responsible for causing anisotropies in the direction of incidence.  Data from Project GRAND, an array of proportional wire chambers, are presented during the time of this Forbush decrease. A simple model for CME propagation is proposed and we present an argument based on gyroradius that shows that a magnetic field of the radius calculated for the ejecta is sufficient to deflect energetic charged particles of an energy detectable by GRAND. ",
        "introduction": "In late October of 2003 the Earth experienced an extraordinary amount of solar and geomagnetic activity originating from solar region NOAA 10486, including an X17 flare which was one of the largest solar flares since 1976 \\cite{IPS03}.  This flare was detected beginning at 9:51~UT on October 28, 2003 (day 301) and peaked at 11:10~UT the same day.  This flare had an associated CME with a transit time from the sun to the Earth of only 19 hours, making it one of the the fastest on record.  The shock from the CME impacted the Earth's magnetic field as a strong sudden impulse (an abrupt increase in the horizontal component of the geomagnetic field) at 6:13 UT on October 29, 2003 (day 301) \\cite{NOAA}.  This CME caused a drop in the counting rate of ground-based cosmic ray detectors.  The counting rate for neutron monitor stations remained suppressed for roughly twelve days after the impact of the shock.  Project GRAND \\cite{Poirieret03}, an extensive air shower array of proportional wire chambers, is a useful tool for studying the ground-level effects of cosmic rays.  GRAND's median energy for vertically incident cosmic rays is 56 GeV, higher than that for neutron monitors. This allows GRAND to complement neutron monitor data by studying higher energy effects.  GRAND has an angular resolution of $0.26^{\\circ}$ on a projected plane for incoming muon tracks. The angular resolution, primary energy sensitivity, and large detector area of this experiment (82~$m^2$) make it an excellent instrument to to study Forbush decreases. Secondary muon data obtained by GRAND during the time of the October 29 Forbush decrease are discussed.  In addition, a simple model for a CME ejecta is proposed and some of the consequences of this model are discussed. ",
        "conclusions": "Project GRAND sees an 8\\% drop in its secondary muon counting rate during the Forbush decrease of October 29, 2003.  A shift is also observed in the mean angle of incident muons in the north-south plane.  A simple model for the propagation of a CME through the solar system is presented.  Within a range of initial conditions for the CME, the model is  consistent with satellite observations near Earth during the October 29, 2003 storm."
    },
    "0608/astro-ph0608546_arXiv.txt": {
        "abstract": "\\baselineskip = 11pt \\leftskip = 0.65in \\rightskip = 0.65in \\parindent=1pc {\\small We review the current state of observational work on the formation of brown dwarfs, focusing on their initial mass function, velocity and spatial distributions at birth, multiplicity, accretion, and circumstellar disks. The available measurements of these various properties are consistent with a common formation mechanism for brown dwarfs and stars. In particular, the existence of widely separated binary brown dwarfs and a probable isolated proto-brown dwarf indicate that some substellar objects are able to form in the same manner as stars through unperturbed cloud fragmentation. Additional mechanisms such as ejection and photoevaporation may play a role in the birth of some brown dwarfs, but there is no observational evidence to date to suggest that they are the key elements that make it possible for substellar bodies to form. \\\\~\\\\~\\\\~}% ",
        "introduction": "}     Although many of the details are not perfectly understood, stars and giant planets are generally believed to form through the collapse of molecular cloud cores and the accretion of gas by rocky cores in circumstellar disks, respectively. In comparison, the formation of objects intermediate between stars and planets -- free-floating and companion brown dwarfs -- has no widely accepted explanation. A priori, one might expect that brown dwarfs form in the same manner as stars, just on a much smaller scale. However, although self-gravitating objects can form with initial masses of only $\\sim1$~$M_{\\rm Jup}$ in simulations of the fragmentation of molecular cloud cores, these fragments continue to accrete matter from their surrounding cores, usually to the point of eventually reaching stellar masses ({\\em Boss}, 2001; {\\em Bate et al.}, 2003). Thus, standard cloud fragmentation in these models seems to have difficulty in making brown dwarfs. One possible explanation is that the simulations lack an important piece of physics (e.g., turbulence), and brown dwarfs are able to form through cloud fragmentation despite their predictions (e.g., {\\em Padoan and Nordlund}, 2004). Another possibility is that a brown dwarf is born when cloud fragmentation is modified by an additional process that prematurely halts accretion during the protostellar stage, such as dynamical ejection ({\\em Reipurth and Clarke}, 2001; {\\it Boss}, 2001; {\\em Bate et al.}, 2002) or photoevaporation by ionizing radiation from massive stars ({\\em Kroupa and Bouvier}, 2003; {\\em Whitworth and Zinnecker}, 2004). This uncertainty surrounding the formation of brown dwarfs has motivated a great deal of theoretical and observational work over the last decade.    \\begin{figure*} \\epsscale{2.2} \\plottwo{luhman_fig1a.ps}{luhman_fig1b.ps} \\caption{\\small IMFs for Taurus ({\\em Luhman}, 2004c), IC~348 ({\\em Luhman et al.}, 2003b), Chamaeleon I ({\\em Luhman}, in preparation), and the Trapezium Cluster ({\\em Muench et al.}, 2002). The completeness limits for these measurements are near 0.02~$M_{\\odot}$ ({\\it dashed lines}). In the units of this diagram, the Salpeter slope is 1.35. \\label{fig:imf}} \\end{figure*}  In this paper, we review the current observational constraints on the formation process of brown dwarfs (BDs), which complements the theoretical review of this topic provided in the chapter by {\\em Whitworth et al.} By the nature of the topic of this review, we focus on observations of BDs at young ages ($\\tau<10$~Myr), although we also consider properties of evolved BDs that provide insight into BD formation (e.g., multiplicity). A convenient characteristic of young BDs is their relatively bright luminosities and warm temperatures compared to their older counterparts in the solar neighborhood, making them easier to observe. However, because the luminosities and temperatures of young BDs are continuous extensions of those of stars, positively identifying a young object as either a low-mass star or a BD is often not possible. The mass estimates for a given object vary greatly with the adopted evolutionary models and the manner in which observations are compared to the model predictions. Using the models of {\\em Baraffe et al.} (1998) and {\\em Chabrier et al.} (2000) and the temperature scale of {\\em Luhman et al.} (2003b), the hydrogen burning mass limit at ages of 0.5-3~Myr corresponds to a spectral type of $\\sim$M6.25, which is consistent with the dynamical mass and spectral type of the first known eclipsing binary BD ({\\em Stassun et al.}, 2006) and other observational tests ({\\em Luhman and Potter}, 2006). Therefore, we will treat young objects later than M6 as BDs for the purposes of this review. ",
        "conclusions": ""
    },
    "0608/hep-ph0608186_arXiv.txt": {
        "abstract": "Scalar self interactions are known to weaken considerably the current constraints on scalar-mediated fifth forces.  We consider a scalar field with a quartic self interaction and gravitation-strength Yukawa couplings to matter particles.  After discussing the phenomenology of this scalar field, we assess the ability of ongoing and planned experiments to detect the fifth force mediated by such a field.  Assuming that the quartic and matter couplings are of order unity, the current-generation E\\\"ot-Wash experiment at the University of Washington will be able to explore an interesting subset of parameter space. The next-generation E\\\"ot-Wash experiment is expected to be able to detect, or to rule out, the fifth force due to such a scalar with unit quartic and matter couplings at the 3$\\sigma$ confidence level. ",
        "introduction": "Years of effort have been devoted to searching for new macroscopic forces from sub-millimeter to solar system scales~\\cite{fischbach}. From a theoretical standpoint, modern theories of particle physics introduce new scalar fields which can mediate long range forces. This is certainly true of string theory, whose plethora of moduli generically couple to matter with a strength comparable to that of gravity. Provided they remain light, these scalars should cause observable deviations from the gravitational inverse square law and violations of the equivalence principle.  The experimental state of affairs is shown in Fig.~\\ref{expcons}.  Evidently, a fifth force of gravitational strength, $\\alpha\\sim {\\cal O}(1)$, is excluded on all scales ranging from 0.1~mm to $10^2$~AU. A crucial underlying assumption, however, is that the mediating scalar field has negligible self-interactions. As argued in~\\cite{gubserkhoury}, the addition of a quartic term drastically changes the picture. Gravitational-strength coupling to matter is allowed again, even with a quartic coupling as small as $\\sim 10^{-53}$~\\cite{nelson}. The ability for self-interacting scalars to hide from experiments relies on the chameleon mechanism~\\cite{original,gubserkhoury}, which suppresses fifth-force signals in two ways.  \\begin{figure} \\includegraphics[width=5in]{figures/longlamnew.eps} \\caption{Current experimental constraints on the strength $\\alpha$ and range $L$ (in meters) of a Yukawa fifth force, ignoring self-interactions. Reprinted from~\\cite{adellong}.} \\label{expcons} \\end{figure}  Firstly, the presence of ambient matter density generates a tadpole term in the Klein-Gordon equation, which shifts the minimum of the potential. Because of the quartic coupling, the mass of small fluctuations around this effective minimum can be much larger than the mass in vacuum. This is most emphatically illustrated with a massless field with quartic coupling of order unity: in a medium of density 1~g/${\\rm cm}^3$, the effective mass is $0.1$~mm.  Secondly, the fifth force is further suppressed by the thin shell effect, another startling consequence of the nonlinearity of the field equations. Within a dense macroscopic body, the effective mass of the scalar is large. As a result, the contribution of the core to the external field is exponentially small.  Only a thin shell near the surface exerts a significant pull on an exterior test particle. In fact, for an infinite plate, the force at the surface eventually saturates as the thickness is increased, causing the strength relative to gravity to fall off rapidly.  It follows that an ideal experiment to detect a chameleon-like scalar must: 1.~use sufficiently small test masses to minimize the thin shell suppression; and 2.~probe the force at distances $\\lsim 1$~mm to avoid the exponential damping from the effective mass. These considerations point towards the E\\\"ot-Wash experiment at the University of Washington~\\cite{hoyle}.  In this paper we carefully assess the ability of the E\\\"ot-Wash experiment to detect or exclude chameleon scalar fields with quartic self-interactions. We find that the current apparatus with two disks of 42 holes each is not sensitive enough to detect a chameleon force with dimensionless quartic coupling of order unity, $\\lambda\\sim {\\cal O}(1)$, and gravitational strength coupling to matter, $\\beta =1$. It should nevertheless place significant constraints on models with larger values of $\\beta$ or smaller $\\lambda$. The next generation E\\\"ot-Wash experiment, whose design involves two disks with 120 wedges removed, is much more promising for detecting self-interacting scalars. Our calculations show it should detect or exclude chameleon fields with $\\lambda=\\beta=1$ at the 3$\\sigma$ level. Unfortunately it looks unlikely that the next generation E\\\"ot-Wash will be able to distinguish between a chameleon-mediated force and a Yukawa force with suitable strength and range.  We begin in Sec.~\\ref{phenom} with a summary of the phenomenology of self-interacting scalar fields. While it is well-known that free scalars mediate attractive forces, in Sec.~\\ref{attract} we argue that this is also the case for chameleon scalars. The chameleon mechanism and thin-shell effect are reviewed in Secs.~\\ref{subsec:chameleonmech} and~\\ref{subsec:thinshell}, respectively. We describe in Sec.~\\ref{predicts} the expected fifth force signals for self-interacting scalars for the current and next generation E\\\"ot-Wash experiments. Section~\\ref{conclu} summarizes our results and discusses prospects for distinguishing chameleon-mediated forces from Yukawa forces.    \\newcommand\\TL{\\hfil$\\displaystyle{##}$} \\newcommand\\TR{$\\displaystyle{{}##}$\\hfil} \\newcommand\\TC{\\hfil$\\displaystyle{##}$\\hfil} \\newcommand\\TT{\\hbox{##}} \\newcommand\\JOT{\\noalign{\\vskip1\\jot}} \\def\\seqalign#1#2{\\vcenter{\\openup1\\jot \\halign{\\strut #1\\cr #2 \\cr}}} \\def\\lbldef#1#2{\\expandafter\\gdef\\csname #1\\endcsname {#2}} \\newcommand{\\eqn}[3][]{\\lbldef{#2}{(\\ref{#2})}% \\def\\@eqnstyle{#1}% \\ifx\\@eqnstyle\\@empty% \\begin{equation} \\eqalign{#3} \\label{#2} \\end{equation}% \\else% \\begin{equation} \\seqalign{\\span\\TC}{#3} \\label{#2} \\end{equation}% \\fi} \\def\\eqalign#1{\\vcenter{\\openup1\\jot \\halign{\\strut\\span\\TL & \\span\\TR\\cr #1 \\cr }}} \\def\\eno#1{(\\ref{#1})} ",
        "conclusions": "\\label{conclu}  We have shown that a chameleon field with a quartic self interaction mediates an attractive force, which falls off rapidly with separation between two massive objects, and is sensitive only to the outer shell of matter in a large object.  Given the size of its test masses and the length scales that it probes, the E\\\"ot-Wash experiment is a promising instrument for searching for this fifth force. For unit values of the matter coupling constant $\\beta$ and the quartic coupling constant $\\lambda$, the current E\\\"ot-Wash apparatus is not expected to find evidence for the chameleon fifth force.  However, the experiment is capable of constraining a very interesting region of parameter space; the expected chameleon signal $N_{21} = 0.11 \\beta^{0.91} \\lambda^{-0.55}$ will be detectable for certain values of $\\beta$ and $\\lambda$ not much different from unity.  The next-generation E\\\"ot-Wash experiment will be several times more sensitive to chameleon-mediated fifth forces.  We expect E\\\"ot-Wash to detect, or to rule out, a chameleon force with unit couplings at the $3\\sigma$ level.  If a small-scale deviation from Newtonian gravity is observed, this deviation may be compared to our predictions for $\\beta=\\lambda=1$ in Fig.~\\ref{fig:fourier}.  The torque Fourier coefficients at other values of the coupling constants, and at a pendulum-attractor separation of $\\Delta z = 0.05$mm, are $N_{120} = 0.38 \\beta^{1.34}\\lambda^{-0.33}$ fNm and $N_{360} = 0.031 \\beta^{1.68} \\lambda^{-0.16}$ fNm.  Thus, if $\\beta$ is large enough for the chameleon force to be detected by the current E\\\"ot-Wash experiment, then the next-generation experiment should detect a peak flattening $N_{360} > 0$.  If a short-range fifth force is observed by the E\\\"ot-Wash experiment, the next goal will be to distinguish between chameleon and Yukawa fifth forces.  Although unit-strength Yukawa fifth forces have already been ruled out by E\\\"ot-Wash, a Yukawa force tuned to be a few orders of magnitude smaller than gravity could conceivably fit the observed fifth force.  As argued in Sec. \\ref{subsec:chameleonmech}, chameleon and Yukawa forces that are identical at small scales begin to differ substantially at length scales of the order of $m_{\\text{eff,}\\rho}^{-1}$.  Unfortunately, a chameleon force with $\\beta=\\lambda=1$ is only visible over a range of pendulum-attractor separations $0.05\\textrm{mm} \\leq \\Delta z \\leq 0.1\\textrm{mm}$, a range whose width is approximately equal to the length scale $m_{\\text{eff,}\\rho}^{-1} = 0.047\\textrm{mm}$ inside the pendulum and attractor wedges.  Distinguishing between a chameleon force and a Yukawa force will be challenging, and unless $\\beta$ is sufficiently greater than unity, a more sensitive experiment will be necessary."
    },
    "0608/nucl-th0608008_arXiv.txt": {
        "abstract": "Spin polarized states in nuclear matter with Skyrme effective forces are studied on the base of  a Fermi liquid theory  for a wide range of isospin asymmetries and densities. There are a few possible scenarios of spin ordered phase transitions: (a) nuclear matter with  SLy4 interaction undergoes at some critical density a phase transition to a spin polarized state with the oppositely directed spins of neutrons and protons; (b) for  SkI5 interaction, a spin polarized state with the like-directed neutron and proton spins is formed; (c) nuclear matter with SkI3 interaction under increasing density, at first, undergoes a phase transition to the state with the opposite directions of neutron and proton spins, which goes over at larger density to the state with the same direction of nucleon spins. Spin polarized states at strong isospin asymmetry are accompanied by the long tails in the density profiles of the neutron spin polarization parameter near the critical density, if the energy gain of the transition from the nonpolarized state to a polarized one is the decreasing function of isospin asymmetry (SLy4 force). If the energy gain is increased with isospin asymmetry, there are no long tails in the density distribution of the neutron spin polarization parameter (SkI3, SkI5 forces). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608220_arXiv.txt": {
        "abstract": "% During a program of spectroscopic monitoring of V838~Mon, we serendipitously discovered that a neighboring 16th-mag star is of type~B\\null. We then carried out a spectroscopic survey of other stars in the vicinity, revealing two more B-type stars, all within $45''$ of V838~Mon. We have determined the distance to this sparse, young cluster, based on spectral classification and photometric main-sequence fitting of the three B~stars. The distance is found to be $6.2\\pm1.2$~kpc, in excellent agreement with the geometric distance to V838~Mon of 5.9~kpc obtained from {\\it Hubble Space Telescope\\/} polarimetry of the light echoes. The cluster's age is less than 25~Myr.  The absolute luminosity of V838~Mon during its outburst, based on our distance measurement, was very similar to that of M31~RV, an object in the bulge of M31 that was also a cool supergiant throughout its eruption in 1988. However, there is no young population at the site of M31~RV.  It does not appear possible to form a nova-like cataclysmic binary system within the young age of the V838~Mon cluster, and the lack of a young population surrounding M31~RV suggests that the outburst mechanism does not require a massive progenitor. These considerations appear to leave stellar-collision or \\hbox{-merger} scenarios as one of the remaining viable explanations for the outbursts of V838~Mon and M31~RV. ",
        "introduction": "A variety of explanations for the outburst of V838~Mon  have been proposed, many of them mutually exclusive (see recent summary in Tylenda \\&  Soker 2006, and the  proceedings of this conference). These explanations involve either thermonuclear processes (an unusual nova-like outburst on a white dwarf, a thermonuclear event in a massive star, or a helium shell flash in a post-AGB star), or the release of gravitational energy (through stellar or planetary mergers or collisions).  A possible new constraint on the nature of V838~Mon came from the discovery of a B3~V companion to the star (Munari \\& Desidera 2002; Wagner \\& Starrfield 2002). The companion is unresolved even at {\\it Hubble Space Telescope\\/} ({\\it HST\\/}) resolution.  Another eruptive object has attracted attention as a possible analog of V838~Mon. This is the ``M31 red variable,\" or ``M31 RV,'' which underwent an outburst in mid-1988 that was remarkably similar to that of V838~Mon (Bond \\& Siegel 2006 and references therein), although not as well observed.  M31~RV occurred in the nuclear bulge of the Andromeda Galaxy, and is thus at a known distance.  In this paper, we present our serendipitous discovery that V838~Mon is a member of a small, young open cluster. We will use spectral classification and photometry of the cluster members to determine a distance. We will also derive a limit to the cluster's age, and will compare the stellar populations surrounding V838~Mon and M31~RV\\null. We will close with a brief discussion of the V838~Mon progenitor object and some new constraints on the outburst mechanism that result from our observations. ",
        "conclusions": ""
    },
    "0608/astro-ph0608016_arXiv.txt": {
        "abstract": "The study of detached eclipsing binaries in open clusters can provide stringent tests of theoretical stellar evolutionary models, which must simultaneously fit the masses, radii, and luminosities of the eclipsing stars and the radiative properties of every other star in the cluster. We review recent progress in such studies and discuss two unusually interesting objects currently under analysis. GV\\,Carinae is an A0\\,m + A8\\,m binary in the Southern open cluster NGC\\,3532; its eclipse depths have changed by 0.1\\,mag between 1990 and 2001, suggesting that its orbit is being perturbed by a relatively close third body. DW\\,Carinae is a high-mass unevolved B1\\,V + B1\\,V binary in the very young open cluster Collinder\\,228, and displays double-peaked emission in the centre of the H$\\alpha$ line which is characteristic of Be stars. We conclude by pointing out that the great promise of eclipsing binaries in open clusters can only be satisfied when both the binaries and their parent clusters are well-observed, a situation which is less common than we would like. ",
        "introduction": "Detached eclipsing binary stars (dEBs) are of fundamental importance to stellar physics because they are, apart from the few closest objects to the Earth, the only stars for which we can accurately measure basic quantities such as mass, radius and surface gravity \\cite{and91}. The realism and reliability of the current generation of theoretical stellar evolutionary models ranks as one of the great achievements of modern astrophysics, but this success would have been much more difficult without the ability to check the effects of particular physics against the accurate physical properties of stars in dEBs.  Given good photometric and spectroscopic data, it is possible to derive masses and radii of stars in a dEB to accuracies better than 1\\% and surface gravities to within 0.01\\,dex \\cite{me5}. However, the predictions of theoretical models can usually match even properties as accurate as this, both because of their sophistication and because there are several important unconstrained parameters, e.g., metal and helium abundance and age. More constraints are needed to investigate the success or otherwise of a number of physical parameters which are only very simplistically treated in theoretical models, e.g., convective core overshooting, mass loss, mixing length and rotational effects. For example, small changes in the mixing length can change the derived ages of the oldest globular clusters, which constrain the age of the Universe, by 10\\% \\cite{cha95}.  An answer to this problem is to study dEBs which are members of stellar clusters \\cite{me1,tho01}. Because the dEB and the other cluster members have the same age and chemical composition, theoretical models must be able to simultaneously match the masses, radii and luminosities of the two stars in the dEB and the radiative parameters of every other cluster member, for one age and chemical composition. This allows much more detailed tests to be made of the success or otherwise of different physical ingredients in models. Alternatively, if the cluster is poorly studied, a comparison of the properties of the dEB with model predictions allows the cluster metal abundance and age to be derived.  Detached eclipsing binaries are also excellent distance indicators \\cite{cla04}, through a variety of methods such as using bolometric corrections or surface brightness calibrations (see \\opencite{me4} for a detailed analysis). dEBs in clusters therefore give an accurate distance to the cluster without the problems and theoretical dependence which affect the main sequence fitting method.  An unique advantage of studying dEBs in clusters is that it can be possible to place four or more stars with the same age and chemical composition onto one mass--radius or \\Teff--\\logg\\ plot, if several dEBs are members of one cluster (\\opencite{me1}; \\citeyear{me3}). Some Galactic open clusters are known to contain four or more dEBs, e.g., NGC\\,7086 (Robb et al., these proceedings). ",
        "conclusions": ""
    },
    "0608/astro-ph0608366_arXiv.txt": {
        "abstract": "We investigate the suggestion that there are stellar populations in some globular clusters with enhanced helium ($Y \\sim$ 0.28 to 0.40) compared to the primordial value.  We assume that a previous generation of massive Asymptotic Giant Branch (AGB) stars have polluted the cluster. Two independent sets of AGB yields are used to follow the evolution of helium and CNO using a Salpeter initial mass function (IMF) and two top-heavy IMFs.  In no case are we able to produce the postulated large $Y \\sim 0.35$ without violating the observational constraint that the CNO content is nearly constant. ",
        "introduction": "\\par Star-to-star abundance variations of the light elements C, N, O, Na, Mg and Al have been observed in every well studied globular cluster (GC) to date \\citep[][and references therein]{kraft94,gratton04} but are not found in field stars of the same metallicity \\citep{gratton00}. Hence these abundance anomalies are somehow the result of the cluster environment. The variations of the elements follow a common pattern: C-N,  O-Na and Mg-Al are all anti-correlated \\citep{shetrone96a,kraft97,cannon98,gratton01,cohen05a,cohen05b}. The abundances of iron-peak, s and r-process elements do not show the same star-to-star scatter nor do they vary {\\em with} the light elements \\citep{gratton04,james04,yong06b}, although new observations by \\citet{wylie06} suggest that there is real star-to-star scatter amongst heavy element abundances in stars in the metal-rich cluster 47 Tucanae. The other important exception is the massive cluster $\\omega$ Centauri, whose age and metallicity spread, along with a rise in s-element abundance with increasing [Fe/H], suggests that it evolved very differently from other GCs and may even have an extragalactic origin \\citep{smith00}. The key points are that O has been destroyed in some stars by up to one dex, the C$+$N$+$O abundances remain almost constant and there is no evidence for large-scale variation of the neutron-capture elements. \\par Two hypotheses had been proposed to explain these observed abundances. The first was deep mixing, where the abundance anomalies are produced by internal mixing during the ascent of the giant branch, after the first dredge-up \\citep{sweigart79,pin97,charbonnel94}. However, star-to-star abundance variations in C, N, O and Na were subsequently observed in stars at or near the main-sequence turn-off \\citep{gratton01,ramirez03,james04,cohen05a,cohen05b}. These observations support the self-pollution scenario, first proposed by \\citet{cottrell81}. Here a previous generation of stars polluted the atmospheres of stars we observe today or provided part of the material from which those stars formed. Because [Fe/H] is roughly  constant in stars in a given GC it has been assumed that polluters were intermediate-mass AGB stars with initial masses between $\\sim$3 to 8$\\Msun$ rather than supernovae, which produce Fe. The hot bottom burning experienced by these stars provides the hydrogen burning environment (at least qualitatively) that can alter the abundances of the light elements. One consequence is the production of a significant quantity of helium \\citep[$Y$, \\iso{4}He;][]{lattanzio04}. The mass lost via the slow winds of AGB stars could, in principle, have been retained by the cluster from which new stars may have been born \\citep{thoul02}. Detailed AGB models have so far mostly failed to match the observed abundance trends \\citep{herwig03b,fenner04,campbell04,cohen05b}, but major uncertainties that affect the models undermine the reliability of the predictions and leave room for an AGB solution \\citep{ventura05a,ventura05b}. \\par Horizontal branch (HB) and main-sequence color-magnitude diagrams (CMD) \\citep{norris04,dantona04,lee05,piotto05} provide increasing evidence for helium enrichment in some GC stars. The unusual HB morphology of NGC 2808, which exhibits an extended blue tail and a gap separating the red and blue clumps \\citep{bedin00}, can be most easily explained if the blue stars have a higher helium content ($Y \\sim 0.32$) compared to those in the red clump which presumably have primordial $Y \\approx 0.24$ \\citep{dantona04}. Furthermore, NGC 2808 also has a peculiar main sequence \\citep{dantona05} where the bluer stars are inferred to have $Y \\sim 0.4$ from fitting theoretical isochrones to the observed data.  To reproduce the CMD, \\citet{dantona05} note the absolute necessity of including a small population ($\\sim 20$\\%) of stars with $Y = 0.40$ and assume a spread of $Y$ between 0.24 -- 0.29 to fit the main fraction ($\\sim 80$\\%).  While the exact maximum value of $Y$ required to reproduce the CMD is dependent on the color-$T_{\\rm eff}$ transformations used in the analysis and therefore rather uncertain, it seems that the most plausible explanation is that the bluest stars have enhanced amounts of helium with $Y \\gtrsim 0.35$.  Observations by \\citet{piotto05} show that the blue main-sequence of $\\omega$ Centauri is more metal-rich than the red sequence, contrary to what is expected from stellar models and \\citet{norris04} showed that isochrones with $Y=0.40$ best fit the bluest stars.  Until a better explanation for these intriguing observations is found, we take them as motivation to study the AGB self-pollution scenario from a global perspective. \\par In this paper, we use the \\citet{fenner04} GC chemical evolution model to follow the evolution of helium in the intracluster gas. We probe AGB model uncertainties by using two independent sets of yields, including those used in the previous study which were tailor made for NGC 6752 (with a metallicity [Fe/H] $\\approx -1.4$\\footnote{Slightly more metal-rich than the average of [Fe/H] $\\sim -1.6$ for NGC 6752, M13, M3 and $\\omega$ Cen with metallicities from \\citet{harris96}; NGC 2808 has [Fe/H] $\\sim -1.15$.}). The evolution of C, N and O is also followed since they impose important empirical constraints, i.e. C$+$N$+$O $\\approx$ constant, that must be met by the model. ",
        "conclusions": "Our investigation into the chemical evolution of helium in GCs highlights the difficulty the AGB self-pollution scenario suffers in trying to explain the large postulated helium enrichment required to fit the horizontal branch of clusters like NGC 2808. With a standard Salpeter IMF, the largest predicted helium abundance in the cluster gas is $Y \\approx 0.29$ but this is accompanied by a large increase in the C$+$N$+$O abundance. Using an independent set of AGB yields from \\citet{ventura02} we find a maximum $Y \\approx 0.26$ and only a modest increase of C$+$N$+$O $\\approx 0.4$~dex,  probably within the observational errors. We conclude that with a standard IMF it does not seem likely that the AGB self-pollution mechanism alone produced the enormous amounts of helium inferred from observations of the bluest HB and main-sequence stars of clusters like NGC 2808 and $\\omega$ Centauri.  Simulations that employ the IMS-enhanced IMF show larger helium enhancements  of $Y \\approx 0.35$ but only  when accompanied by enormous increases in the total C$+$N$+$O content of the cluster gas, in violation of observations. The Ventura et al. yields predict a maximum $Y \\approx 0.28$ and the total CNO abundance stays constant to within $\\sim 0.3$\\,dex, although we again point out that this maximum $Y$ is made uncertain by extrapolating the yields to higher masses. Even with such an extreme IMF we have a problem fitting the observational constraints.  Indeed, the use of such an IMF does not help the difficulties faced by the self-pollution scenario in matching the constraints that we have considered  in this study i.e. $Y \\gtrsim 0.30$ and C$+$N$+$O $\\approx$ constant. Both sets of yields considered here also fail to match the observed spread of abundances between O, Na, Mg and Al, even though the quantitative predictions of the models are different. For example, the Fenner et al. models produce too much Na whereas the Ventura et al. yields destroy sodium, resulting in a O-Na correlation.  \\citet{bekki06} discussed the consequences of a number of top-heavy IMFs on the evolution of $\\omega$ Centauri and concluded that the ones most suitable for producing the large helium enrichment  would also tear the cluster apart.  While $\\omega$ Cen is an unusual cluster indicated by the spread in Fe, and clearly had a very different chemical enrichment history to the other lower-mass GCs, the conclusions reached about the top-heavy IMF on the cluster evolution are applicable to smaller mass, less tightly bound systems. Given the difficulties associated with the self-pollution scenario, we may need to look to other solutions such as pollution from outside the cluster.  Indeed, perhaps the most unusual of all clusters, $\\omega$ Cen, is actually just an extreme member and is telling us something useful about all GCs."
    },
    "0608/astro-ph0608685_arXiv.txt": {
        "abstract": "We study correlations in spatial orientation between galaxy clusters and their host superclusters using a Hubble Volume N-body realization of a concordance cosmology and an analytic model for tidally-induced alignments. We derive an analytic form for distributions of the alignment angle as functions of halo mass ($M$), ellipticity ($\\epsilon$), distance ($r$) and velocity ($v$) and show that the model, after tuning of three parameters, provides a good fit to the numerical results.  The parameters indicate a high degree of alignment along anisotropic, collapsed filaments. The degree of alignment increases with $M$ and $\\epsilon$ while it decreases with $r$ and is independent of $v$.  We note the possibility of using the cluster-supercluster alignment effect as a cosmological probe to constrain the slope of the initial power spectrum. ",
        "introduction": "Superclusters are collections of galaxy groups and clusters that represent the largest, gravitationally bound structures in the universe \\citep{sha30,kal-etal98}.  If the dark energy is a cosmological constant, then the collapse of these systems over the next few billion years of the cosmic future will mark the end of hierarchical structure formation in our universe \\citep{nag-loeb03,busha05}. A conspicuous feature of locally observed superclusters is the strong tendency of member clusters to be elongated in their major axis orientations \\citep{pli02,pli04}, which is in turn closely related to their filamentary shapes \\citep[e.g.,][]{bas03}. To describe the structure distribution on the largest scale in the universe,  it will be quite essential to understand this effect of cluster-supercluster alignments from first principles.  The effect of structure-substructure alignment is in fact observed on all different scales in the universe. On the subgalactic scale the galaxy satellites are observed to be located preferentially near the major axes of their host galaxies \\citep{val-etal78,kne-etal04,bra05,agu-bra06}. On the galactic scale the major axes of cluster galaxies are observed to be aligned with that of their host clusters \\citep{pli-bas02,pli-etal03}. The cluster galaxies are also observed to have a strong tendency toward radial alignment \\citep{per-kuh05}.  Although this alignment effect has been shown to be a natural outcome in the currently favored concordance $\\Lambda$CDM cosmology \\citep{onu-tho00,lib-etal05,kan-etal05,lee-etal05,zen-etal05,kas-evr05, bas-etal06}, its detailed origin remains a subject of debate between those who emphasize the importance of anisotropic merging and those who stress tidal interaction.  The anisotropic merging scenario explains that the effect of substructure-structure alignment is induced by the anisotropic merging and infall of matter along filaments \\citep{wes89}. It was indeed shown by N-body simulations that the merging and infall of matter to form bound halos indeed occur preferentially along filaments, which provided supporting evidences for this scenario \\citep[e.g.,][]{wes-etal91,van-van93,dub98,fal-etal02,kne-etal04,zen-etal05}.  The tidal interaction theory explains that the correlations between the substructure angular momentum vectors and the principal axes of the host tidal fields induce the alignment effect. \\citet{lee-etal05} constructed an analytic model for the effect of substructure alignment in the frame of the tidal interaction theory, and showed that their analytic predictions are in good agreement with the numerical results from N-body simulations.  In fact, the above two theories are not mutually exclusive since the anisotropic merging and infall itself is a manifestation of the primordial tidal field \\citep{bon-etal96}. The difference between the two scenarios, however, lies in the question of whether the connection to filaments is a major contribution or not.  Very recently, \\citet{atl-etal06} have quantified the influences of both the tidal interaction and the anisotropic infall through the analysis of data from recent high-resolution N-body simulations. What they confirmed is the following: (i) For the majority of halos the alignment effect is caused by the tidal field but not by the anisotropic infall; (ii) Only for the cluster-size halos the alignment effect is dominantly due to the anisotropic merging and infall of matter along filaments. In other words, they made it clear that the filaments are important marker of local orientations on the cluster halo scale.  Now that the cluster-supercluster alignment turns out to be due to the anisotropic merging along filaments, it is desirable to have a theoretical frame work within which one can provide physical answers to the remaining questions such as how the alignment effect depends on the cluster properties such as mass, shape, and etc.  Our goal here is to construct such a theoretical framework by using both analytical and numerical methods. Analytically we adopt the standard cosmic web theory, and numerically we use the large Hubble volume simulation data.  The organization of this paper is as follows. In $\\S 2$ we provide a brief description of the Hubble volume simulation and summarize the numerical results. In $\\S 3$ we present an analytic model and compare its predictions with the numerical results. In $\\S 4$ we discuss our results and draw final conclusions. ",
        "conclusions": "In the context of the standard cosmic web picture of large-scale structure, we have constructed a parametric model for the alignment of cluster-sized halos with their host superclusters. The underlying assumption is that cluster-supercluster alignment reflects the spatial distribution of matter as it is organized along filamentary structures by the primordial tidal field. The parameters of the analytic model represent the dominance of filaments and the spatial coherence of the initial tidal field.  We show that the analytic model provides a good fit to orientation data derived from mass-limited halo samples of a $\\Lambda$CDM Hubble volume simulation.  After fitting the three free parameters using the overall distribution of cluster-supercluster alignment angles, the model then simultaneously matches the behavior of the mean alignments as a function of relative mass, cluster position within the supercluster, and cluster ellipticity.  No trend with cluster velocity is predicted or measured in the simulation.  It is worth discussing a couple of simplified assumptions on which our theoretical model is based. First,we have used the FOF algorithm to identify superclusters in the N-body simulation data. Unlike the case of virialized clusters, however, there is no established consensus on how to define superclusters. Different supercluster-identification algorithms could result in different multiplicity, mass, and shape of superclusters which would in turn affect our results.  Second, we have assumed that the filaments correspond to the Lagrangian regions where only the largest and the second largest eigenvalues of the local tidal tensor are positive. Although this definition of a filament is consistent with the picture of the Zel'dovich approximation, it is obviously an oversimplification of the reality. A more realistic definition and treatment of cosmic filaments will be necessary to refine the model.  Another issue that we would like to discuss here is the possibility of using the cluster-supercluster alignment effect as a cosmological probe. We have shown that the phenomena of cluster-supercluster alignments are closely related to the dominance of filaments, the web-like distribution of galaxies on very large scales. The dominance of filaments is in turn related to the spatial correlations of the primordial tidal field, which depends sensitively on the slope of the initial power spectrum on the supercluster scale. Thus, by measuring the degree of cluster-supercluster alignment, it might be possible to constrain the slope of the initial power spectrum in a complementary way.  Finally, we conclude that our model for the cluster-supercluster alignments will provide a theoretical framework within which the distribution of cosmic structures on the largest scales can be physically understood and quantitatively described."
    },
    "0608/gr-qc0608117_arXiv.txt": {
        "abstract": "\\noindent Loop quantum cosmology homogeneous models with a massless scalar field show that the big-bang singularity can be replaced by a big quantum bounce. To gain further insight on the nature of this bounce, we study the semi-discrete loop quantum gravity Hamiltonian constraint equation from the point of view of numerical analysis. For illustration purposes, we establish a numerical analogy between the quantum bounces and reflections in finite difference discretizations of wave equations triggered by the use of nonuniform grids or, equivalently, reflections found when solving numerically wave equations with varying coefficients. We show that the bounce is closely related to the method for the temporal update of the system and demonstrate that explicit time-updates in general yield bounces. Finally, we present an example of an implicit time-update devoid of bounces and show back-in-time, deterministic evolutions that reach and partially jump over the big-bang singularity. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608150_arXiv.txt": {
        "abstract": "We present the Fundamental Plane (FP) in the $z = 0.28$ cluster of galaxies \\rxj. There is no evidence for a difference in the slope of the FP when compared with the Coma cluster, although the internal scatter is larger. On average, stellar populations in \\rxj \\ have rest-frame $V$-band mass-to-light ratios ($M/L_V$) $0.29 \\pm 0.03$ dex lower than in Coma. This is significantly lower than expected for a passively-evolving cluster formed at $z_{f} = 2$. Lenticular galaxies have lower average $M/L_V$ and a distribution of $M/L_V$ with larger scatter than ellipticals. Lower mass-to-light ratios are not due to recent star formation: our previous spectroscopic observations of \\rxj \\ E/S0 galaxies showed no evidence for significant star-formation within the past $\\sim 4$ Gyr. However, cluster members have enhanced $\\alpha$-element abundance ratios, which may act to decrease $M/L_V$. The increased scatter in the \\rxj \\ FP reflects a large scatter in $M/L_V$ implying that galaxies have undergone bursts of star formation over a range of epochs. The seven easternmost cluster galaxies, including the second brightest member, have $M/L_V$ consistent with passive evolution and $z_{f} = 2$. We speculate that \\rxj \\ is a cluster--cluster merger where the galaxies to the east are yet to fall into the main cluster body or have not experienced star formation as a result of the merger. ",
        "introduction": "The Fundamental Plane (FP) of elliptical (E) and lenticular (S0) galaxies is a log-linear relation between effective radius, surface brightness, and velocity dispersion (\\eg Djorgovski \\& Davis 1987; Dressler \\etal 1987; J{\\o}rgensen \\etal 1996, hereafter JFK1996)\\nocite{djorgovski87,dressler87,jorgensen96}. Its ubiquity in nearby clusters of galaxies has been noted by many investigators (\\eg Bender \\etal 1992; JFK1996; Colless \\etal 2001)\\nocite{bender92,jorgensen96,colless01}.  Coefficients in the observed FP differ from those predicted by the virial theorem, and this `tilt' implies that mass-to-light ratio ($M/L$) is not a constant for E/S0s but depends to some extent on galaxy properties (see Cappellari \\etal 2006\\nocite{cappellari06}).  The FP must confront studies showing that the growth of galaxies in clusters involves complex processes like galaxy mergers, cluster mergers, and AGN activity (\\eg Fabian \\etal 2000; Tran \\etal 2005)\\nocite{fabian00,tran05}.  If E/S0 galaxies form in a short burst and their stellar populations subsequently evolve passively, one would expect a gradual decrease in $M/L$ with cosmic epoch.  This induces an offset in the FP as a function of redshift, with no change in slope. However, the tendency of star formation to occur over a longer period in lower-mass galaxies (\\eg Juneau \\etal 2005; Treu \\etal 2005)\\nocite{juneau05,treu05}, would be manifest as a steepening of the empirical FP with redshift. Evidence for changes in the slope of the FP are seen in the latest studies of the FP at $z>0.8$ (\\eg Holden \\etal 2005; J{\\o}rgensen \\etal 2006)\\nocite{holden05,jorgensen06}, but this has not been observed at $0.2 < z < 0.8$ (\\eg van Dokkum \\& Franx 1996; Kelson \\etal 2000; Wuyts \\etal 2004)\\nocite{vandokkum96,kelson00b,wuyts04}, most likely because these studies do not probe far enough down the mass function. The internal scatter of the FP reported at $0.2 < z < 0.8$ is consistent with that of the Coma cluster (JFK1996)\\nocite{jorgensen96}. Since scatter in the FP is a reflection of the variation in $M/L$, this is taken as evidence that cluster E/S0s at a given mass are coeval.  In this letter we examine the FP in \\rxj, a cluster of galaxies at $z=0.28$ for differences in slope and internal scatter with respect to that of Coma. We use \\cosmology, and convert previous work to these values where necessary. The look-back time to \\rxj \\ in this cosmology is 3.2 Gyr. ",
        "conclusions": "\\label{sec:dis}  We have established the FP for E/S0 galaxies in \\rxj. There is no evidence that the slope of this FP is different from that of the Coma cluster. On average, $M/L_V$ for galaxies in \\rxj \\ is $0.28 \\pm 0.03$ dex lower than in the Coma cluster.  The internal scatter of the FP in \\rxj \\ is larger than for Coma. This can also be visualised as a larger scatter in mass-to-light ratio for galaxies in \\rxj. Lenticular galaxies have lower average $M/L_V$ with a larger scatter than Es. Spectroscopically-confirmed members east of 01h42m06s, including the second brightest cluster galaxy, have $M/L_V$ ratios consistent with passively-evolving galaxies formed at $z_f \\sim 2$ and an internal scatter consistent with the Coma FP.  The large scatter in the FP can be explained by spatially-segregated galaxies with different average $M/L_V$, consistent with galaxies in \\rxj \\ having undergone bursts of star formation over a range in epochs. This effect is not evident when looking at age-sensitive absorption line indices. However, galaxies in \\rxj \\ have enhanced \\afe \\ which might serve to push any putative star formation episode to higher redshift. If this were the case our data suggest that increased \\afe \\ acts to reduce $M/L_V$. Improved modeling of the effect of age and \\afe \\ on $M/L_V$ is required before more accurate statements about star-formation epochs can be made.  We speculate that the large scatter in galaxies' $M/L_V$ is caused by a cluster--cluster merger. This scenario is consistent with the large cluster velocity dispersion and implies that gaseous or galactic substructure should be detectable with X-ray imaging or spectroscopy of hundreds of cluster galaxies."
    },
    "0608/astro-ph0608199_arXiv.txt": {
        "abstract": "The remarkable improvement in the estimates of different cosmological parameters in recent years has been largely spearheaded by accurate measurements of the angular power spectrum of Cosmic Microwave Background (CMB) radiation.  This has required removal of foreground contamination as well as detector noise bias with reliability and precision. Recently, a novel {\\em model-independent} method for the estimation of CMB angular power spectrum from multi-frequency observations has been proposed and implemented on the first year WMAP (WMAP-1) data by Saha et al.~2006. We review the results from WMAP-1 and also present the new angular power spectrum based on three years of the WMAP data (WMAP-3). Previous estimates have depended on foreground templates built using extraneous observational input to remove foreground contamination.  {\\em This is the first demonstration that the CMB angular spectrum can be reliably estimated with precision from a self contained analysis of the WMAP data}.  The primary product of WMAP are the observations of CMB in $10$ independent difference assemblies (DA) distributed over $5$ frequency bands that have uncorrelated noise.  Our method utilizes maximum information available within WMAP data by linearly combining DA maps from different frequencies to remove foregrounds and estimating the power spectrum from the $24$ cross power spectra of clean maps that have independent noise.  An important merit of the method is that the expected residual power from unresolved point sources is significantly tempered to a constant offset at large multipoles (in contrast to the $\\sim l^2$ contribution expected from a Poisson distribution) leading to a small correction at large multipoles. Hence, the power spectrum estimates are less susceptible to uncertainties in the model of point sources. ",
        "introduction": "Remarkable progress in cosmology has been made due to the measurements of the anisotropy in the cosmic microwave background (CMB) over the past decade.  The extraction of the angular power spectrum of the CMB anisotropy is complicated by foreground emission within our galaxy and extragalactic radio sources, as well, as the detector noise~\\cite{bouc_gisp99,Tegmark96}. It is established that the CMB follows a blackbody distribution to high accuracy~\\cite{FIRAS}. Hence, foreground emissions may be removed by exploiting the fact that their contributions in different spectral bands are considerably different while the CMB power spectrum is same in all the bands~\\cite{Dodelson, Tegmark2000a, Bennett1,Tegmark98}.  Different approaches to foreground removal have been proposed in the literature~\\cite{bouc_gisp99,Tegmark96,hob98,main0203,erik06}.  The Wilkinson Microwave Anisotropy Probe (WMAP) observes in $5$ frequency bands at $23~\\ghz$~(K), $33~\\ghz$~(Ka), $41~\\ghz$~(Q), $61~\\ghz$~(V) and $94~\\ghz$~(W).  In the first data release, the WMAP team removed the galactic foreground signal using a template fitting method based on a model of synchrotron, free free and dust emission in our galaxy \\cite{Bennett1}. The sky map around the galactic plane and around known extragalactic point sources were masked out and the CMB power spectrum was then obtained from cross power spectra of independent difference assemblies in the $41~\\ghz$, $61~\\ghz$ and $94~\\ghz$ foreground cleaned maps~\\cite{hin_wmap03}.  A model independent removal of foregrounds has been proposed in the literature~\\cite{Tegmark96}.  The method has also been implemented on the WMAP data in order to create a foreground cleaned map~\\cite{Tegmark}. The main advantage of this method is that it does not make any additional assumptions regarding the nature of the foregrounds.  Furthermore, the procedure is computationally fast.  The foreground emissions are removed by combining the five different WMAP bands by weights which depend both on the angular scale and on the location in the sky (divided into regions based on `cleanliness'). However, the power spectrum recovered from the single foreground cleaned map has severe excess power at large multipole moments due to amplification of detector noise bias in the autocorrelation power spectrum beyond the beam resolution.  The prime objective of our method~\\cite{sah06} is to adapt and extend the foreground cleaning approach to power spectrum estimation by retaining the ability to remove detector noise bias exploiting the fact that it is uncorrelated among the different Difference Assemblies (DA) \\cite{hin_wmap03,Jarosik}. The WMAP data uses 10 DA's \\cite{Bennett, Bennett2, WMAP, Hinshaw1}, one each for K and Ka bands, two for Q band, two for V band and four for W band.  We label these as K, Ka, Q1, Q2,V1, V2, W1, W2, W3, and W4 respectively. We eliminate the detector noise bias using cross power spectra and provide a model independent extraction of CMB power spectrum from WMAP first year data.  So far, only the three highest frequency channels observed by WMAP have been used to extract CMB power spectrum and the foreground removal has used foreground templates based on extrapolated flux from measurements at frequencies far removed from observational frequencies of WMAP~\\cite{hin_wmap03,Pablo,Patanchon}.  We present a more general procedure where we use observations from all the five frequency channels of WMAP and do not use any observational input extraneous to the WMAP data set. ",
        "conclusions": "The rapid improvement in the sensitivity and resolution of the CMB experiments has posed increasingly stringent requirements on the level of separation and removal of the foreground contaminants.  Standard approaches to foreground removal, usually attempt to incorporate the extra information about the amplitude, spatial structure and distribution of foregrounds available at other frequencies, in constructing a foreground template at the frequencies of the CMB measurements. These approaches could be susceptible to uncertainties and inadequacies of modeling involved in extrapolating from the frequency of observation to CMB observations.  We carry out an estimation of the CMB power spectrum from the WMAP data that is independent of foreground model and evades these uncertainties. The novelty is to make clean maps from the difference assemblies and exploit the lack of noise correlation between the independent channels to eliminate noise bias.  {\\em This is the first demonstration that the angular power spectrum of CMB anisotropy can be reliably estimated with precision solely from the WMAP data (difference assembly maps) without recourse to any external data.} Our work is a clear demonstration that the blind approach to foreground cleaning is comparable in efficiency to that from template fitting methods and certainly adequate for a reliable estimation of the angular power spectrum. Moreover, this method enjoys the advantage of eliminating the uncertainties of modeling required to create the foreground templates that estimate the contamination at the CMB dominated frequencies by extrapolating from the observed emission flux at very different frequencies where the foregrounds dominate.  Further, the tempered, flat, large $l$ behavior of residual from unresolved point sources in this method holds promise for experiment with finer angular resolution.  The understanding of polarized foreground contamination in CMB polarization maps is rather scarce. Hence modeling uncertainties could dominate the systematics error budget of conventional foreground cleaning.  The blind approach extended to estimating polarization spectra after cleaning CMB polarization maps could prove to be particularly advantageous."
    },
    "0608/astro-ph0608144_arXiv.txt": {
        "abstract": "\\baselineskip=18pt An analysis of $n=3 \\to 2$ transition lines of carbon-like silicon reveals that some ratios of line intensities are sensitive to the electron density. The ratio between two group of $3d\\to2p$ transition lines at 55.246~\\AA\\, and 55.346~\\AA\\, is a good $n_{\\rm e}$-diagnostic technique, due to its insensitivity to the electron temperature. Using this property, a lower limit of the density of 0.6$\\times10^8$cm$^{-3}$ is derived for Procyon, which is consistent with that constrained by C V and Si X emissions. Significant discrepancies in ratios of $3s\\to2p$ lines to $3d\\to2p$ lines between theoretical predictions and observed values, are found, by the spectral analysis of Procyon observed with the {\\it Chandra} High Resolution Transmission Grating spectra. The difference exceeding a factor of 3, cannot be explained by the uncertainty of atomic data. The opacity effect is also not a choice as reported by Ness and co-workers. For the $3s\\to2p$ line at 61.611~\\AA\\,, present work indicates that the large discrepancy may be due to the contamination from a S VIII line at 61.645~\\AA\\,. For the lines at 61.702 and 61.846~\\AA\\,, we suggest that the discrepancies may be attributed to contaminations of unknown lines. \\newline \\vspace{2cm} {\\it PACS:} 95.20.Jg; 97.10.Ex; 95.85.Nv ",
        "introduction": "Since the launch of the new generation X-ray observatories, {\\it Chandra} and {\\it XMM-Newton}, X-ray spectra with unprecedented spectral resolution and high statistical quality can be obtained for astrophysical objects. The High Energy Transmission Grating Spectrometer (HETGS) and Low Energy Transmission Grating Spectrometer (LETGS) on {\\it Chandra} can provide resolving power of $\\lambda/\\Delta\\lambda~\\sim~100-1000$ and $\\ge1000$ in wavelength ranges of 1.5-30\\AA\\, and 50-160\\AA, respectively. The effective areas can be as large as 200~cm$^2$ in some energy regions. Using these high-quality spectra, physical conditions of stellar coronae such as electron density, thermal structure, and element abundances, can be derived. Such information further helps to constrain models of the coronal heating and structuring. In the past, the information about the spatial distribution of coronal plasma was inferred by indirect means such as modelling of eclipses and rotational modulation (G$\\ddot{\\rm u}$del et al., 1995, 2003; Siarkowski et al., 1996), and such analyses can only be carried out for very special systems with advantageous geometries. Presently, the electron density ($n_{\\rm e}$) of stellar coronae can be disentangled from the emission measure ($EM=\\int n_{\\rm e}^2dV$), and determined separately using the X-ray spectra with high-resolution for various late-type stars.  The complexes of He-like ``triplets\" of C, N, O, Ne, Me, and Si, which include the {\\it resonance} ($r$: $1s^2~^1S_0 - 1s2p~^1P_1$), {\\it inter-combination} ($i$: $1s^2~^1S_0 - 1s2p~^3P_{2,1}$), and {\\it forbidden} ($f$: $1s^2~^1S_0 - 1s2s~^3S_1$) lines, are covered by the HETGS and/or LETGS on {\\it Chandra} and Reflection Grating Spectrometer (RGS) on {\\it XMM-Newton}, and are resolvable. Gabriel \\& Jordan (1969) pointed out the potential diagnostic application of these features for hot plasmas. Recently, Porquet \\& Dubau (2000) recalculated line intensities by accounting for other processes such as dielectronic recombination and radiative recombination. In the HETGS observation of Capella, Canizares et al. (2000) estimated a density range of 0.8-2$\\times10^{10}$~cm$^{-3}$ from O~VII, and upper limits near 7$\\times10^{11}$ and 1$\\times10^{12}$~cm$^{-3}$ by Mg~XI and Si~XIII. Using the LETGS observation of Capella, Brinkman et al. (2000) derived a density of $\\sim2.6\\times10^9$~cm$^{-3}$ by C V for the lower temperature component of a multi-temperature structure. Later, Ness et al. (2002) and Testa et al. (2004) performed coronal density diagnostics with He-like triplets for other late-type stars with various activities such as Algol, Procyon, $\\alpha$~Cen~A\\&B, $\\epsilon$~Eri, and HR~1099 etc. Using the derived densities, Ness et al. (2004) further estimated the sizes of stellar X-ray coronae, and concluded that the cooler plasma component cannot cover a large fraction of the stellar surface, and there must be spatially separate plasma components at different temperatures. They also argued that the hotter plasma loops fill the space between cooler loops until much of the corona is dominated by the hotter plasma.  The parameter of the electron density is also used to judge the mechanism of the X-ray production. For T~Tauri stars, a density of $>1\\times10^{12}$~cm$^{-3}$ has been obtained from O~VII for TW~Hya (Stelzer \\& Schmitt, 2004) and BP~Tau (Schmitt et al., 2005), which is higher than typical coronal densities by at least two orders of magnitude. It was assumed that the X-ray emission is produced in the accretion shocks of T~Tauri stars, instead of stellar coronae. Additionally, the UV radiative field from the inner layer also plays an important role on the population of levels $^3P_{2,1}$, which directly influences the estimation of $n_{\\rm e}$. For lighter elements, this effect is more sever (Ness et al., 2001). In the coronal X-ray spectra, in addition of the emissions of H-, He-like ions and Fe L-shell ions, many lines from highly charged silicon have been observed in the LETGS observations, such as that of Procyon. These emissions may give further insight into the physical conditions of stellar coronae. Recently, we estimated the electron densities and/or temperatures for the cooler components using Si~X and Si~XI features (Liang et al. 2006a). The emissions from $3d\\to2p$ and $3s\\to2p$ transitions of Si~IX have also been identified by Raassen et al. (2002) in the LETGS observation of Procyon. In this paper, we investigate the soft X-ray emissions of Si~IX in detail.  This paper is organized as followings, observations and measurements of line fluxes are described in Sect.2. Theory model of C-like Si~IX is presented in Sect.3. Sect.4 outline the results and discussions, and the conclusion is given in the last section. ",
        "conclusions": "A detailed analysis of soft X-ray emissions of carbon-like silicon reveals that the ratio of Si~IX lines around 55.246 and 55.346~\\AA\\, is a good $n_{\\rm e}$-diagnostic method. By comparing the observed ratio with theoretical prediction, a lower limit of the density (0.6$\\times10^8$~cm$^{-3}$) was obtained for Procyon, which is the typical value of the solar quiescent corona. The constructed spectrum of Si~IX suggests that the major contribution at 55.083\\AA\\,, and 61.611\\AA\\, should be assigned to Si~IX.  In intensity ratios of $3s\\to2p$ transition lines of Si~IX relative to that of $3d\\to2p$ line at 55.346~\\AA\\,, significant discrepancies between observations and theoretical predictions are found in Procyon. For the $3s\\to2p$ line at 61.611~\\AA\\,, the observation and theoretical prediction agree within a factor of 2, when the contamination arising from S~VIII line (61.645~\\AA\\,) has been considered. Here, an experimental ratio of 0.75 (Lepson et al. 2005) between S~VIII lines at 61.645 and 52.781~\\AA\\, was adopted to disentangle the blending. However, for other two possible $3s\\to2p$ lines of 61.792 and 61.848~\\AA\\,, large differences (nearly an order of magnitude) cannot be explained by the uncertainty of atomic data and the resonant scatter effect. So we suggest that the most possible reason is contributions of unknown emissions, which may be addressed by further improvements in the soft X-ray spectral analysis."
    },
    "0608/astro-ph0608691_arXiv.txt": {
        "abstract": "We explore the possibility of putting constraints on dark energy models with statistical property of large scale structure in the non-linear region. In particular, we investigate the $w$ dependence of non-Gaussianity of the smoothed density distribution generated by the nonlinear dynamics. In order to follow the non-linear evolution of the density fluctuations, we apply N-body simulations based on $P^3 M$ scheme. We show that the relative difference between non-Gaussianity of $w=-0.8$ model and that of $w=-1.0$ model  is $0.67 \\%$ (skewness) and $1.2 \\%$ (kurtosis) for $R=8h^{-1}$ Mpc. We also calculate the correspondent quantities for $R=2h^{-1}$ Mpc, $3.0 \\%$ (skewness) and $4.5 \\%$ (kurtosis), and the difference turn out to be  greater, even though non-linearity in this scale is so strong that the complex physical processes about galaxy formation affect the galaxy distribution. From this, we can expect that the difference can be tested by all sky galaxy surveys with the help of mock catalogs created by selection functions, which suggests that non-Gaussianity of the density distribution potentially plays an important role for extracting information on dark energy. ",
        "introduction": "\\label{sec:intro}     The formation of large scale structure in the Universe is a central element of astrophysics. The key idea in understanding the formation of the structure in the Universe is gravitational instability \\cite{Peebles80,Peacock,Liddle,Coles}. Roughly stated, if the material in the Universe distributed irregularly, then the overdense regions provide extra gravitational attraction, draw material toward them and become more overdense. At present, this simple picture dramatically succeeds in explaining what is observed and current attention is focused entirely on the details of the gravitational instability process which strongly depends on the cosmological models.  One of the useful means to extract the detail of structure formation is the statistical property of observable quantities. Among various statistical quantities, the density probability distribution function (PDF) of the density contrast plays a very important role because it includes information of non-Gaussianity. In the standard picture, the primordial density fluctuations originated from quantum fluctuations are stretched to large comoving scales during the inflation phase and  assumed to be random Gaussian \\cite{Starobinsky:1982ee}. It is, however, well known that even though the density fluctuations remain Gaussian in the linear regime, it significantly deviates from Gaussian once the non-linear stage is reached \\cite{Kofman94,Yokoyama96}. Since structure formation is competition between the cosmic expansion and gravity, the growth of the density contrast as well as the evolution of its statistical quantities are intimately tied to the cosmic expansion law of the Universe.  On the other hand, according to recent observations for type Ia Supernovae~\\cite{Perlmutter99}, the expansion of the Universe is accelerating. Combining measurements of Cosmic Microwave Background Radiation~\\cite{WMAP03} and recent galaxy redshift survey ~\\cite{Tegmark04}, we are forced to recognize the existence of a cosmological constant, or kind of dark energy whose value is almost the same order of magnitude as the present density of the Universe~\\cite{WMAP}. (For reviews about the cosmological constant problem, see \\cite{Weinberg89}.) From the phenomenological viewpoint, in order to clarify the property of dark energy, it is important to constrain the effective equation of state of dark energy $w$ by observations. Even though we have obtained the constraint as $w < -0.90$ (95 \\% confidence limit assuming $w \\ge -1$)~\\cite{WMAP} by combining WMAP data with other astronomical data, in order to pin down the value of $w$, it is necessary to propose other independent and complementary methods.  For this purpose, as stated at the beginning, we suppose the statistical property of large scale structure of the Universe is helpful. Actually, we were the first to analyze the $w-$dependence of non-Gaussianity of the PDF of the smoothed density contrast in \\cite{Tatekawa06} where we considered simple constant-$w$ dark energy models with $w=-0.5, -0.8, -1.0, -1.2$. (The case for the cosmological constant ($w=-1$) had already been considered in such as \\cite{Plionis95,Borgani95,Kayo01}.) There, in order to see the effects of the non-linear dynamics analytically, as a first step, we adopted Lagrangian linear perturbation \\cite{Zeldovich70,Arnold82,Tatekawa05}  which describes the evolution of the density fluctuation in the quasi-nonlinear region better than second order Eulerian perturbation \\cite{Munshi:1994zb,Yoshisato:1997eb}.  As a natural extension to the previous one, in order to go into the strongly nonlinear region where the difference of non-Gaussianity depending on $w$ is expected to be large, numerical simulations seem to be necessary. Even though the simulations including dark energy are important, because of computational expense of the simulations, the first simulations including dark energy other than the cosmological constant appeared only in 2003. By now the aim of N-body simulations in the presence of dark energy are limited to the mass function of dark matter halo \\cite{Linder:2003dr,Klypin:2003ug,Kuhlen:2004rw} as well as the matter power spectrum \\cite{McDonald:2005gz}. Therefore, in this paper, we analyze the $w$-dependence of non-Gaussianity of the PDF of the smoothed density contrast with N-body simulations to cover the strongly nonlinear region where the results of \\cite{Tatekawa06} remain invalid.  This paper is organized as follows. In section \\ref{sec:dark_energy}, we briefly summarize the effects of dark energy on the structure formation scenario and define the quantities which are necessary to estimate non-Gaussianity of the density distribution in section \\ref{sec:non-Gaussianity}. The set up and results of N-body simulations are described in Section \\ref{sec:N-body} and Section \\ref{sec:results}, respectively. Section \\ref{sec:summary} is devoted to conclusions. ",
        "conclusions": "\\label{sec:summary}  To clarify the nature of dark energy is, without doubt, one of the most important tasks in modern cosmology. From the phenomenological viewpoint, it is important to constrain the effective equation of state of dark energy $w$ by observations. Even though we have obtained the constraint as $w < -0.90$ (95 \\% confidence limit assuming $w \\ge -1$)~\\cite{WMAP} by combining WMAP data with other astronomical data, to pin down the value of $w$, it is necessary to propose other independent observational tools.  For this purpose, we consider non-Gaussianity of the density distribution which is generated by the non-linear dynamics. In this paper, as a natural extension to our previous work \\cite{Tatekawa06} where we follow a semi-analytic approach based on Lagrangian linear perturbation theory, we apply N-body simulations based on $P^3 M$ codes. In terms of the density fluctuation obtained by N-body simulations, we compute the skewness and kurtosis of the smoothed density distribution with the appropriate definition in that they become constant in extremely weakly nonlinear stage. They are important statistical quantities denoting the deviation from the Gaussian distribution. By considering several constant-$w$ dark energy models ($w=-0.5, -0.8, -0.9, -1.0, -1.2$), we present the $w$-dependence of the time evolution of non-Gaussianity of the density distribution.   First we compare the results for the weakly nonlinear region at $z=2$ obtained by N-body simulations and Lagrangian linear perturbation by choosing the same smoothing length $R = 8 h^{-1}$ Mpc. In this region where non-Gaussianity become smaller as $w$ becomes larger, we regarded the results by Lagrangian linear perturbation as reliable in the previous work. As a result, even though the dispersion obtained by Lagrangian linear perturbation agrees well with the ones obtained by N-body simulations until $z\\sim2$, non-Gaussianity tend to be underestimated by Lagrangian linear perturbation.  This result is attributed to the fact that second order Lagrangian perturbation is necessary to express second order Eulerian perturbation \\cite{Grinstein87,Bouchet92}, mathematically and the peculiar acceleration perpendicular to the peculiar velocity can not be picked up by Lagrangian linear perturbation, physically. Therefore, we learn the fact that ``Lagrangian linear perturbation is not necessarily better than second order Eulerian perturbation''. However, regardless of the underestimation of non-Gaussianity, the conclusions of the previous paper that the relative difference of non-Gaussianity between $w=-1.0$ and $w=-0.8$ or $w=-1.2$ is about  $1 \\%$ at $z=2$ for $R=8 h^{-1} \\mbox{Mpc}$ remains unchanged because the degree of the underestimation in Lagrangian linear perturbation is almost $w$-independent.  Next, we analyze the statistical quantities in the strongly nonlinear region where the previous perturbative approach is obviously invalid. We find that the relative difference of non-Gaussianity among several constant-$w$ models spread as $z\\to 0$, even though the growth of non-Gaussianity is not monotonous in that region and the order changes around $z=0.5$ except $w=-0.5$ model whose dispersion is not so large enough. For example, at $z=0$ the relative difference between non-Gaussianity of $w=-0.8$ model and that of $w=-1.0$ model is $0.67 \\%$ for the skewness and $1.2 \\%$ for the kurtosis, while at $z=2$ they are $-0.52 \\%$ for the skewness and  $-1.3 \\%$ for the kurtosis. Similarly, at $z=0$ the difference between non-Gaussianity of $w=-1.2$ model and that of $w=-1.0$ model is $0.25 \\%$ for the skewness and $-0.19 \\%$ for the kurtosis, while at $z=2$ they are $0.15 \\%$ for the skewness and $0.34 \\%$ for the kurtosis.  Making use of N-body simulations, we also analyze the smaller structure by setting the smoothing scale $R=2 h^{-1} \\mbox{Mpc}$, even though the nonlinearity in this scale at $z=0$ is so strong that the galaxy distribution is influenced by the process other than gravity. For this case, since we look at the smaller scale, the dispersion and non-Gaussianity of the density distribution is larger than that in the case of $R=8 h^{-1} \\mbox{Mpc}$ for a fixed $z$ and $w$. Relative difference of non-Gaussianity among several constant-$w$ models are also greater than the case of $R=8 h^{-1} \\mbox{Mpc}$.  For example, at $z=0$ the relative difference between non-Gaussianity of $w=-0.8$ model and that of $w=-1.0$ model is $3.0 \\%$ for the skewness and $4.5 \\%$ for the kurtosis. Similarly, the relative difference between the non-Gaussianity of $w=-1.2$ model and that of $w=-1.0$ model is $-2.4 \\%$ for the skewness and $-9.0 \\%$ for the kurtosis. Furthermore, on the contrary to the weakly nonlinear region, for a fixed $z$ we confirm non-Gaussianity become larger as $w$ becomes larger after $z=0.5$, including the case of $w=-0.5$.  This suggests that in the sufficiently nonlinear region, say $\\sigma > 2$, there exists some scaling between the higher order cumulants and the dispersion $<\\delta^n>_c \\propto \\sigma^{p(n)}$, where $p(n) < 2n-2$, at least inside the scope of our calculations. Academically, not limited to the constraint for the dark energy models, this result seems to worth investigating related with the nature of the nonlinear gravitational clustering in the expanding Universe.  It is necessary to mention the possibility of examining our results. Recently, several galaxy redshift survey projects have been progressing~\\cite{Hawkins03,Abazajian04,Gerke04}. In these projects, many galaxies within a region ($z<0.3$ for 2dF, $z<0.5$ for SDSS, and $0.7 < z < 1.4$ for DEEP2) have been observed. These projects show the latest distribution of galaxies, which form strongly nonlinear structure. Actually, the skewness and kurtosis of the distribution of SDSS Early Data Release had been computed \\cite{Szapudi:2001mh}. The results suggested the SDSS imaging data can enable to determine the skewness and kurtosis up to $1 \\%$ and less than $10 \\%$, respectively. Therefore, in future projects we can expect to determine the skewness and kurtosis with high accuracy. However, based on our error analysis, the variance of the non-Gaussianity turned out to be about $50 \\%$ for 10 samples and it does not increased even by increasing the number of samples up to 50. This suggests that our analysis, by itself, is still immature and further improvements are necessary for the practical use to constrain the property of dark energy.  As mentioned above, for smaller scale like $R=2 h^{-1} \\mbox{Mpc}$ now, the complex physical processes about galaxy formation affect the galaxy distribution. In order to use observational data such as the PSCz survey \\cite{Saunders:2000af} which is an all sky survey at the present epoch, it is necessary to create mock PSCz catalogs from our results making use of the PSCz selection function \\cite{Rowan-Robinson:1999mx}. Even though this is from excessively academical interest, $w$-dependence of the non-Gaussianity of the actual galaxy distribution seems to be worth investigating. On the other hand, it is known that weak lensing surveys can potentially provide us with precision maps of the projected density up to redshifts around $1$ \\cite{vanWaerbeke:2000rm,Bacon:2000sy,Wittman:2000tc,Kaiser:2000if}. Even though we need another step of obtaining the convergence field which can be written as the projection of the matter density along the line of sight, the skewness and kurtosis of the convergence field can be tested by weak lensing surveys. Obviously, our results obtained in this paper help to derive these quantities.  Finally, it is worth noting that primordial non-Gaussianity which is generated by inflation \\cite{BKMR}. From WMAP observations, a limit for the non-linear coupling parameter has been established \\cite{Komatsu:2003fd} which does not deny the existence of non-Gaussianity in the primordial density fluctuations. Even though we cannot disentangle them generally, we can calculate non-Gaussianity of large scale structure for a given non-Gaussian initial condition and compare the results with the one obtained by the Gaussian initial condition. This requires further investigations and we hope to report results in a separate publication.      \\ack We thank to Masahiro Morikawa for useful discussion. For usage of COSMICS and $P^3M$ codes, we would like to thank Edmund Bertschinger and Alexander Shirokov. SM is supported in part by the Japan Society for Promotion of Science (JSPS) Research Fellowship."
    },
    "0608/astro-ph0608372_arXiv.txt": {
        "abstract": "Persistent high-energy emission of magnetars is produced by a plasma corona around the neutron star, with total energy output of $\\sim 10^{36}$~erg/s. The corona forms as a result of occasional starquakes that twist the external magnetic field of the star and induce electric currents in the closed magnetosphere. Once twisted, the magnetosphere cannot untwist immediately because of its self-induction. The self-induction electric field lifts particles from the stellar surface, accelerates them, and initiates avalanches of pair creation in the magnetosphere. The created plasma corona maintains the electric current demanded by ${\\rm curl}\\bB$ and regulates the self-induction e.m.f. by screening. This corona persists in dynamic equilibrium: it is continually lost to the stellar surface on the light-crossing time $\\sim 10^{-4}$~s and replenished with new particles. In essence, the twisted magnetosphere acts as an accelerator that converts the toroidal field energy to particle kinetic energy. The voltage along the magnetic field lines is maintained near threshold for ignition of pair production, in the regime of self-organized criticality. The voltage is found to be about $\\sim 1$~GeV which is in agreement with the observed dissipation rate $\\sim 10^{36}$~erg/s. The coronal particles impact the solid crust, knock out protons, and regulate the column density of the hydrostatic atmosphere of the star. The transition layer between the atmosphere and the corona is the likely source of the observed 100~keV emission from magnetars. The corona also emits curvature radiation up to $10^{14}$~Hz and can supply the observed IR-optical luminosity. ",
        "introduction": "\\label{intro}  At least 10\\% of all neutron stars are born as magnetars, with magnetic fields $B>10^{14}$~G. Their activity is powered by the decay of the ultrastrong field and lasts about $10^4$~years. They are observed at this active stage as either Soft Gamma Repeaters (SGRs) or Anomalous X-ray Pulsars (AXPs) \\cite{WT06}.  Besides the sporadic X-ray outbursts, a second, persistent, form of activity has been discovered by studying the emission spectra of magnetars. Until recently, the spectrum was known to have a thermal component with temperature $\\kB T\\sim 0.5$~keV, interpreted as blackbody emission from the star's surface. The soft X-ray spectrum also showed a tail at $2-10$~keV with photon index $\\Gamma=2-4$. This deviation from pure surface emission already suggested that energy is partially released above the star's surface. Recently, observations by RXTE and INTEGRAL have revealed even more intriguing feature: magnetars are bright, persistent sources of 100 keV X-rays \\cite{Kuiper06}. This high-energy emission forms a separate component of the magnetar spectrum. It becomes dominant above 10~keV, has a very hard spectrum, $\\Gamma\\simeq 1$, and peaks above 100~keV where the spectrum is unknown. Its luminosity, $L\\sim 10^{36}$~erg~s$^{-1}$, even exceeds the thermal luminosity from the star's surface. The observed hard X-rays can be emitted only in the exterior of the star and demonstrate the presence of an active plasma corona.  This corona can form as a result of starquakes that shear the magnetar crust and its external magnetic field. The persistence of magnetospheric twists can explain several observed properties of magnetars. Thompson, Lyutikov, \\& Kulkarni~\\cite{TLK} investigated the observational consequences using a static, force-free model, idealizing the magnetosphere as a globally twisted dipole. They showed that a twist affects the spindown rate of the neutron star: it causes the magnetosphere to flare out slightly from a pure dipole configuration, thereby increasing the braking torque acting on the star. Although the calculations of force-free configurations are independent of the plasma behavior in the magnetosphere, they rely on the presence of plasma that can conduct the required current $\\bj_B=(c/4\\pi)\\nabla\\times\\bB$. This requires a minimum particle density $n_c=j_B/ec$. Thompson et al.~\\cite{TLK} showed that even this minimum density can modify the stellar radiation by multiple resonant scattering.  The problem of plasma dynamics in the closed twisted magnetosphere has been formulated in \\cite{BT}, and a simple solution has been found to this problem. It allows one to understand the observed energy output of the magnetar corona and the evolution of magnetic twists. It is worth mentioning that plasma behavior around neutron stars has been studied for decades in the context of radio pulsars, and that problem remains unsolved. The principle difference with radio pulsars is that their activity is caused by {\\it rotation}, and dissipation occurs on {\\it open magnetic lines} that connect the star to its light cylinder. Electron-positron pairs are then created on open field lines \\cite{RS75,AS79}. By contrast, the formation of a corona around a magnetar does not depend on its rotation. All observed magnetars are slow rotators ($\\Omega=2\\pi/P\\sim 1$~Hz), and their rotational energy is unable to feed the observed coronal emission. Practically the entire plasma corona is immersed in the {\\it closed} magnetosphere and its heating must be caused by some form of dissipation on the closed field lines. It is this closure that facilitates the corona problem by providing both boundary conditions at the two footpoints of a field line.  The basic questions that one would like to answer are as follows. How is the magnetosphere populated with plasma? The neutron star surface has a temperature $k_{\\rm B}T\\simlt 1$~keV and the scale-height of its atmospheric layer (if it exists) is only a few cm.  How is the plasma supplied above this layer and what type of particles populate the corona? How is the corona heated, and what are the typical energies of the particles? If the corona conducts a current associated with a magnetic twist $\\nabla\\times\\bB\\neq 0$, how rapid is the dissipation of this current, i.e. what is the lifetime of the twist? Does its decay imply the disappearance of the corona?  An outline of the proposed model is as follows. The key agent in corona formation is an electric field $E_\\parallel$ parallel to the magnetic field. $E_\\parallel$ is generated by the self-induction of the gradually decaying current and in essence measures the rate of the decay. It determines the heating rate of the corona via Joule dissipation. If $E_\\parallel=0$ then the corona is not heated and, being in contact with the cool stellar surface, it will have to condense to a thin surface layer with $k_{\\rm B}T\\simlt 1$~keV. The current-carrying particles cannot flow upward against gravity unless a force $eE_\\parallel$ drives it. On the other hand, when $E_\\parallel$ exceeds a critical value, $e^\\pm$ avalanches are triggered in the magnetosphere, and the created pairs screen the electric field. This leads to a ``bottleneck'' for the decay of a magnetic twist, which implies a slow decay.  Maintenance of the corona and the slow decay of the magnetic twist are intimately related because both are governed by $E_\\parallel$. In order to find $E_\\parallel$, one can use Gauss' law $\\nabla\\cdot\\bE=4\\pi\\rho$ where $\\rho$ is the net charge density of the coronal plasma. This constraint implies that $\\bE$ and $\\rho$ must be found self-consistently. The problem turns out to be similar to the classical double-layer problem of plasma physics with a new essential ingredient: $e^\\pm$ creation.  A direct numerical experiment can be designed that simulates the creation and behavior of the plasma in the magnetosphere. The experiment shows how the plasma and electric field self-organize to maintain the time-average magnetospheric current $\\bar{\\bj}=\\bj_B$ demanded by the magnetic field, $(4\\pi/c)\\bar{\\bj}=\\nabla\\times\\bB$. The electric current admits no steady state on short timescales and keeps fluctuating, producing $e^\\pm$ avalanches. This state may be described as a self-organized criticality. Pair creation is found to provide a robust mechanism for limiting the voltage along the magnetic lines to $e\\Phi_e\\simlt 1$~GeV and regulate the observed luminosity of the corona. ",
        "conclusions": ""
    },
    "0608/astro-ph0608658_arXiv.txt": {
        "abstract": "{}{We investigate the effect of including a proper energy balance on the interaction of a low-mass planet with a protoplanetary disk.}{We use a three-dimensional version of the RODEO method to perform hydrodynamical simulations including the energy equation. Radiation is included in the flux-limited diffusion approach.}{The sign of the torque is sensitive to the ability of the disk to radiate away the energy generated in the immediate surroundings of the planet. In the case of high opacity, corresponding to the dense inner regions of protoplanetary disks, migration is directed \\emph{outward}, instead of the usual inward migration that was found in locally isothermal disks. For low values of the opacity we recover inward migration and show that torques originating in the coorbital region are responsible for the change in migration direction.}{} ",
        "introduction": "Migration has become a standard ingredient in planet formation theory, providing an explanation for the extrasolar giant planets that were found orbiting very close to their central star, the so-called Hot Jupiters. It is believed that these planets were formed at a distance of several Astronomical Units (AU), after which they moved in due to tidal interaction with the surrounding protoplanetary disk.  Depending on the mass of the planet, three different migration modes can be distinguished: \\begin{enumerate} \\item Type I migration \\citep{1979ApJ...233..857G,1997Icar..126..261W}, valid for low-mass planets. The disk response is linear in this case, and the resulting torques can be calculated in a semi-analytic way \\citep{2002ApJ...565.1257T}. Planets that are less massive than a few Earth masses are subject to this mode of migration. \\item Type II migration, for which the planet orbits in a deep annular gap \\citep{1986ApJ...309..846L}. In this case, the planet is tidally locked inside the gap, and the time scale for inward migration equals the viscous time scale, which approximately equals $10^6$ yr. This mode of migration holds for planets comparable to or more massive than Jupiter. See, however, \\cite{2002ApJ...572..566R} on gap formation for   low-mass planets. \\item In between these two regimes, interesting things happen. The onset of non-linear behavior can influence the torque to a large extent \\citep{2006astro.ph..7155M}. However, when the disk is massive enough, the planet enters a new migration mode that can be very fast and that relies on dynamical corotation torques \\citep{2003ApJ...588..494M,2004ASPC..324...39A}. This is called Type III migration, and it holds for planets that open up a partial gap. \\end{enumerate} In this Letter, we focus on deeply embedded low-mass planets and, therefore, on the regime of Type I migration.  The time scale for Type I migration is inversely proportional to the mass of the disk and the planet \\citep{2002ApJ...565.1257T}, and it can be much shorter than the disk lifetime of approximately $10^7$ yr. A planet of 1 Earth mass (1 \\mearth), for example, embedded in the minimum mass solar nebula (MMSN) at 5 AU, would migrate into the central star within $10^6$ yr. Within the widely accepted core-accretion model of giant planet formation \\citep{1996Icar..124...62P}, the gas giants also pass through a stage where the core is a few \\mearth, which thus should disappear into the central star even faster. Therefore, some sort of stopping mechanism is required for planets to survive Type I migration.  Proposed stopping mechanisms include magnetic turbulence \\citep{2004MNRAS.350..849N}, a toroidal magnetic field \\citep{2003MNRAS.341.1157T}, and an inner hole in the disk \\citep{2002ApJ...574L..87K,2006ApJ...642..478M}. However, most analytical and numerical work on planet migration has focused on disks with a fixed radial temperature profile. For these disks, hydrodynamical simulations can reproduce Type I migration in both two and three dimensions \\citep{2002A&A...385..647D, 2003ApJ...586..540D, 2003MNRAS.341..213B}.  The isothermal assumption is only valid if the disk can radiate away all excess energy efficiently. When the disk is optically thick, the radiative energy flux through the surface of a sphere of radius $H$ around the planet is given by $F_\\mathrm{R}=\\sigma T^4/\\tau$, where $\\sigma$ is the Stefan-Boltzmann constant and $\\tau$ is the optical depth over a distance $H$. Therefore, when the opacity increases, less energy can be radiated away. This makes the cooling time scale much longer than the dynamical time scale in the dense inner regions of protoplanetary disks, and therefore the isothermal assumption is likely to be invalid there.  One may wonder what effects releasing the isothermal assumption may have on the torque. The effects of a more realistic temperature profile on the Lindblad torques were investigated by \\cite{2005ApJ...619.1123J} and \\cite{2004ApJ...606..520M}, who found that the migration rate is very sensitive to the temperature and opacity structure and that in some cases migration may be very slow compared to the isothermal Type I case. In this Letter, we present the first results of including a detailed energy balance on the migration behavior of low-mass planets in a radiation-hydrodynamical context. In Sect. \\ref{secMod} we discuss the disk model, in Sect. \\ref{secRes} we present the results, and we give a discussion of the results in Sect. \\ref{secDis}. We then conclude in Sect. \\ref{secCon}. ",
        "conclusions": "\\label{secCon} In this Letter, we have presented the first radiation-hydrodynamical simulations of low-mass planets embedded in a protoplanetary disk. We find that the migration behavior of a 5 \\mearths planet depends critically on the ability of the disk to radiate away the heat generated by compression close to the planet. When radiative cooling is not efficient, the torque on the planet is \\emph{positive}, indicating outward migration. Analysis of the torque distribution showed that it is the coorbital region that generates the large positive contribution to the torque, while the Lindblad torque is reduced in strength, but is still negative. Local, adiabatic simulations showed that compression at the turn-over point of the horseshoe orbits creates a warm region behind the planet. As the disk tries to maintain pressure equilibrium, the density behind the planet will become lower, which leads to a positive corotation torque. This happens in the dense inner parts of protoplanetary disks, where the opacity is high. In regions of low opacity, we recover inward Type I migration. Therefore, low-mass planets do not migrate inward all the way to the central star, bun instead stop migrating when radiative cooling becomes inefficient. In the MMSN, the transition occurs at approximately 15 AU. As the disk slowly accretes onto the central star, this radius will move inward; as a result, low-mass planets will also eventually continue their inward migration, but on the much longer viscous time scale. This way, low-mass planets are saved from fast inward Type I migration into the central star."
    },
    "0608/astro-ph0608615.txt": {
        "abstract": "{} {Planet formation has been known for many years to be tied to the spatial distribution of gas and dust in disks around young stars. To constrain planet formation models, imaging observations of protoplanetary disks are required. } { In this framework, we have undertaken a mid-infrared imaging survey of Herbig Ae stars, which are pre-main sequence stars of intermediate mass still surrounded by a large amount of circumstellar material. The observations were made at a wavelength of 20.5 $\\mu$m with the CAMIRAS camera mounted at the Cassegrain focus of the Canada France Hawaii Telescope.} { We report the observations of three stars, HD135344, CQTau and HD163296. The circumstellar material around the three objects is spatially resolved. The extensions feature a disk like shape. The images provide direct information on two key parameters of the disk : its inclination and its outer radius. The outer radius is found to be quite different from the one deduced from disk models only constrained by fitting the Spectral Energy Distribution of the object. Other parameters of the disk, such as flaring, dust mass have been deduced from fitting both the observed extension and the spectral energy distribution with sophisticated disk models. } {Our results show how important imaging data are to tighten constraints on the disk model parameters.} %{\\bf attention on parle de la masse de petits grains ici}The disk dust mass found here is quite different from previously calculated : 10 times lighter in the case of CQ~Tau, %5 times lighter for HD~163296, twice heavier in the case of HD~135344. These results show how important imaging data % are to tighten constraints on the disk model parameters.  %All three objects are spatially resolved.  These %extensions feature a disk like shape. We compare the observed %spatial distribution at 20.5 $\\mu$m with the predictions of a disk %model based on a 2-D radiative transfer code in which the gas %and dust phases are well mixed. It is demonstrated how imaging data %at 20.5 $\\mu$m as those presented in this paper are valuable to determine a %minimum outer radius and an inclination of the disk and definitively tighten the constraints %on the model parameters. ",
        "introduction": "The formation of circumstellar disks is a natural outcome of the star formation process by which a molecular core collapses to form a star \\citep{SHU87}.  Circumstellar disks can outlive the period during which stars form and still be present when the star is in its Pre-Main-Sequence (PMS) phase. In these disks, composed of gas and dust, various physical processes can lead to the growth of dust grains and eventually to the formation of planets.  Understanding the physical conditions that prevail in these objects is of crucial importance when studying planet formation.\\\\ The study of circumstellar disks is a field in fast development both from the observational and the modeling point of view (e.g. \\citet{NATTA04} and references there-in). The Infrared Space Observatory (ISO) has given clues on the dust composition of a sample of isolated HAeBe systems \\citep{BOUW01, MEEUS01}.  While these spectra reveal the composition of the dust, no direct information concerning the spatial distribution of the different dust species can be inferred from the ISO data. Most studies so far have used the Spectral Energy Distributions (SEDs) to put constraints on the spatial distribution of the circumstellar material.  %Although great progress had been made in modeling the disk structure with radiative transfer models able to reproduce the %SEDs (Chiang, E. I., \\& Goldreich, P. 1997, ApJ, 490, 368; Natta, %A., Prusti, T., Neri, R., et al. 2001, A\\&A, 371, 186; Dullemond, %C.P., \\& Dominik, C., \\& Natta, A. 2001, ApJ, 560, 957), the structure %of the disks is not uniquely constrained.   Models of protoplanetary disks are increasingly successful at accounting for much of the observed properties. For instance, they can justify that disks SEDs are generally rather flat in $\\nu F_{\\nu}$, where $\\nu$ is the frequency and $F_{\\nu}$ the flux \\citep{KEN87}. Furthermore, models can explain that dust features are almost all seen in emission (\\citet{CAL91}, \\citet{CG97}, hereafter CG97), the presence of a near-infrared excess in the SEDs of Herbig Ae stars (\\citet{NATTA01}, \\citet{DDN01}, hereafter DDN01) and interpret the differences observed in the far-IR excesses \\citep{DD02, DD04a}. Fitting the SED only allow to make conjectures on how the disks' geometry looks like; spatially resolved imaging data of those disks are absolutely necessary to verify theories and models' assumptions. % But, spatially resolved observations are also needed to further %constrain the geometry of the disks. For instance, key parameters, such as the disk surface density profile with radius, are still very poor constrained when fitting SEDs.\\\\ Mid-infrared imaging observations from a large ground-based telescope are potentially well suited to bring spatial information on disk around Herbig Ae (HAe) stars. HAe stars represent the middle stage of PMS evolution of intermediate-mass stars ($\\sim$ 2-3 $M_{\\sun}$); they are bright enough to heat sub-micron dust grains at 100 AU to a temperature of about 150 K. %up to an equilibrium temperature around 150 K sub-micron size dust grains located at a typical distance of 100 AU from the star. Grains at such a temperature have their peak of thermal emission in the mid-InfraRed (mid-IR). The diffraction limited angular resolution achievable with a 3.6 meters class telescope in the mid-IR, 0.6/1.2 arcsec at 10/20 $\\mu$m, corresponds to a distance of 60/120 AU for a star located at a typical distance of 100 pc and thus allows a relatively good sampling of disks whose sizes range in the several hundreds of AU. The first attempts to resolve the spatial structure of the circumstellar material around HAe stars were performed with multi-aperture observations with a single bolometer; emission extending up to large distances from the star (more than 5 arcsec) were found around 3 objects HD97048, HD97300 and HD176386 \\citep{PRUSTI94}. Such extended emissions were attributed to the emission from large molecules and small grains transiently heated by star-light and distributed in a dust shell with a large inner radius, probably a remnant from the cloud from which the star was born. It is only with the advent of mid-IR cameras that the full potential of mid-IR observations to study disk structures was achieved. Single dish observations have revealed the disk structure in the 500 AU (Astronomical Units) range around two HAe stars: AB Aur \\citep{MAR95, PAN05} and HD100546 \\citep{GRA01, LIU03}. Interferometric observations have allowed to probe the innermost regions of the disk (1-10 AU) and numerous objects have been observed \\citep{MILLAN01, TUTHILL02, Wilkin, LEINERT04, LIU05}. When mid-IR single dish observations have now clearly demonstrated on a few examples that they can provide unique information on the disk structure, the constrains brought on the disk modeling by such spatial information has not yet been fully exploited. %Several objects such as AB Aur, HD100546, HD~135344, %HD97048, HD97300 and HD176386 have been already resolved in %mid-infrared (\\citet{MARSH95}, \\citet{GRA01}, \\citet{JAYA01}, \\cite).  Prusti et %al. (1994) modeled the SED between 10 and 20 $\\mu$m with small grains %in dust shell with large inner radius, taken into account transiently %heated PAHs. Others authors have imaged objects like HD100546 and AB %Aur at 10 and 20 $\\mu$m in which a spherical cloud and a disk are %detected; but they did not model the systems (?? references ???).  A pilot study, %conducted in 1999 with the Saclay mid-Infrared camera CAMIRAS, has %demonstrated the potential of the mid-IR to study disk: the dust disk %around the HAe star AB Aur has already been resolved with this %instrument (Pantin, Bouwman \\& Lagage 2005), showing an elliptical %ring-like emission structure at 280 AU from the star which could be %explained by very small grains transiently heated by star light.\\\\  In this paper, we were interested in both increasing the number of HAe for which information on the mid-IR spatial extension is available, and in using this information, in combination with the already existing SED measurements between 1 and 100 $\\mu$m, to constrain the parameters of the disk models developed these last years. The paper is organized as follows : the observations and data reduction are described in Sect. 2. Sect. 3 deals with the results, in terms of spatial extensions. In Sect. 4, a first, simple approach is used to establish the disk inclination, and in Sect. 5, we describe the 2-D radiative transfer code used to reproduce the observations. A discussion of the results follows in Sect. 6. Conclusions and perspective are drawn in Sect. 7. ",
        "conclusions": "We have shown the strength of mid-IR imaging to constrain the disk properties. There are still a very limited number of objects with extended emission spatially resolved and it is not yet possible to draw statistical conclusions about the spatial structure of disk around HAe stars. With the advance of mid-IR instruments on 8 meter class telescope, such as the VISIR \\citep{LAGA04} instrument available on the MELIPAL Very Large Telescope (VLT) at the European Southern Observatory (ESO), higher angular resolution will be available and the field will further develop; more quantitative studies will be possible and the goal of retrieving detailed disk surface density profiles from the observations should be achievable.\\\\ Observations should not be limited to the 20 $\\mu$m atmospheric window; observations in the 10 $\\mu$m atmospheric window are also very promising, especially for those HAe stars whose spectrum features the so-called PAH bands features at 3.3, 6.2, 7.7, 8.6, 11.2, 12.7 $\\mu$m. PAH bands are attributed to vibrational relaxation of UV-pumped Polycyclic Aromatic Hydrocarbon molecules containing about 50-100 carbon atoms \\citep{ALL89, PUGET89}. Their emission, as a function of the distance $r$ to the star, drops with a $r^{-2}$ power-law, much slower than thermal emission from large grains in thermal equilibrium. PAHs emission is thus a promising probe to study flaring disks at large distances from the star, with the good angular resolution achieved now on large ground-based telescopes at 10 microns. Fitting by sophisticated models a combination of interferometric observations in the near- and mid-IR, which probe the inner-most disk regions, with single dish observations in the near and mid-IR which probe intermediate disk regions, is clearly the way to clear up the field in the next few years. \\\\ %\\begin{figure}[!b] % \\resizebox{\\hsize}{!}{\\includegraphics[scale=0.9,angle=0]{residuals_HD~135344.ps}} % \\caption{The residuals of HD~135344 and % the model associated with the parameters found on the Table.2.}\\label{residual_hd135344} % \\end{figure}  \\clearpage{\\pagestyle{empty}}  \\newpage \\begin{figure}[!t] \\resizebox{\\hsize}{!}{\\includegraphics[angle=0,scale=0.48]{hd163296_camiras2.ps}} \\caption{HD~163296 extended emission (central point source removed). West is up, North on the left. The pixel size is 0.29 arcsec. The disk has a surface brightness of 0.59 Jy/$^{\"2}$ (S/N=27 for $\\sigma \\sim$ 0.002 Jy) for the brightness part in the direction east/west and 0.23 Jy/$^{\"2}$ (S/N=11) for the less bright part in the direction north/south. The noise is calculated using a sigma-clipping technique. %The total flux of the disk is 2 $\\pm$ 0.5 Jy. }\\label{hd163296} \\end{figure}  \\begin{figure}[!h] \\resizebox{\\hsize}{!}{\\includegraphics[scale=0.1]{cqtau_camiras2.ps}} \\caption{CQ~Tau extended emission (central point source removed). The same orientation as Fig~\\ref{hd163296}. The pixel size is 0.29 arcsec. At 0.9\" from the center, the surface brightness is 0.94 Jy/$^{\"2}$ (S/N=15 for $\\sigma \\sim$ 0.004 Jy). At 2\" from the center, the surface brightness is 0.23 Jy/$^{\"2}$ (S/N=5). %The total flux of the disk is 5.9 $\\pm$ 0.3 Jy. }\\label{CQtau} \\end{figure}  \\begin{figure}[!h] \\resizebox{\\hsize}{!}{\\includegraphics[scale=0.5,angle=0]{hd135344_camiras2.ps}} \\caption{HD~135344 extended emission (central point source removed).The same orientation as Fig~\\ref{hd163296}. The pixel size is 0.29 arcsec. At 0.9\" from the center, the surface brightness is 0.25 Jy/$^{\"2}$ (S/N=15 for $\\sigma \\sim$ 0.0015 Jy). %The total flux of the disk is 0.9 $\\pm$ 0.3 Jy. }\\label{hd135344} \\end{figure}   \\emph{Acknowledgments.} We are gratefully indebted to P. Masse, R. Jouan and M. Lortholary for their assistance with CAMIRAS instrument, A. Claret in efficiently supporting us in our observations, as well as to the staff of CFHT/Hawaii for their support during the observing runs. CD wishes to thank J. Bouwman and E. Habart for very helpful discussions and advice."
    },
    "0608/astro-ph0608002_arXiv.txt": {
        "abstract": "We note that galaxy-galaxy lensing by non-spherical galaxy halos produces a net anti-correlation between the shear of background galaxies and the ellipticity of foreground galaxies. This anti-correlation would contaminate the tomographic cosmological weak lensing two point function if the effect were not taken into account. We compare the size of the galaxy-galaxy lensing contribution to the change in the cosmic shear two-point cross-correlation function due to a change in the dark energy equation of state $w$ of 1\\%. We find them comparable on scales $\\lesssim 5'$ for NFW galaxy profiles, and out to much larger scales for SIE profiles. However the galaxy-galaxy lensing signal has a characteristic spatial and redshift pattern which should allow it to be removed. ",
        "introduction": "\\label{sec:introduction}  Cosmic shear shows great promise for testing the cosmological model and measuring cosmological parameters. It measures the gravitational bending (lensing) of light by all the intervening mass in the Universe. In cosmic shear, distant (background) galaxies provide a convenient screen of objects with potentially simple statistical properties, from which the light bending can be inferred. The distortion depends on the lens geometry and thus on the curvature and expansion history of the Universe; it also depends on the distribution of matter, which itself depends on most aspects of the cosmological model.  Cosmic shear was first detected in 2000 \\citep{wittmanea00,baconre00,vanwaerbekeea00,kaiserwl00} and has been used to constrain cosmological parameters in subsequent surveys \\citep[recently][]{hoekstraea05}. Cosmic shear has the potential to become the most powerful probe of cosmology because it observes a non-Gaussian three dimensional field in the local Universe, where the mysterious dark energy dominates. However a number of details inevitably remain to be investigated further whilst planning for future experiments.  The simplest statistic is the two point angular correlation function \\citep[or integrals of this quantity such as the power spectrum or aperture mass,][]{schneider96}. If the background galaxies can be separated into populations at different redshifts the power of this statistic is greatly increased by cross-correlating galaxies in different redshift bins \\citep{hu99}. This technique will dominate for future surveys which aim to measure the equation of state of the dark energy, or determine a new gravitational theory.  Cosmic shear is particularly simple if distant background galaxies can be assumed to have random orientations. The distortion by gravitational lensing can then be extracted statistically, for example by averaging over galaxy orientations in a given patch of sky. Unfortunately it is unlikely to be straightforward since when galaxies form they tend to align pointing towards dark matter concentrations due to tidal interactions. This leads to two complicating effects: (i) neighbouring galaxies at a given redshift are intrinsically aligned \\citep{hoyle49,peebles69,heavensp88,crittendennpt01} and (ii) a pair of galaxies at two different redshifts have correlated observed ellipticities because a dark matter concentration close to the nearer object may tidally align the nearer object \\emph{and} simultaneously gravitationally lens the more distant object \\citep{hiratas04}.  These intrinsic alignment effects are complicated because they require an understanding of tidal alignments of galaxies, however they do have a distinctive signature and can therefore be removed without requiring a detailed understanding of galaxy formation \\citep{kings02,kings03,heymansh03,king05}.  In this paper we present another complicating effect which is much simpler than the intrinsic alignment effects: cross-correlation of a pair of galaxies at different redshifts produces a contribution to the usual cosmic shear two-point statistic even if there are no tidal effects. This is because the nearer galaxy will gravitationally lens the more distant galaxy (galaxy-galaxy lensing). This effect produces a net anti-correlation in the realistic case when the nearer galaxy (i) is not circularly symmetric and (ii) has a correlation between the asymmetry of the light and the mass. In this paper we present this new factor and quantify its effect for elliptical dark matter halos for which mass and light have the same ellipticity and orientation.  First we review ellipticity and shear correlation function notation. We then illustrate the galaxy-galaxy lensing contribution to the ellipticity correlation function and calculate the effect due to a population of galaxy lenses at a fixed redshift. We investigate how this effect depends on the source and lens redshift. Finally we compare this to the size of the tomographic cosmic shear cross-correlation signal. ",
        "conclusions": "\\begin{figure} \\epsfig{file=f4.ps,width=8.5cm,angle=0} \\caption{ \\emph{Solid lines:} Dependence of the galaxy-galaxy lensing shear-ellipticity correlation at 4 arcmin on lens redshift for source redshifts 0.4, 0.8 and 1.2 (lower, middle and upper lines respectively). \\emph{Dashed lines:} Cosmic shear cross-correlation between a sample at a redshift $z_l$ and samples at redshifts $z_s= 0.4$, 0.8 and 1.2 (lower, middle and upper lines respectively); for a change in $w$ of 0.01 (the result for $w=-0.99$ is subtracted from the result for $w=-1$). } \\label{fig:z dep} \\end{figure}    In Figs.~\\ref{fig:cf cosmic shear} and~\\ref{fig:z dep} we compare the galaxy-galaxy lensing contribution to the ellipticity correlation function with that from cosmic shear cross-correlation tomography. We calculate the non-linear matter power spectrum as a function of redshift using CAMB~\\citep{lewiscl00} with the HaloFit \\citep{smithea03} option. This is converted into the cosmic shear cross power spectrum using equations from \\cite{hu99}. We assume delta function redshift distributions corresponding to the galaxy-galaxy lens and source redshifts.  The cosmic shear correlation function is flatter than the galaxy-galaxy lensing contribution Above $\\sim0.2'$ the galaxy-galaxy lensing contribution is smaller than the cosmic shear signal for NFW lenses. However it must be much smaller than the cosmic shear signal if it is not to interfere with cosmological parameter analyses.  To illustrate this we also show the change in cosmic shear signal for a small change ($\\Delta w=0.01$) in the equation of state of dark energy, $w$, leaving the other cosmological parameters (including $\\sigma_8$) fixed. The galaxy-galaxy lensing contribution is not negligible below 4 arcmin for NFW galaxy profiles. We find that this scale will strongly depend on the lens profile since it is greater than 30 arcmin for SIE profiles.  We assumed that the lens ellipticity is the same size and orientation as the mass ellipticity. This is consistent with measurement attempts \\citep{hoekstrayg04,mandelbaumea06}, but the signal to noise is low. If the mass ellipticity were a factor $f$ smaller than the light then this would simply scale the galaxy-galaxy lensing contribution by $f$. If there is a misalignment then the effective $f$ on stacking would be a first approximation to the change. \\cite{heymanswhvv06} consider a more sophisticated alignment model which can reduce the signal by a factor of roughly three.  It would be interesting to see whether the contribution is increased or reduced on considering clustering of the lens galaxies. Further, it is not clear whether deviations from ellipsoidal symmetry of the lens (such as substructures) would average out.  A rough comparison of Fig.~\\ref{fig:cf cosmic shear} with \\cite{heymanswhvv06} (Table 2) indicates that at 1 arcminute the galaxy-galaxy NFW lensing contribution is a significant fraction of the total shear-intrinsic alignment correlation, whereas at 10 arcmin the shear-intrinsic alignment correlation is dominated by tidal alignments. Note that this comparison is only approximate since \\cite{heymanswhvv06} integrate over lens redshift.  Fortunately it may be relatively easy to remove the galaxy-galaxy lensing contribution using the characteristic tangential nature of galaxy-galaxy lensing. For example the ellipticity two-point function could be measured as a function of two dimensional separation $(\\theta_1, \\theta_2)$, where the position angle $\\alpha$ is measured from the major axis of the foreground lens. \\citep[This ``stack and rotate\" method was advocated by][to measure ellipticity of galaxy dark matter halos.]{natarjanr00} The galaxy ellipticity and profile could then be fitted simultaneously with extracting the cosmic shear two point function. Alternatively one could use the method of \\cite{king05}, which uses the redshift signature of the effect. Also \\cite{heymanswhvv06} suggest removing the most luminous galaxies in the lower redshift slice.  This effect should also ultimately be taken into account for higher order statistics of cosmic shear, since they probe smaller scales than the two-point statistic \\citep[see ][for a calculation for the three point function]{how01}. Furthermore the galaxy-galaxy flexion signal could be a bigger contaminant of cosmic flexion because the latter is more sensitive than cosmic shear at small angular scales."
    },
    "0608/astro-ph0608528_arXiv.txt": {
        "abstract": "The Milky Way satellite dwarf spheroidal (dSph) galaxies are the smallest dark matter dominated systems in the universe.  We have underway dynamical studies of the dSph to quantify the shortest scale lengths on which Dark Matter is distributed, the range of Dark Matter central densities, and the density profile(s) of DM on small scales. Current results suggest some surprises: the central DM density profile is typically cored, not cusped, with scale sizes never less than a few hundred pc; the central densities are typically $10-20$GeV/cc; no galaxy is found with a dark mass halo less massive than $\\sim5.10^7M_{\\odot}$. We are discovering many more dSphs, which we are analysing to test the generality of these results. ",
        "introduction": "Determining the nature of the dark matter is one of the key goals of contemporary astronomy.  The extremely large mass to light ratios for certain Local Group dwarf spheroidal (dSph) galaxies suggest that these are the most dark matter dominated stellar systems known in the Universe. Given the apparent absence of dark matter in globular clusters (length scale $\\sim10$pc), and the direct evidence that there is no dark matter associated with the galactic thin disk (length scale $\\sim 100$pc; \\cite{KG89,KG91})  the dSphs (characteristic radii $\\sim300+$ pc) are both the smallest systems to contain dynamically significant quantities of dark matter, and the most nearby systems where we may look for characteristic properties of the distribution of dark matter (maximum density, etc ...), which may provide some knowledge of its nature.  Additionally, simulations of dark halo formation in a cosmological context predict that dSph galaxies are (unmodified?) survivors from the earliest structures formed, so that understanding their nature and structure also has general implications for galaxy formation and evolution.  Is it possible reliably to determine the dark matter distribution in a dSph galaxy? Fortunately, some dSphs are astrophysically simple, are old, are apparently in equilibrium, have no dynamically-significant gas, but contain a (small) number of stars, which provide ideal kinematic collisionless tracer particles.  With the new generation of wide-field multi-object spectrographs on 4-8m telescopes, it is now viable to obtain sufficient high-quality kinematic data to determine the gravitational potentials in all the Galactic satellites.  There are two types of analysis feasible with current datasets: a straighforward pressure-gravity balance analysis, based on the velocity dispersion moment, using the Jeans' equations, and an analysis of the full projected kinematic distribution function, using the full information set but requiring more sophisticated analyses. These analyses are made more robust, and simplified, by stellar population studies of the dSph galaxies. Ultra-deep HST imaging for Draco \\cite{Gr} and UMi \\cite{Wy} show their stellar populations are indeed indistinguishable from those of old metal-poor globular star clusters, and hence have  $M/L_V\\sim2$.  The Jeans' equations for a spherical stellar system lead directly to a model-independent mass estimator (see \\cite{BT} Eq. 4-55 \\& 4-56) requiring only knowledge of the velocity anisotropy and the true radial velocity dispersion $\\langle v_r^2 \\rangle$. Assuming spherical symmetry, it is straightforward to obtain $\\langle v_r^2 \\rangle$ from the line of sight velocity dispersion $\\langle v^2 \\rangle$ using Abel integrals. In applying the Jeans' equation analyses, the spatially-binned dispersion profile is fit by an appropriate smooth function (figure 1), while the light is fit by a Plummer law, as this is an excellent fit to available star count data in all available cases (cf \\cite{K01}).   \\begin{figure}[!h] \\begin{center} \\includegraphics[width=0.42\\textwidth]{gg-fig-dra1.ps} \\includegraphics[width=0.42\\textwidth]{gg-fig-dra2.ps} \\caption{Draco surface brightness profile, with a fitted Plummer model, and the observed line of sight velocity dispersion profile, from \\cite{K01}.} \\end{center} \\end{figure}  The Jeans' equation moment analysis involves two radial functions, the mass density and a possibly variable anisotropy in the stellar orbits, this second generating an anisotropic stress tensor. It is this anisotropic stress which gives the dSph galaxies their shapes, as direct observational limits exclude any angular momentum support against gravity. Equivalently, radially variable anisotropy is degenerate with mass (figure 2) making any deductions on the inner mass profile being cored or cusped in general model-dependent. Further information is needed to break this degeneracy, and fortunately is sometimes available, as we discuss below. In general however, full multi-component distribution-function models are required to use the information in the data to break this degeneracy.  \\begin{figure}[!h] \\begin{center} \\includegraphics[width=0.42\\textwidth]{gg-fig-aniso.eps} \\caption{An illustration of the possible effects of stellar orbital anistropy variations on kinematic modelling, from \\cite{W04}.} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0608/astro-ph0608234_arXiv.txt": {
        "abstract": "We present a grid of radiation transfer models of axisymmetric young stellar objects (YSOs), covering a wide range of stellar masses (from $0.1$\\,$\\msun$ to $50$\\,$\\msun$) and evolutionary stages (from the early envelope infall stage to the late disk-only stage). The grid consists of 20,000 YSO models, with spectral energy distributions (SEDs) and polarization spectra computed at ten viewing angles for each model, resulting in a total of 200,000 SEDs. We made a careful assessment of the theoretical and observational constraints on the physical conditions of disks and envelopes in YSOs, and have attempted to fully span the corresponding regions in parameter space. These models are publicly available on a dedicated WWW server\\,\\footnote{http://www.astro.wisc.edu/protostars}. In this paper we summarize the main features of our models, as well as the range of parameters explored. Having a large grid covering reasonable regions of parameter space allows us to shed light on many trends in near- and mid-IR observations of YSOs (such as changes in the spectral indices and colors of their SEDs), linking them with physical parameters (such as disk and infalling envelope parameters). In particular we examine the dependence of the spectral indices of the model SEDs on envelope accretion rate and disk mass. In addition, we show variations of spectral indices with stellar temperature, disk inner radius, and disk flaring power for a subset of disk-only models. We also examine how changing the wavelength range of data used to calculate spectral indices affects their values. We show sample color-color plots of the entire grid as well as simulated clusters at various distances with typical {\\it Spitzer Space Telescope} sensitivities. We find that young embedded sources generally occupy a large region of color-color space due to inclination and stellar temperature effects. Disk sources occupy a smaller region of color-color space, but overlap substantially with the region occupied by embedded sources, especially in the near- and mid-IR. We identify regions in color-color space where our models indicate that only sources at a given evolutionary stage should lie. We find that while near-IR (such as JHK) and mid-IR (such as IRAC) fluxes are useful in discriminating between stars and YSOs, and are useful for identifying very young sources, the addition of longer wavelength data such as \\mips is extremely valuable for determining the evolutionary stage of YSOs. ",
        "introduction": "\\label{s:intro} An explosion of infrared data on star formation regions is being collected by the {\\it Spitzer Space Telescope} and will continue to be collected by future missions such as the {\\it Herschel Space Observatory} and the {\\it James Webb Space Telescope}. Despite its modest size ($0.85$\\,m), the {\\it Spitzer Space Telescope} is much more sensitive and has higher spatial resolution than previous infrared observatories \\citep{werner04}. It is also very efficient, and has mapped a substantial fraction of the star formation regions in the Galaxy using the IRAC \\citep[3.6, 4.5, 5.8, and 8\\micronsns;][]{fazio04} and MIPS \\citep[24, 70, and 160\\micronsns;][]{rieke04} instruments. For example, the GLIMPSE \\citep{benjamin03} and MIPSGAL (PI Carey) Legacy surveys are mapping over 150 square degrees of the inner Galactic plane; the Cores-to-Disks (c2d) Legacy project has surveyed five nearby large molecular clouds \\citep{evans03}; and several projects are surveying the Orion molecular cloud \\citep{megeath_iau}, the Taurus molecular cloud \\citep{hartmann05}, the LMC \\citep{chu_letter,meixner06} and SMC, and many other well-known star formation regions \\citep[e.g][]{gutermuth04,megeath_etacam,allen_ppv,sicilia06}. These data can be combined with other surveys, such as the 2MASS near-infrared all-sky survey \\citep{skrutskie06}, expanding the wavelength coverage of the observed YSOs.  Our particular goal is to characterize YSOs in the Galaxy using the GLIMPSE and MIPSGAL surveys, as well as in the LMC using the SAGE survey, and to determine the timescales of various evolutionary stages as a function of stellar mass and location in the Galaxy. Ultimately, we hope to provide an independent estimate of the star formation rate and efficiency in the Galaxy and the LMC.  To help analyze the SEDs of YSOs we developed radiation transfer models \\citep{whitney03p2,whitney03p1,whitney04}. The well-known classification for low-mass YSOs uses the spectral index $\\alpha$ (the slope of $\\log_{10}{\\lambda F_\\lambda}$ vs $\\log_{10}{\\lambda}$ longward of $2$\\micronsns) to classify a source as embedded ($\\alpha > 0$; Class~I), a disk source ($-2 < \\alpha < 0$; Class~II) or a source with an optically thin or no disk ($\\alpha < -2$; Class~III) \\citep{lada87}. However, our models indicate that the spectral index and the colors of a YSO in certain wavelength ranges are not always directly related to its evolutionary stage. For example, an edge-on embedded protostar can have a decreasing slope in the narrow wavelength range of $2-10$\\microns if the flux is dominated by scattered light (edge-on disks exhibit similar behavior, \\citealt{wood02_hh30}, \\citealt{grosso03}). In addition, the temperature of the central source also affects the colors in this wavelength range, as does the location of the inner radius of the disk:  hot stellar sources and large inner disk holes can produce red colors in a star+disk source (see Section~\\ref{s:color}).  In an attempt to improve our physical understanding gained from interpreting the SEDs of the many sources found in these surveys, many of which only have JHK, IRAC, and \\mips data, we plan to fit observed SEDs using a large grid of pre-computed model SEDs. This grid attempts to encompass a large range of stellar masses and YSO evolutionary stages. We sampled the model parameters based on both theory and observations (see Section \\ref{s:params}). This paper describes our publicly available grid of models. A companion paper (Robitaille et al. 2006, in preparation) describes the method used to fit observed SEDs using the grid of models. The advantages of fitting pre-computed SEDs to data, even in the fairly narrow wavelength range mentioned above, are that 1) one makes use of all available data simultaneously without loss of information 2) the uniqueness or non-uniqueness of a fit is immediately apparent from the range of model parameters that can fit a given SED, and 3) it is an efficient technique when used with large datasets as the model SEDs do not have to be computed for each source. In Section~\\ref{s:models} we describe the grid of models. Section~\\ref{s:results} shows results from the grid, including sample SEDs, polarization spectra, and color-color plots, as well as an analysis of spectral index and color-color plot classifications. Finally, in Section~\\ref{s:conclusion} we make concluding remarks. ",
        "conclusions": "\\label{s:conclusion}  We have computed a large grid of SEDs from axisymmetric YSO models using a Monte-Carlo radiation transfer code. These models span a large range of evolutionary stages, from the deeply embedded protostars to stars surrounded only by optically thin disks.  We have made the 20,000 models publicly available on a dedicated web server\\footnote{http://www.astro.wisc.edu/protostars}. For each model, the following output are available for download: \\begin{itemize} \\item The SEDs for 10 inclinations (from pole-on to edge-on) and integrated in 50 different circular apertures (with radii between $100$ to $100,000$\\,AU). \\item Three separate SEDs, constructed from the energy packets whose last point of origin is the star, the disk, and the envelope respectively (see Section~\\ref{s:seds}), for each inclination and aperture. In addition, an SED constructed for each inclination and aperture from photons who last scattered before escaping, and an SED constructed for each inclination from photons which escaped directly from the stellar surface are available. \\item A polarization spectrum for each inclination and aperture. \\item Convolved fluxes and magnitudes for each inclination and aperture for a wide range of filters. \\end{itemize} Users have the option of downloading the convolved fluxes and the model parameters for all models as single files. These files can be used for example to carry out the analysis of color-color spaces or spectral indices not covered in this paper, or to compare the models to data directly.  We request feedback from the community, observers and theorists alike, on how to improve our grid of models, as we plan to run new grids of models over the next few years. For example, we plan to produce higher signal to noise SEDs in the next grid of models, as well as produce resolved images, in addition to the planned model improvements described in Section~\\ref{s:caveats}.  We have presented typical model SEDs from our grid (Figure~\\ref{f:seds}), including separate SEDs constructed from the energy packets whose last point of origin are the star, the disk and the envelope (Figure~\\ref{f:seds_indiv}). We have also presented polarization spectra for these models in UBVRI and JHKL bands (Figure~\\ref{f:seds_pol}). In addition, we have shown the variation of the K-band polarization for all of our models with disk mass and envelope accretion rate (Figure~\\ref{f:k_pol}). Stage~I sources show a high polarization ($>5$\\,\\%) for envelope accretion rates above $\\mdote/\\mstar\\sim2\\times10^{-6}$\\,yr$^{-1}$, over a large range of viewing angles ($45-90^{\\circ}$) and a broad wavelength range ($0.5-10$\\,\\micronsns). Stage~I sources can show a $90^{\\circ}$  position angle rotation with wavelength, as discussed by previous authors \\citep{bastien87,kenyon93p2,whitney97} due to the dominance of scattering in either the cavity or the disk. Stage~II sources show high polarization ($>3$\\,\\%) only when viewed edge-on, but show significant polarization ($>0.5$\\,\\%) over a range of viewing angles.   We have carried out an analysis of JHK, IRAC and MIPS spectral indices and color-color plots using our model SEDs, and have constructed `virtual clusters' in order to understand how a cluster of young stars might look in these color-color spaces. Our results indicate that: \\begin{itemize} \\item How well a given spectral index indicates the evolutionary stage of a source is dependent on the range of wavelengths of the fluxes used to calculate the spectral index (Figure~\\ref{f:sivsparams}). Our results suggest that the use of fluxes beyond $\\sim20$\\microns (e.g. \\mipsns), in addition to $1-10$\\microns fluxes, in calculations of the spectral index of YSOs is valuable for deriving information about the evolutionary stage of these sources. \\item For optically thin disks, the traditional $2.0\\rightarrow25.0$ spectral index increases with disk mass, while for optically thick disks, the same spectral index is insensitive to disk mass, as the emergent flux at these wavelengths originates only in the outer layers of the disk (Figure~\\ref{f:sivsparams}). \\item The near- and mid-IR colors of a disk-only source are sensitive to stellar temperature, disk inner radius, and disk flaring power (Figure~\\ref{f:siadditional}). In particular, we find that for temperatures above $5,000$\\,K, these colors become redder as the temperature increases, while the presence of an inner hole decreases the flux at near- and mid-IR wavelengths. For large enough holes, the near- and mid-IR colors are the same as those of a stellar photosphere. The effect of a larger disk flaring power is to increase the amount of flux intercepted by the disk, and thus to increase the amount of reprocessed flux at mid- and far-IR wavelengths. \\item For a given source, the spectral index is sensitive to the wavelength range of the broadband fluxes used to calculate the spectral index (Figure~\\ref{f:siranges}), and to a lesser extent is also dependent on whether the spectral index is calculated from the slope of a line joining two points, or as the slope of a least-squares fit line to a larger number of fluxes (Figure~\\ref{f:sidensity}). \\item (J-H) vs. (H-K) color-color plots can be used to identify sources as YSOs if they lie redward in (H-K) of the locus for reddened stellar photospheres, although their evolutionary stage cannot be reliably found simply from the JHK colors. In addition, many YSOs are likely to be indistinguishable from reddened stellar photospheres in regions of high extinction. \\item Stage~I models tend to occupy large regions in IRAC and IRAC+\\mips color-color plots, due to inclination effects, scattering in the bipolar cavities (which makes some edge-on sources blue), and variations in stellar temperature. \\item $[3.6]-[4.5]$ vs. $[5.8]-[8.0]$ IRAC color-color plots contain a large region that is likely to be occupied only by Stage~I sources. The region corresponding to \\citet{allen04}'s ``disk domain'' is likely to contain mostly Stage~II sources, but also a number of Stage~I sources. The remaining regions in this color-color space can be occupied by sources at {\\it any} evolutionary stage. \\item $[3.6]-[5.8]$ vs. $[8.0]-[24.0]$ IRAC+\\mips color-color plots provide a good separation of YSOs at different evolutionary stages. We note that even if a given source is not visible in \\mipsns, an upper limit on the flux can still constrain its evolutionary stage. \\end{itemize}  In Figure~\\ref{f:schematic} we show approximate regions in IRAC and IRAC+\\mips color-color space where models of a given evolutionary stage lie, irrespective of whether we are using the whole grid of models or the virtual clusters. In summary we find that near-IR fluxes (such as JHK fluxes) can be used to discriminate between reddened stars and many (but not necessarily all) YSOs, but cannot reliably determine the evolutionary stage of a YSO. Mid-IR fluxes (such as IRAC fluxes) should be efficient in separating YSOs from stellar photospheres, and allow some of the youngest sources to be unambiguously classified as such, although the evolutionary stage of the remaining YSOs is likely to be ambiguous. Finally, including fluxes at wavelengths beyond $\\sim20$\\microns (such as \\mipsns) in addition to near- and mid-IR fluxes substantially improves the ability to distinguish between the various evolutionary stages of YSOs."
    },
    "0608/astro-ph0608078_arXiv.txt": {
        "abstract": "{Very massive stars pass through the Wolf-Rayet (WR) stage before they finally explode. Details of their evolution have not yet been safely established, and their physics are not well understood. Their spectral analysis requires adequate model atmospheres, which have been developed step by step during the past decades and account in their recent version for line blanketing by the millions of lines from iron and iron-group elements. However, only very few WN stars have been re-analyzed by means of line-blanketed models yet.} {The quantitative spectral analysis of a large sample of Galactic WN stars with the most advanced generation of model atmospheres should provide an empirical basis for various studies about the origin, evolution, and physics of the Wolf-Rayet stars and their powerful winds.} {We analyze a large sample of Galactic WN stars by means of the Potsdam Wolf-Rayet (PoWR) model atmospheres, which account for iron line blanketing and clumping. The results are compared with a synthetic population, generated from the Geneva tracks for massive star evolution.} {We obtain a homogeneous set of stellar and atmospheric parameters for the Galactic WN stars, partly revising earlier results.} {Comparing the results of our spectral analyses of the Galactic WN stars with the predictions of the Geneva evolutionary calculations, we conclude that there is rough qualitative agreement. However, the quantitative discrepancies are still severe, and there is no preference for the tracks that account for the effects of rotation. It seems that the evolution of massive stars is still not satisfactorily understood.} ",
        "introduction": "Very massive stars pass through the Wolf-Rayet (WR) stage before they finally explode as supernovae or, possibly, $\\gamma$-ray bursts. The WR stars are important sources of ionizing photons, momentum, and chemical elements. However, the evolution of massive stars has not yet been safely established, and their physics is not well understood.  The empirical knowledge about WR stars is hampered by difficulties with their spectral analysis. Adequate model atmospheres, which account for the non-LTE physics and their supersonic expansion, have been developed step by step during the past decades. In a previous paper we presented a comprehensive analysis of the Galactic WN stars from their helium, hydrogen, and nitrogen spectra (Hamann \\& Koesterke \\cite{HK98}, hereafter quoted as Paper\\,I). The major deficiency of this generation of model atmospheres was its neglect of line blanketing by the millions of lines from iron and iron-group elements. Hillier \\& Miller (\\cite{HiMi98}) were the first to  include this important effect in their code {\\sc cmfgen}. The line-blanketed version of the Potsdam Wolf-Rayet (PoWR) model atmosphere code became available with Gr\\\"afener et al.\\ (\\cite{blanketing02}).  The improved models provide a much better fit to the observed spectra, and hence lead to more reliable determination of the stellar parameters. For many WN stars, a substantial revision of previous results can be expected. However, only very few WN stars have been re-analyzed so far by means of line-blanketed models (Herald et al.\\ \\cite{Herald01}; Marchenko et al.\\ \\cite{Marchenko+al2004}). In the present paper we now re-analyze all Galactic WN stars for which observed spectra are available to us.  The best way to analyze a larger sample of stars systematically is first to establish ``model grids'', i.e.\\ sets of model atmospheres and  synthetic spectra that cover the relevant range of parameters. For WN stars we have already prepared such grids of iron-line blanketed models (Hamann \\& Gr\\\"afener \\cite{blanket-WN}), which will be used in the present paper.  The main objective of analyzing the Galactic WN stars is to understand the evolution of massive stars. The empirical results from our previous comprehensive analyses were in conflict with the evolutionary calculations. However, {\\changed with the application of the more advanced models, we will revise the empirical stellar parameters. Moreover, the evolutionary models have also been significantly improved recently, since the effects of stellar rotation are now taken into account (Meynet \\& Maeder \\cite{MM03}). Hence, the question of whether stellar evolution theory still conflicts with the empirical stellar parameters should be reassessed. For this purpose, we use these new evolutionary tracks for generating synthetic stellar populations and compare them with the results from our analyses of the Galactic WN sample.}  The paper is organized as follows. In Sect.\\,2 we briefly characterize the applied model atmospheres. The program stars and observational data are introduced in Sect.\\,3. The spectral analyses and their results are given in Sect.\\,4, together with a long list of detailed comments on individual WN stars. Our empirical results are compared with evolutionary models in Sect.\\,5, especially by means of a population synthesis. ",
        "conclusions": ""
    },
    "0608/astro-ph0608552_arXiv.txt": {
        "abstract": "We study hard states of the black-hole binary XTE J1550--564 during its 2000 outburst. In order to explain those states at their highest luminosities, $L\\sim 10\\%$ of the Eddington luminosity, $L_{\\rm E}$, we propose a specific hot accretion flow model. We point out that the highest values of the hard-state $L$ are substantially above the $L$ an advection-dominated accretion flow (ADAF) can produce, $\\sim 0.4\\alpha^2 L_{\\rm E}$, which is only $\\sim (3$--$4)\\%L_{\\rm E}$ even for $\\alpha$ as high as $0.3$. On the other hand, we successfully explain the hard states with $L\\sim (4$--$10)\\%$ using the luminous hot accretion flow (LHAF) model. As $10\\%L_{\\rm E}$ is also roughly the highest luminosity an LHAF can produce, such an agreement between the predicted and observed highest luminosities provides by itself strong support for this model. Then, we study multi-waveband spectral variability during the 2000 outburst. In addition to the primary maxima in the optical light curves, secondary maxima were detected after the transition from the very high state to the hard state. We show that the secondary maxima are well modeled by synchrotron emission from a jet formed during the state transition. We argue that the absence of the corresponding secondary peak in the X-ray light curve indicates that the X-ray jet emission, regardless of its radiative process, synchrotron or its Comptonization, is not important in the hard state compared to the emission from the accretion flow. ",
        "introduction": "\\label{intro}  Black hole X-ray binaries (BHXBs) appear in five main spectral states, namely quiescent, low/hard, intermediate, high/soft, and very high ones (e.g., McClintock \\& Remillard 2006, though they prefer a different nomenclature). The origin of the X-ray spectrum of the hard state presents a highly interesting problem. The presence of a universal high-energy spectral cutoff above $\\sim$100--200 keV in the hard state points out to the X-ray emission coming from thermal Comptonization in a hot accretion corona with an electron temperature of $\\sim 10^9\\,{\\rm K}\\simeq 10^2\\,{\\rm keV}/k$ (see, e.g., Zdziarski 2000; Zdziarski \\& Gierli\\'{n}ski 2004 for reviews). We note that an alternative jet model (Markoff, Nowak \\& Wilms 2005), though also relying on thermal Comptonization as the dominant component at photon energies $\\ga 10$ keV, requires that the electron temperature is several tens higher, $kT\\sim 3$--4 MeV. Then, the origin of the $\\sim$100 keV cutoff is not clear; if it is due to the 1st order scattering by those thermal electrons, very strong fine tuning is required to achieve the observed cutoff energy. On the other hand, higher order scatterings will yield another spectral cutoff at $\\ga 3kT\\sim 10$ MeV, which presence remains to be observationally tested.  There are two main different models of the hot corona. One relies on electron heating by magnetic reconnection above the standard thin disk (Liang \\& Price 1977; Galeev, Rosner, \\& Vaiana 1979). Since the process of magnetic reconnection remains poorly understood, this model has not yet been elaborated in detail, although recent numerical simulations have shed some light on this problem (Hirose, Krolik \\& Stone 2006). The other model is a hot accretion flow, which is advection dominated (ADAF) in most of its parameter space (Narayan \\& Yi 1994, 1995; Abramowicz et al.\\ 1995; Rees et al.\\ 1982). We need to illustrate two points here when we use the term ADAF. First, it is found that in most of the parameter space of ADAF, outflow is moderately strong, as emphasized by Blandford \\& Begelman (1999, the ADIOS solution) and other authors. Second, as we will show later, another hot accretion solution may be responsible for the X-ray emission of the relatively luminous hard states. In comparison with the magnetic reconnection, ADAF has clear dynamics, and its astrophysical applications can be worked out in detail. The specific suggestion of applying the ADAF to the hard state is due to Narayan (1996) and Narayan, McClintock \\& Yi (1996) (see also Shapiro, Lightman \\& Eardley 1976, who, however, did not recognize the importance of advection). First detailed calculations applied to data were then done by Esin, McClintock \\& Narayan (1997, hereafter E97).  In the present paper we concentrate on the hot accretion flow. We note that the radiative efficiency of an ADAF increases with the accretion rate, $\\dot M$. Thus an ADAF can explain not only very dim sources like the Galactic center (e.g., Yuan, Quataert \\& Narayan 2003), but also some relatively luminous sources such as hard states of BHXBs. The highest rate, called the critical accretion rate, of an ADAF is determined by the balance between the viscous heating and cooling, $q_{\\rm vis}\\approx q_{\\rm ie}$, which occurs at  $\\dot{M}_{\\rm ADAF} \\sim 10\\alpha^2 \\dot{M}_{\\rm E}$, where $\\alpha$ is the viscosity parameter, $\\dot{M}_{\\rm E} \\equiv L_{\\rm E}/c^2$, $q$ is the rate of energy change per unit volume, and $L_{\\rm E}$ is the Eddington luminosity. The corresponding maximum luminosity is $L_{\\rm ADAF} \\sim 0.4\\alpha^2 L_{\\rm E}$ (E97)\\footnote{Thus the highest radiative efficiency of an ADAF is $\\simeq 0.04$, not much lower than that of a standard thin disk around a Schwarzschild black hole, $\\simeq 0.057$.}. The actual value of $\\alpha$ remains uncertain, but recent three-dimensional magnetohydrodynamic (MHD) numerical simulations suggest $\\alpha\\sim 10^{-2}$--$10^{-1}$ in most of the disk midplane (Hawley \\& Krolik 2001). Even for $\\alpha$ as high as 0.3 (a common value used in ADAF studies), the maximum ADAF luminosity is still only $\\simeq (3$--$4)\\%L_{\\rm E}$. It has been thought that above the $L_{\\rm ADAF}$ there is no hot accretion solution, with the thin disk remaining the only viable solution.  However, observations of the hard state in BHXBs often show higher luminosities during the rising part of an outburst. For example, the hard state of XTE J1550--564 reached $20\\%L_{\\rm E}$ (Sobczak et al.\\ 2000) and $\\sim 10\\%L_{\\rm E}$ (Rodriguez, Corbel \\& Tomsick 2003; this work) during its outbursts in 1998 and 2000, respectively. The highest hard-state luminosity of GX 339--4 was $\\simeq  25$--$30\\%L_{\\rm E}$ (Zdziarski et al. 2004; Done \\& Gierli\\'nski 2003). These observational results have long presented a puzzle for the hot accretion flow models.  Of importance in this context was the finding of a new hot accretion flow solution, the so-called luminous hot accretion flow (LHAF; Yuan 2001, hereafter Y01). Different from an ADAF, this solution corresponds to accretion rates above the critical rate of ADAF, $\\dot{M}_{\\rm ADAF}$, up to another critical accretion rate which is about 3-5 times higher than $\\dot{M}_{\\rm ADAF}$ for parameters adopted in Y01. In addition to accretion rates, the radiative efficiency of an LHAF is also higher than that of a typical ADAF. Therefore, an LHAF can produce much higher luminosity than an ADAF. Yuan \\& Zdziarski (2004) studied some luminous hard states of BHXBs and AGNs and found that their relatively low electron temperatures and high luminosities do require an LHAF to be present. In the first part of this work, we apply this approach to hard-state spectra of XTE J1550--564 during its 2000 outburst.  In addition to studying X-ray spectra alone, more and more efforts are devoted to multi-wavelength simultaneous observations and modeling of BHXBs, as illustrated, e.g., by studies of XTE J1118+480. Its nearly simultaneous observations were conducted from radio to X-ray wavelengths (e.g., Hynes et al.\\ 2000; McClintock et al.\\ 2001; Frontera et al.\\ 2001, 2003; Chaty et al.\\ 2003). These have allowed important theoretical conclusions to be drawn. Esin et al.\\ (2001) showed that the EUV data rule out models with a thin disk extending to the last stable orbit. Using a coupled jet-ADAF model, Yuan, Cui \\& Narayan (2005) and Malzac, Merloni \\& Fabian (2004) successfully explained the radio to X-ray spectrum and most of the timing properties of the source.  Here, we apply the same approach to XTE J1550--564. We use simultaneous multi-wavelength observations at radio, infrared, optical and X-rays during the 2000 outburst (Corbel et al.\\ 2001; Jain et al.\\ 2001, hereafter J01; Tomsick, Corbel \\& Kaaret 2001; Rodriguez et al.\\ 2003). Compared to XTE J1118+480, XTE J1550--564 is much more luminous. During the outburst, transitions between the hard and very high states occurred, unlike the case of XTE J1118+480, which was always in the hard state. Therefore some new interesting phenomena were observed. One issue we discuss below is that of secondary maxima in the infrared/optical light curves. J01 reported infrared ($H$), optical ($I$, $V$) and X-ray light curves for the 2000 outburst. After reaching the peak of the outburst, the IR, optical, and X-ray fluxes declined, and a transition from the very high state to the hard state occurred around MJD 51,680. Afterwards, the X-ray flux continued to decline whereas the IR/optical light curves began to rise and reached secondary maxima at MJD 51,699--51,705, after that they declined again. A contemporaneous spectrum in the radio band was obtained by Corbel et al.\\ (2001) at MJD 51,697. Here, we model the simultaneous radio, IR, optical, and X-ray data and explain the secondary IR/optical maxima.  In \\S \\ref{hot_flow}, we model the X-ray emission of the hard states, paying special attention to the dynamics of the hot accretion flow. In \\S \\ref{jet}, we study the simultaneous radio, IR/optical and X-ray data during the decline phase and explain the secondary IR/optical maxima. We conclude in \\S \\ref{sum} with a summary and discussion. ",
        "conclusions": "\\label{sum}  We have investigated the hard states of XTE J1550--564 during its 2000 outburst. We have modeled the X-ray spectra with accretion flow models (\\S \\ref{hot_flow}). Our model consists of an inner hot flow within a transition radius and an outer thin disk, as in E97. Different from E97, we have taken into account two new developments of the hot accretion flow models. One is the inclusion of outflow, the other is the extension of the ADAF model to higher accretion rates using the LHAF model (Y01). Our model fits well the lower three of the four characteristic hard-state spectra, see Figure \\ref{spectra}. The two extensions of the model are necessary for the obtained good fits. First, the inclusion of the outflow allows us to fit the X-ray spectral slopes. Without outflow, the emitted spectrum hardens with the increasing accretion rate (E97), which disagrees with the observations, where the slopes of the two lowest data sets are  almost identical in spite of their different luminosities. To resolve this discrepancy, the outflow becomes weaker with the increasing accretion rate. This is physically motivated as the radiative energy loss is smaller at lower accretion rates and thus the Bernoulli parameter is larger.  The extension of hot accretion flow models from ADAFs to LHAFs enables us to model the luminous hard states with luminosities up to $\\sim 8\\%L_{\\rm E}$ (while the ADAF is limited to $\\la 3\\%L_{\\rm E}$). The LHAF fits the spectral slope and the cut-off energy of those states, as shown in Figure \\ref{spectra}, confirming the prediction of Yuan \\& Zdziarski (2004).  The value of the highest possible luminosity of the LHAF model, $L_{\\rm LHAF}$, depends on $\\alpha$ and other model parameters, and the outer boundary condition (see Y01). For our assumed parameters, we find $L_{\\rm LHAF}\\simeq 8\\%L_{\\rm E}$, somewhat lower than $L$ of the most luminous hard state considered here shown by the blue symbols in Figure \\ref{spectra}. However, we can see that this state appears intermediate between the hard and the very high states, and perhaps some prcesses responsible for the very high state are already for importance here. {\\em Thus, the overall agreement between the maximum observed $L$ and the prediction of the LHAF model provides strong support for it as a model of the X-ray emission of the hard states of BHXBs}.  The highest $L$ of the hard state varies among various sources as well as different outbursts of a given source. As discussed in \\S \\ref{intro}, the maximum hard-state $L$ of XTE J1550--564 and GX 339--4 varies in the $\\sim$10--30\\% range. On the other hand, the corresponding value in Cyg X-1 is only $\\sim 0.01$--$0.02 L_{\\rm E}(d/2\\,{\\rm kpc})^2(10\\msun/M)$ (e.g., Zdziarski et al.\\ 2002). We discuss now possible reasons for this variability.  First, $L_{\\rm LHAF}$ is a function of the model parameters, e.g., $\\alpha$, and parameter variations among different sources will contribute to its dispersion. However, numerical simulations of accretion flow indicate that $\\alpha<0.3$ (e.g., Hawley \\& Krolik 2001). Thus, this effect can explain a low value of the maximum $L$, but not values $\\ga 10\\%L_{\\rm E}$. Still, we can obtain $\\sim 15\\%L_{\\rm E}$ by adjusting the outer boundary condition (determined by the physics of the transition region between the hot and cool accretion flows, which is complex and currently unclear).  An efficient way of increasing the maximum hard-state $L$ can be due to the angular momentum of the black hole being $>0$. In the present calculations, we have assumed a Schwarzschild black hole. If the black hole is rapidly rotating, substantially more gravitational energy will be released because of the deeper potential well (e.g., Gammie \\& Popham 1998; Popham \\& Gammie 1998). Shafee et al.\\ (2006) indeed have found that a number of BHXBs appear to have the spin parameter of $\\sim$0.6--0.8.  Yet another effect that could increase $L_{\\rm LHAF}$ is a magnetic coupling between the rotating black hole and the accretion flow. If the black hole rotates faster than the disk, energy and angular momentum will be extracted from the black hole and transferred to the disk by the magnetic field (Li 2000, 2002; Wang et al.\\ 2003). However, we are not able to estimate quantitatively the importance of this effect here because the existing studies have only considered the standard thin disk rather than a hot flow\\footnote{Recently attention has been paid to the effect of the presence of magnetic stresses near and inside the last stable circular orbit on the radiative efficiency of the standard thin disk (e.g., Agol \\& Krolik 2000; Beckwith, Hawley \\& Krolik 2006). This effect does not apply to the hot accretion flow since the no-torque boundary condition is not adopted in the global solution of hot accretion flow.}.  As we point out above, the most luminous hard states border the very high state. The nature of the latter remains an open question. Numerical calculations by Y01 (see also Begelman, Sikora \\& Rees 1987) indicate that above $\\dot{M}_{\\rm LHAF}$ the hot accretion flow will collapse within a certain radius and form a cold annulus (see Liu, Meyer \\& Meyer-Hofmeister 2006 for a similar result in the context of their evaporation model). The inner cold annulus will emit black body radiation which may explain the thermal component of the very high state. At the collapse region, strong magnetic reconnection events should occur due to the overflow of the magnetic field (Begelman et al.\\ 1987; Y01). This is likely to accelerate some electrons into a power-law distribution. Compton scattering of the soft photons from the inner annulus by both the thermal and power-law electrons may then be responsible for the X-ray emission of the very high state. We note that Gierli{\\'n}ski \\& Done (2003) have found that the shape of the soft $\\gamma$-ray spectrum of XTE J1550--564 in the very high state indeed requires the presence of both thermal and nonthermal (power-law like) electrons.  A possible problem with the above scenario is that the hot accretion flow is strongly thermally unstable at the collapse region (Yuan 2003). As a result of the instability, the accretion flow may form a two-phase flow structure, with cold clumps surrounded by the hot phase (Yuan 2003; Guilbert \\& Rees 1988; Celotti, Fabian \\& Rees 1992; Kuncic, Celloti \\& Rees 1997; Krolik 1998; Beloborodov 1999; Merloni et al.\\ 2006). The radiation from the cold clumps can then explain the thermal component of the very high state, whereas Comptonization of the soft photons by the high-energy electrons in the hot phase may be responsible for the power law tail. Thus, this represents an alternative scenario for the very high state. (However, a hybrid electron distribution in the hot phase is required as stated above, see Gierli{\\'n}ski \\& Done 2003.) We speculate that even some luminous hard-state flows with $\\dot{M}\\la \\dot{M}_{\\rm LHAF}$ may also contain such a two-phase configuration. The formation of the cold clumps in this case could be due to the strong instability in the transition region between the hot accretion flow and the standard thin disk (Kato 1999, 2000). This scenario is promising in explaining the observed power spectra of the rapid aperiodic variability of X-ray emission of BHXBs (B\\\"ottcher \\& Liang 1999). It also may explain the claimed relativistic iron K emission line profiles in BHXBs (e.g., Miller et al.\\ 2006). We plan to investigate these issues in future work.  We have also investigated the nature of the secondary peaks in the $HIV$ light curves of XTE J1550--564 (J01). They began to emerge after the transition from the very high state back to the hard state, and reached the maximum $\\sim$20 d later. However, there have been no corresponding peak in the X-ray light curve. Based on a detailed modeling to the multi-waveband spectrum from the radio to X-rays, we conclude that the secondary peaks are due to the emission from the jet formed during the state transition. The mass loss rate in the jet kept increasing for $\\sim$20 d after its birth, and then decreased. The secondary peaks were then due to the corresponding maximum of the relative contribution of the jet ($HIV$) emission with respect to the accretion emission. The absence of the corresponding peak in the X-rays strongly indicates that the jet emission (either synchrotron or Compton) is negligible in the X-ray band in the hard state. This differs from the scenarios with the dominance of the jet in X-rays (e.g., Markoff et al.\\ 2005). In principle, it is possible that a drop in the X-ray flux from the disk is exactly compensated by a rise in the X-ray flux from the jet, with the total flux changing smoothly. As this scenario requires fine tuning, we consider it to be highly unlikely.  We also point out that Xue et al.\\ (2007) present another argument against the jet dominance in X-rays in the hard state. By modeling the spectrum of the jet blob, they obtain the parameters of the electron energy distribution. Assuming the electron energy distribution is the same for all blobs in the jet, they model the overall radio spectrum of the source and extrapolate their fit to higher energies. They find that the synchrotron emission from the jet contribute negligibly to the observed X-ray emission."
    },
    "0608/astro-ph0608287_arXiv.txt": {
        "abstract": "We consider the question of whether it is worth building an experiment with the sole purpose of bringing the detectable limit on the tensor-to-scalar ratio $r$ down to $10^{-3}$. We look at the inflationary models which give a prediction in this region and recap the current situation with the tensor mode, showing that there are only three known models of inflation which give definitive predictions in the region $10^{-3}<r<10^{-2}$. ",
        "introduction": "Inflation predicts a unique spectrum of gravitational waves with a tensor spectral index $n_T\\lesssim{}0$ a detection of which would provide the ``smoking'' gun that would affirm inflation \\cite{smoke}. Currently, model fitting of the CMB data is used to obtain the allowed range on the tensor-to-scalar ratio $r$ , the updated fit by Ref.~\\cite{wmap3} places the upper limit on $r$ from WMAP as $r<0.90$ with no running in the spectral index. % The PLANCK satellite promises the ability to detect $r\\gtrsim0.005$ \\cite{planck}, and the ground based CLOVER claims that it will bring this limit down to $r\\gtrsim10^{-2}$ \\cite{clover}.   There are in fact two methods to measure the primordial gravitational wave background (pGWB), either indirectly for large scales via the B-mode of the CMB-polarisation such as ESA's B-mode satellite \\cite{bsat} or NASA's inflation probe \\cite{probe}, or directly for small scales using space based laser interferometers such as NASA's Big Bang Observer (BBO) or Japan's Deci-Hertz Interferometer Gravitational Wave Observatory (DECIGO). The contribution of these experiments to deciphering inflation's signatures have been considered \\cite{Smith:2006xf,Smith:2005mm,EC,Boyle:2005se}. Refs.\\cite{Smith:2005mm,Smith:2006xf} use both model dependent and independent methods to evaluate the pGWB amplitude predicted by inflation, they conclude that direct detection methods would not further constrain models, but that if $r\\lesssim{}0.01$ would be useful in constraining the scale of inflation. Ref.~\\cite{Boyle:2005se} concur with this and further note the degeneracy in measuring the tensor power spectrum at small (BBO) scales between models which predict notably different tensor power spectra at large (CMB) scales. However they do point out that depending on results of the BBO, there is potential to gain deeper insight into the physics of reheating. We note that this does have (indirect) implications in constraining inflationary models, since various models of inflation predict various lower limits on the number of $e-$folds of inflation when combined with observational limits on the spectral index \\cite{AL1,AL2}.  Judging by these papers, any advances in model discrimination from the pGWB signal will have to come from measuring $r$ at CMB scales, and the onus is on CMB-polarisation experiments. However, the future experiments focusing on this are the B-mode satellite and the inflation probe, with the ultimate lower limit on $r$ being $r\\sim10^{-3}$, yet we can measure $r\\gtrsim10^{-2}$ with PLANCK and CLOVER. So the question is whether it is worth building new experiments solely to measure $10^{-2}<r\\lesssim10^{-3}$.  \\par{}This question has been raised for measuring $10^{-2}<r<10^{-4}$ in Ref.~\\cite{EC}. They noted that this level of improvement in the detectability of $r$ corresponded to an improvement on the limit on the energy scale on the order of $\\log(\\Delta{}V)\\sim4$. They argue that there is no reason to expect the energy range to be so narrow, and conclude by advising that the next generation of CMB probes focus on other signatures of inflation, such as the non-gaussian signal.  \\par{}In this paper we reconsider the purpose of such an improvement in $r$ detectability by looking at the specific predictions for $r$ from models of inflation. We revisit the $\\log(r) \\rm{vs.} n$ plot in Ref.~\\cite{AL1}, and add the logarithmic bounds on $r$ from a WMAP year 3 % data fit. We also discuss the assumptions made to obtain the $r$ predictions and make amends to the prediction for mutated hybrid inflation so that it conforms to particle-physics requirements.  \\par{}This paper is laid out as follows, in section \\textbf{II} we recap the definition of $r$, then in section \\textbf{III} we review the inflation models giving an $r$ prediction, in section \\textbf{III A} we re-calculate the upper bound on the $r$ prediction for the mutated hybrid inflation model, and in section \\textbf{III B} we briefly introduce brane induced inflation. In section \\textbf{IV} we discuss the possibility of detecting the gravitational wave background (GWB) signature in the region $-3<\\log(r)<-2$. ",
        "conclusions": "There are three currently known models which give definitive predictions in the desert, five if we take into account F/D-term inflation and a near upper-limit of brane inflation; the other models either give predictions in the $\\log(r)\\ll-3$ range or will have been ruled in or out by experiments measuring up to $\\log(r)=-2$. We thus concur with Ref.~\\cite{EC} that a singular mission to detect the gravitational wave background on angular scales of $\\sim2^o$ \\footnote{$2^o$ is the angular scale on which the signature of primordial gravitational waves can be found} is impractical, as we do not expect a detection in this range.  \\par{}The beyond Einstein inflation probe and B-mode satellite, overcome the unpromising nature of detecting the tensor-to-scalar ratio by including a wider range of multipoles to get a fuller spectrum of the B-mode polarisation signature. This means that both the scales on which the primordial gravitational wave background and those on which the re-ionization signatures are included, hence enabling a wider range of research interests to benefit (see for example Ref.~\\cite{KM}), while still probing the tensor desert.   \\par{}On a final note, this particular signature of inflation is crucial in extending the work on discriminating between models of inflation \\cite{AL1,AL2}. If a detection at $\\log(r)>-3$ \\emph{is} made, all models with $\\Delta\\phi<M\\sub{pl}$ would be eliminated, and even if a detection is not made, it would automatically rule out models of inflation with $\\Delta\\phi\\geqslant{}M\\sub{pl}$. The latter case would also set a a more definitive upper limit on $\\mu$ in brane inflation, as well a slightly stronger constraint on the $g$ term in the logarithmic potential."
    },
    "0608/astro-ph0608414_arXiv.txt": {
        "abstract": "I touch upon some of the discoveries made by the Swift Team during the first 18 months of operation focusing on a few critical points. In addition to the early afterglows and complete coverage of the light curves, we mention the discovery of the location of Short bursts, the statistics on flares that is under completion and refer to their explanation as due to late internal shocks. The full understanding of the connection with supernovae and the study of high z flares will ultimately lead to the detailed understanding of the mechanism of explosion both for GRBs and for Supernovae. ",
        "introduction": "Following the launch of the Swift satellite, Gehrels et al. (2005) the Swift Team, and the community at large, witnessed an unprecedented collection of data. The instrument performed excellently and more than anything else many of us felt that the key advantage consisted in the rapidity by which the spacecraft was able to point the target on the axis of the Narrow Field Instruments (NFI) and in a perfect worldwide organization that allowed the real time data analysis and follow up both from the ground and from space. In addition the performance of the spacecraft, the ASI (Agenzia Spaziale Italiana) ground station was far beyond specifications thanks to the Malindi staff so that all the data were, and are being, received quickly and without losses. Now, at about 18 months from launch, we have a reasonable understanding of the data and we can tackle the statistical studies since the sample is large enough (we detect about 102 bursts per year) to allow it. It is also time to improve theory and interpretation. The data are now good enough that in spite of the fact that the fireball model (Rees \\& Meszaros, 1992) has been confirmed in its more general features, many details are lacking and such details, once they are related to the observations, are fundamental to understand the physics at work. The work carried out with Swift has been, so far, a Team work where everybody participated to the research we carried out in different roles and capabilities. The input of the theoretician has been extremely useful since thanks to them the team was greatly facilitated in the understanding of the models and of the observations and move swiftly since the very beginning.  It is an impossible task to evidence in a few pages the richness of the data we obtained in this period, data that relate to a large number of bursts and involve other satellites and the largest ground based facilities of the world. In addition very many review articles and excellent journal articles are being published in the field. We will simply try to touch upon the most important issues and briefly illustrate some of the data we obtained.  In the following we will discuss some of the highlights of the Swift observations and discoveries.  In the second section we describe the theoretical framework.  In the third section we will discuss the morphology of the light curves.  To give a broader understanding of the capability of the Swift mission and to illustrate the power of combining observations by different satellites, in Section 4 we briefly discuss the burst GRB060124.  In the Section 5 we will deal with flares. One of the major discoveries carried out by Swift has been the location of the short GRBs. This will be discussed in Section 6. Finally we will touch upon two fundamental issues again related to discoveries made by Swift, the connection between GRBs and SNe, Section 7, and the detection of high z objects, Section 8. We conclude with the host galaxies and progenitors with the conclusion that the real research work and deep understanding of what is going on just started.  \\section {Generalities on the model}   The framework of the discussion in this paper is that of the collapsar (Woosley 1993, MacFadyen \\& Woosley 1999, Zhang \\& Woosley, 2003) while for the emission we refer to the fireball model by Rees \\& Meszaros (1992) and following modifications. In the collapse of the iron core of a massive rotating star we likely end up with a black hole and an accreting disk. In the process a large amount of energy is produced that finds its way out the star as a well collimated jet, $\\sim$10 degrees, and later impact in the interstellar medium. The decrease of the external pressure after the break out causes the front of the jet to expand. The central engine, black hole and accretion disk, has a scale length of about 10$^6$cm and the energy released is of the order of 10$^{50}$ -- 10$^{52}$ erg. The emission is turbulent and relativistic pellets move at different Lorentz factors causing an internal shock when a faster pellet (shell) catches up with a slower pellet. The physics of the encounter between two pellets, conversion of kinetic energy in internal energy and emission of electromagnetic waves, is rather straightforward (see for instance Kobayashi et al., 1997). The scenario becomes now somewhat more complex. In the immediate vicinities of the progenitor we certainly have a small ionized region of the interstellar medium or a cocoon formed by highly absorbing material. Wind, as observed in massive stars and Wolf Rayet stars, may be present. It is important to disentangle the various components of the environment in order to better understand the physics, the characteristics of the environment and of the progenitor. A high Lorentz factor ($\\gamma$) pellet may be expelled first and this could interact with the ISM before the internal shock occurs. While detailed work and simulations are going on in order to fully clarify the evolution of the ejected material, it suffices to remind now that in addition to the clashing pellets (shells) causing the so called internal shock we have an external shock as soon as the material encounters the interstellar medium. Shocks, as it is well known, are characterized by a forward and reverse shock with Synchrotron and Inverse Compton emission. It is striking that such a complex set of phenomena will originate, as we will see in the following, a quite simple observed output with rather standard characteristics.  As we know, and we will describe later, Swift located also the short bursts. The short duration of the prompt emission rules out immediately the collapse of a massive star for which the free fall time is too long ,t$_{ff} \\propto M^{-1/2} R^{-3/2}$. These seem to be due to the merging of a NS-NS or NS-BH binary and the events show characteristics in the light curve that are very similar to those observed in the long bursts. The basic physical phenomena must be simple and similar. ",
        "conclusions": "There is no doubt that with the observations we obtained with the Swift satellite we made, and are making, a fundamental step forward in understanding observationally the evolution of the event and in collecting data for statistical studies. At the same time it has been possible to test theories and at least have an indication of what is acceptable and what needs to be revised. Indications that came earlier from observations from GRBs detected by Beppo SAX or other satellites, were confirmed, as the GRBs-SNe connection, or denied, as the simple power law decaying light curve, and the theory has now a more solid ground on which to proceed. So far the interpretation of the data has been done, by a large extent, on the existing theory that, in their basic features, hold.  We showed that the field touches almost in all branches of astronomy and is very demanding for the state of the art instrumentation. Knowledge of the performance of the central engine and of the progenitors is intimately related to the complex physics of the death of massive stars and to all those phenomena accompanying the evolution of relativistic stars. The evolution, and physics of the collapse, is intimately connected to the GRBs - SNe connection and to the abundance of the progenitors and of the host galaxies. This open the field to evolution in relation to cosmology and at the end the high luminosity and the possibility to be detected at very large reshifts make them a fantastic beacon to probe the high z Universe and the interstellar medium.  The most important observations seem to be those capable of detecting high z GRBs that could be immediately followed up by ground based observations with large and robotic telescopes. Robotic telescopes are essential to ease such search. We may need to think about desgning a few two meters aperture robotic telescopes. While all the data tend to suggest we might have more long bursts at high redshifts either because of the characteristics of the progenitors and host galaxies or because of statistics based on what we observed so far or simply because the phenomenon seems to prefer low abundances galaxies it remain a puzzle why so far we detected only one GRB at 6.29 (the detection of a possible object at higher redshift, GRB060116 remains uncertain). How many did we loose because of limitations related to the zone of the sky in which the bursts have been detected and the difficulty of a proper and immediate follows up?"
    },
    "0608/astro-ph0608308_arXiv.txt": {
        "abstract": "Analyses of spectra obtained with the {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)} satellite, together with spectra from the {\\it Copernicus} and {\\it IMAPS} instruments, reveal an unexplained very wide range in the observed deuterium/hydrogen (D/H) ratios for interstellar gas in the Galactic disk beyond the Local Bubble. We argue that spatial variations in the depletion of deuterium onto dust grains can explain these local variations in the observed gas-phase D/H ratios. We present a variable deuterium depletion model that naturally explains the constant measured values of D/H inside the Local Bubble, the wide range of gas-phase D/H ratios observed in the intermediate regime ($\\log N({\\rm H~I}) =$ 19.2--20.7), and the low gas-phase D/H ratios observed at larger hydrogen column densities. We consider empirical tests of the deuterium depletion hypothesis: (i) correlations of gas-phase D/H ratios with depletions of the refractory metals iron and silicon, and (ii) correlation with the H$_2$ rotational temperature. Both of these tests are consistent with deuterium depletion from the gas phase in cold, not recently shocked, regions of the ISM, and high gas-phase D/H ratios in gas that has been shocked or otherwise heated recently. We argue that the most representative value for the total (gas plus dust) D/H ratio within 1 kpc of the Sun is $\\geq 23.1\\pm 2.4$ (1~$\\sigma$) parts per million (ppm). This ratio constrains Galactic chemical evolution models to have a very small deuterium astration factor, the ratio of primordial to total (D/H) ratio in the local region of the Galactic disk, which we estimate to be $f_d \\leq 1.19^{+0.16}_{-0.15}$ ($1\\sigma$) or $\\leq 1.12\\pm 0.14$ ($1\\sigma$) depending on the adopted light element nuclear reaction rates. ",
        "introduction": "Accurate measurements of the deuterium/hydrogen (D/H) ratio (by number) provide critically important tests for models of primordial nucleosynthesis, Galactic chemical evolution, and the chemical properties of the intergalactic medium (IGM). Among all of the light elements, deuterium arguably provides the most stringent tests of our understanding of these three critically important topics, because the evolution of deuterium is predicted to be rather simple and the dependence of D/H varies steeply with baryonic density. \\citet{eps76} and others have argued that deuterium was formed only in the Big Bang and that D is easily converted to $^3$He, $^4$He, and heavier elements by nuclear reactions in stars, a process commonly called astration. Thus, as stars evolve, the interstellar medium (ISM) receives deuterium-depleted and metal-enhanced material from stellar winds and supernova explosions, leading to a monotonic decrease in the D/H ratio with increased processing through stars. Other light elements (e.g., $^3$He, $^4$He, $^7$Li) have more complex creation and destruction histories or otherwise provide less sensitive tests of these theories \\citep[cf.][]{boe85,wil94}. Infall of less astrated gas, which is deuterium rich and metal poor, from the IGM and small galaxies captured by the Milky Way Galaxy, complicates models of Galactic chemical evolution, because the rate of infall, its time and spatial dependence, and the timescale for mixing with other gas are poorly known.  Although based on different physical processes and relating to very different times in the early universe, the two main approaches for determining the primordial D/H ratio, (D/H)$_{\\rm prim}$, are now in agreement. The first approach is to measure the column densities of D~I and H~I in quasar absorption line systems. \\citet{kir03} have determined a mean value of  D/H = $27.8^{+4.4}_{-3.8}$ parts per million (ppm) by averaging the log (D/H) values using column densities $N$(H~I) and $N$(D~I) obtained from absorption line measurements of gas in the lines of sight toward five quasars. Most of these measurements use high resolution spectra obtained with the {\\it Keck} telescope. Since the observed gas has very low metallicity, these values of D/H are likely close to, but perhaps slightly smaller than, the primordial (Big Bang) value, (D/H)$_{\\rm prim}$. The measurement of D/H = $22\\pm 7$ ppm for the high velocity cloud Complex C \\citep{sem04} likely refers to gas located 6--10 kpc into the Galactic halo with low metallicity, 10--30\\% solar abundance \\citep{col03}, and this value should be somewhat smaller than (D/H)$_{\\rm prim}$.  The second approach involves the analysis of data from the {\\it Wilkinson Microwave Anisotropy Probe (WMAP)} satellite and other cosmic microwave background experiments. Assuming a flat $\\Lambda$-dominated universe and other standard assumptions, several authors \\citep{spe03,cyb03,coc04,san06} have derived tight constraints on the baryon closure parameter, $\\Omega_{\\rm b}$, and the Hubble constant. Together with other constraints from the power spectrum of galaxy clustering and L$\\alpha$ forest data, they find that $\\Omega_b h^2 = 0.0224$ with a 1$\\sigma$ error of about 4\\%. These results are summarized in Table~1. The ratio of baryons to photons (multiplied by $10^{10}$), $\\eta_{10}$, is simply related to $\\Omega_b h^2$ \\citep{bur01}. In the standard big bang nucleosynthesis (SBBN) model, $\\eta_{10}$ is the single parameter that predicts the primordial light element abundances including (D/H)$_{prim}$. For example, \\citet{spe03}, \\citet{coc04}, and \\citet{san06} obtain (D/H)$_{\\rm prim} \\approx 26$~ppm. Using the same values of $\\Omega_b h^2$ and $\\eta_{10}$ but a more recent compilation of nuclear reaction rates \\citep{ang99}, \\citet{cyb03} obtain (D/H)$_{\\rm prim} = 27.5^{+2.4}_{-1.9}$ ppm. The somewhat different values of (D/H)$_{\\rm prim}$ arise from the different authors treating the light element nuclear reaction rates and error propagation somewhat differently (cf. Nollett \\& Burles 2000). We therefore use in the subsequent analysis two different values for (D/H)$_{\\rm prim}$, $27.5^{+2.4}_{-1.9}$ ppm \\citep{cyb03} and $26.0^{+1.9}_{-1.7}$~ppm \\citep{coc04}. Both values are consistent with the quasar absorption line result for (D/H)$_{\\rm prim}$.  Beginning with observations by the {\\it Copernicus} satellite, data from high-resolution ultraviolet spectrometers on several satellites have been analyzed to measure D/H in the Galactic ISM. These results provide insight into the abundance of deuterium in the present epoch, which with measurements of various metal abundances constrain models of Galactic chemical evolution. The {\\it International Ultraviolet Explorer (IUE)} satellite and more recently the {\\it Goddard High Resolution Spectrometer (GHRS)} and {\\it Space Telescope Imaging Spectrograph (STIS)} instruments on the {\\it Hubble Space Telescope (HST)} have provided UV spectra of the Ly$\\alpha$ absorption lines of H~I (1215.67 \\AA) and D~I (shifted by --82 km s$^{-1}$ relative to the H~I line), which are needed to measure D/H in the ISM along lines of sight to stars located within a few hundred parsecs of the Sun. The same analysis technique has been applied to hydrogen and deuterium lines higher in the Lyman series that lie in the spectral range 912--1025 \\AA, using data obtained with the {\\it Copernicus}, the {\\it Interstellar Medium Absorption Profile Spectrograph (IMAPS)}, and the {\\it FUSE} satellite. For a description of the {\\it FUSE} satellite and its operations, see \\citet{moo00} and \\citet{sah00}.  The far-UV spectral range is essential for sampling more distant lines of sight using hot subdwarfs and white dwarfs as background sources, because the decreasing opacity in the higher Lyman lines allows one to observe the deuterium lines without masking by the nearby very opaque broad hydrogen lines. For example, the dark core of the H~I Ly$\\alpha$ line absorption broadens to mask the D~I Ly$\\alpha$ line when the hydrogen column density $\\log N({\\rm H~I}) > 18.7$, but at Ly$\\gamma$ (972~\\AA) the horizon for detecting the DI Ly$\\gamma$ line extends to $\\log N({\\rm H~I}) \\approx 20.0$. Measurements of D/H using the higher Lyman series lines are feasible to about $\\log N({\\rm H~I}) \\approx 21$ (cf. Hoopes et al. 2003), where the density of overlapping H$_2$ lines and the large widths of the interstellar H~I Lyman lines determine the final horizon for such analyses.  There are several subtle effects that can enter the analysis of D~I and H~I lines in the far UV. One disadvantage is that, for long lines of sight, there is an increased chance of weak features at large velocity displacements. An H~I feature near --82 km s$^{-1}$ can masquerade as D~I absorption. An important advantage is that the far UV includes a large number of Lyman lines with a wide range of optical depths. Thus, measurements of $N$(H~I) and $N$(D~I) can be based on several spectral lines in the Lyman series. Another advantage is that, for the higher Lyman lines, the upper state lifetimes are long, decreasing the damping and better separating the H~I and D~I absorption for each line in the Lyman series and between adjacent lines in the series.  Since the very first of these measurements, a disturbing aspect of the D/H studies has been the wide range of measured values with no apparent pattern in the data or a convincing explanation for it. We address the first part of this problem with the latest data, including many measurements with {\\it FUSE}; we conclude (see \\S~2) that the variations are real, and we show that there is a consistent pattern. We propose that the variations arise from time-dependent deuterium depletion onto dust grains (see \\S~4 and \\S~5), and we examine two critical tests (see \\S~6) of this model. The first test (\\S6.1) uses iron and silicon depletions as proxies for grain destruction, and the second test (\\S6.2) uses correlations of D/H with H$_2$ excitation temperatures. We conclude (see \\S~7 and \\S~9) that the most likely value for the total D/H ratio within 1 kpc of the Sun is about 45\\% larger than the previously assumed value of $\\approx$15 ppm, which provides a revised constraint on models of Galactic chemical evolution. Also, the inferred depletion of deuterium from the gas phase places constraints on the physics and chemistry of interstellar dust. See \\citet{lin06} for an early version of this paper. ",
        "conclusions": "The main goal of the {\\it FUSE} mission has been to obtain accurate measurements of (D/H)$_{\\rm gas}$ for many sightlines in the Milky Way Galaxy and beyond in order to measure (D/H)$_{\\rm prim}$ and to obtain constraints on Galactic chemical evolution. {\\it FUSE} has now obtained the data needed for accurate measurements of (D/H)$_{\\rm gas}$ much further out in the Galaxy and with higher precision than previous instruments. In our summary of the published measurements of (D/H)$_{\\rm gas}$ obtained from {\\it FUSE}, {\\it HST}, {\\it Copernicus}, and {\\it IMAPS}, we find, following \\citet{woo04}, that the (D/H)$_{\\rm gas}$ measurements appear to fall into three distinct groups depending on the neutral hydrogen column density $N$(H~I) to the target star. For lines of sight with $\\log N({\\rm H~I}) < 19.2$, that is within the Local Bubble, the mean value of (D/H)$_{\\rm gas}$ is $15.6\\pm 0.4$ ppm, where the uncertainty is the standard deviation of the mean. At large column densities, $\\log N({\\rm H~I}) > 20.7$, we find that the five lines of sight have much lower values of (D/H)$_{\\rm gas}$ with a mean value of $8.6\\pm 0.8$~ppm. In the intermediate regime there is a factor of 4--5 range in the high precision (D/H)$_{gas}$ measurements with values as low as $5.0^{+2.9}_{-3.4}$~ppm ($\\theta$~Car) and as high as $21.8\\pm 2.1$~ppm ($\\gamma^2$~Vel). Since this large range has been measured from spectra obtained by several instruments, it is not an instrumental artifact, but rather an observational result that requires an explantion.  Both large changes in the local astration rates and infall of D-rich and metal-poor gas from the halo or IGM cannot explain the large range in (D/H)$_{\\rm gas}$ without predicting large variations in the local ISM metal abundances, in particular (O/H)$_{\\rm gas}$, which are not observed. Instead, we describe a time-dependent deuterium-depletion model in which much of the deuterium in the ISM can reside on grains under the appropriate conditions. The theoretical basis for this model is the lower zero-point energy of C--D bonds compared to\\\\ C--H bonds. When the ISM is undisturbed for a long time, grains are very cold, and deuterium can replace hydrogen on carbonaceous grains and perhaps other types of grains. We estimate that D/H in the grains can approach $\\sim$0.2, thereby reducing (D/H)$_{\\rm gas}$ to 6~ppm or lower. When a region of the ISM is shocked or the gas comes close to a hot star, the grains are partially or fully destroyed, sending deuterium atoms back into the gas phase. In this model, the value of (D/H)$_{\\rm gas}$ for a given line of sight depends on the environment and past history of the grains.  Evidence in support of this model includes the measurement of high D/H ratios in interplanetary carbonaceous dust grains, which are likely interstellar in origin. Since the conditions required from deuterium depletion are similar to those for depletion of metals like Fe, Si, and Ti in the ISM, the model predicts a correlation of depletions of these metals with low values of (D/H)$_{\\rm gas}$. We present data that strongly support the correlation for Fe and Si, and \\citet{pro05} present the data in support of correlation with Ti. The model also predicts that (D/H)$_{\\rm gas}$ should be correlated with the recent thermal and ionization history of the gas, whereas the variable astration hypothesis would not be expected to show such a connection; we find that (D/H)$_{\\rm gas}$ has a significant positive correlation with the rotational temperature of the molecular hydrogen on the line-of-sight, further supporting the hypothesis that (D/H)$_{\\rm gas}$ is affected by interstellar processes acting on relatively short time scales. We note that in the most recent discussion of Galactic chemical evolution models, \\citet{rom06} conclude that ``depletion of deuterium on to dust grains is the most likely physical mechanism proposed so far to explain the observed dispersion in the local data.''  If this model is realistic, then the best estimate of the total deuterium abundance, (D/H)$_{\\rm total}$ would be the highest measured values of (D/H)$_{\\rm gas}$ because these would be for lines of sight with minimum depletion of deuterium. We note that there are five lines of sight in the intermediate $N$(H~I) regime with high values of (D/H)$_{\\rm gas}$. The weighted mean value for these four lines of sight is $21.7\\pm 1.7$~ppm (standard deviation of the mean). Since there could be some deuterium depletion even in these lines of sight, we conclude that the best estimate of the total D/H in the local disk region of the Galaxy, after subtracting the Local Bubble foreground column densities, is (D/H)$_{\\rm LDtot} \\geq 23.1\\pm 2.4$~ppm\\footnote{Very recently \\citet{rog05} detected the hyperfine ground-state transition of atomic deuterium at 327~MHz (92~cm) in the Galactic anticenter direction. They find that (D/H)$_{\\rm gas} = 23\\pm 4$~ppm ($1\\sigma$) or $23^{+15}_{-13}$~ppm (3$\\sigma$), and estimate that the column density-weighted mean distance of the emitting gas is about 2~kpc. Since the beamwidth is $14^{\\circ}$, the beam averages an unknown amount of cloud and diffuse gas in and above the Galactic plane. While it is difficult to compare this new result with our line of sight measurements over shorter paths, the value of (D/H)$_{\\rm gas}$ obtained from the hyperfine transitions of D and H is consistent with our proposed value of (D/H)$_{\\rm LDtot}$.}. This new estimate is $\\geq 58$\\% higher than the value of 15~ppm recently used in some Galactic chemical evolution models (e.g., Romano et al. 2003), which is (D/H)$_{\\rm gas}$ for the Local Bubble rather than (D/H)$_{\\rm LDtot}$.  Two methods for determining (D/H)$_{\\rm prim}$ based on data from very different times in the early universe are now in agreement. Measurements of (D/H)$_{\\rm gas}$ in five quasar absorption line systems have a mean value of $27.8^{+4.4}_{-3.8}$~ppm \\citep{kir03}. Analysis of data from the {\\it WMAP} and other cosmic microwave experiments yield $\\Omega_b h^2 = 0.0224$. Depending on the adopted nuclear reaction rates, (D/H)$_{\\rm prim} = 27.5^{+2.4}_{-1.9}$~ppm \\citep{cyb03} or $26.0^{+1.9}_{-1.7}$~ppm \\citep{coc04}.  Galactic chemical evolution models attempt to explain the decrease in D/H from the protogalaxy, for which D/H is presumed to be the primordial value, to the present epoch where D/H = (D/H)$_{\\rm LDtot}$. Nuclear reactions in stars and subsequent return of D-poor and metal-rich gas by SN explosions and stellar winds will reduce deuterium in the ISM by the factor $f_d \\equiv {\\rm (D/H)}_{\\rm prim}/{\\rm (D/H)}_{\\rm LDtot}$. Models of \\citet{chi02}, for example, predict that $f_d \\simeq 1.5$ for a wide range of assumptions concerning the value of (D/H)$_{\\rm prim}$ and amount of mixing in stellar interiors. This result can be compared with our results, $f_d \\leq (27.5^{+2.4}_{-1.9})/(23.1\\pm 2.4) \\leq 1.19^{+0.16}_{-0.15}$ or $f_d \\leq (26.0^{+1.9}_{-1.7})/(23.1\\pm 2.4) \\leq 1.12\\pm 0.14$. These two results differ by $\\geq 1.5\\sigma$ and $\\geq 1.8\\sigma$, respectively, from the predictions of the \\citet{chi02} models. If, on the other hand, the low value of (D/H)$_{\\rm gas} = 8.6\\pm 0.8$~ppm for the five most distant lines of sight is representative of (D/H)$_{\\rm LDtot}$ as \\citet{heb05} proposes, then $f_d$ would be $(27.5^{+2.4}_{-1.9})/(8.6\\pm 0.8) = 3.2^{+0.41}_{-0.37}$ or $(26.0^{+1.9}_{-1.7})/(8.6\\pm 0.8) = 3.0\\pm 0.35$, which are far larger values than the \\citet{chi02} models predict. Very recent Galactic chemical evolution models \\citep{rom06} can now accomodate values of $f_d$ as small as $\\sim$1.3, consistent with our determination $f_d < 1.19^{+0.16}_{-0.15}$ or $1.12\\pm0.14$, depending on the adopted value of (D/H)$_{\\rm prim}$. However, our empirical results for $f_d$ are smaller than $f_d = 1.39$, the lowest value for a model discussed by \\citet{rom06}. We suggest that new Galactic chemical evolution models should be developed to see whether smaller values of $f_d$ are consistent with the other empirical abundance and theoretical constraints that the models must fit.  We also call attention to the need for more high quality data to better constrain (D/H)$_{\\rm LDtot}$ and (D/H)$_{\\rm prim}$. In particular, we need better quality observations of the highly saturated Lyman lines for which measurements of $N$(H~I) may contain systematic errors of an unknown nature. Also, the assumptions underlying the Galactic chemical evolution models should be reexamined to see whether $f_d$ might be as small as $\\sim$1.25 without violating other constraints."
    },
    "0608/hep-ph0608067_arXiv.txt": {
        "abstract": "\\noindent We relate possible cosmological variations in the mass ratio $\\mu\\equiv m_p/m_e$ and the fine structure constant $\\alpha$ to long-range composition-dependent forces mediated by a scalar field. The differential acceleration $\\eta$ in E{\\\" o}tv{\\\" o}s-type experiments is bounded below by $10^{-14}$, except in cases where one or more scalar couplings vanish. We consider what values for these couplings could arise from unified theories. By considering the contribution of the scalar field to the cosmological energy density we use bounds on $\\eta$ to put upper bounds on the current rate of change of $\\mu$ and $\\alpha$. ",
        "introduction": "\\noindent The constancy of the coupling strengths and particle masses in the Standard Model Lagrangian is an assumption that should be tested. If a variation existed it would violate the principle of Local Position Invariance contained in the Einstein equivalence principle. Nevertheless consistent, relativistically covariant theories with varying couplings can be written down by including scalar degrees of freedom, in which the variation arises due to the cosmological solution for the scalar(s). See for example \\cite{Bekenstein,OliveP,Wetterich02,BarrowElectron}.  There are also many situations in cosmology where gravitational effects appear to behave differenty from what a naive application of GR would produce, which may also motivate the existence of cosmological scalar fields. The first is inflation (see for example \\cite{WMAPcosmo}), which seems to require a scalar field with a peculiar kind of potential. Closely related, though happening on an immensely slower timescale, is the current apparent accelerated expansion of the Universe \\cite{supernovae,WMAPcosmo}, which may also be related to scalar field dynamics \\cite{cosmonquint,Wetterich02}.  The violation of the Einstein equivalence principle makes itself via scalar mediated forces, thus gravity becomes a generalised scalar-tensor theory. % So far, precision tests of gravity have revealed no deviation from General Relativity \\cite{Will01,Nordtvedt02}, in particular concerning the question we will discuss here, the equivalence of gravitational and inertial mass (universality of free fall) \\cite{Su:1994gu}. The motivation for continuing to consider deviations from GR then comes into question.  The current motivation comes from astrophysical measurements of various coupling strengths and mass ratios deduced from atomic and molecular absorption (or emission) wavelengths. These probe back to redshifts up to about 4 and have now reached the sensitivity of $10^{-5}$ or better; moreover some measurements of the fine structure constant $\\alpha$ \\cite{Webbalpha}, and most recently the proton-electron mass ratio $\\mu \\equiv m_p/m_e$ \\cite{mu} show potentially significant deviations. The fine structure constant has been the subject of debate, with other recent studies obtaining negative results \\cite{novary,levshakov}, some at a precision greater than or equal to previous evidence for variation.  Bounds on variation of $\\alpha$ from nuclear physics effects \\cite{Oklo,Olive_et_al} have also been derived; however they are subject to considerable uncertainty if one also allows other quantities such as quark masses to vary (see for example \\cite{FlambaumS}).  It was pointed out in \\cite{DvaliZ} that a fractional variation of $\\alpha$ on the order of $10^{-5}$ required a scalar coupling to electromagnetic energy large enough such that the differential acceleration of two bodies in the gravitational field of the Sun (including scalar forces) is of order $10^{-18}$, which may be detectable with the STEP experiment. A more detailed analysis of scalar field models in \\cite{OliveP}, including also the electrostatic energy of nuclei, found a rather more restrictive bound from non-universal free fall, generally referred to as Weak Equivalence Principle (WEP) violation, although. Other works have related a variation in $\\alpha$ to scalar field models which describe the current acceleration of the Universe \\cite{Anchordoqui:2003ij,Parkinson:2003kf,Doran:2004ek}.   As pointed out in \\cite{Wetterich02}, one reason for believing the result of \\cite{DvaliZ} to be an underestimate of WEP violation is that the authors did not include possible scalar couplings to other Standard Model sectors which are likely in unified models (see also \\cite{C+O,Calmet,Langacker,us,Dine,me03}). Clearly also a scalar coupling to electromagnetism alone cannot produce a variation in $\\mu$ at the same level as that in $\\alpha$.  Hence in this paper we redo and extend the analysis to include other possible types of scalar coupling through which varying $\\alpha$ and $\\mu$ can be related to WEP violation. Current limits on differential gravitational acceleration are at the level of $10^{-12}$ -- $10^{-13}$ \\cite{Su:1994gu} % for the parameter $\\eta$ defined as \\beq \\label{etadef} \\eta \\equiv 2\\frac{|a_1-a_2|}{|a_1+a_2|} \\eeq where two test bodies of different composition have accelerations $a_1$, $a_2$ towards a known gravitational source.  In the first section of the paper we make a preliminary estimate of the size of differential acceleration $\\eta$ due to $\\phi$-mediated forces, allowing for varying $\\mu$, and find it may be orders of magnitude greater than $10^{-18}$, and possibly visible with the MICROSCOPE experiment with sensitivity $10^{-15}$. Exceptions occur if some scalar couplings vanish exactly.  In the second section we consider what scalar couplings consistent with varying $\\mu$ may emerge from unified models both with and without supersymmetry, using the detailed analysis of \\cite{me03}. We find it unlikely that the scalar coupling to any particle will vanish. In fact some unification scenarios imply that WEP violation may be at the current limits of detection, or even be ruled out by E{\\\" o}tv{\\\" o}s-type experiments. % Only a restricted range of unified scenarios can be consistent with reported variations of $\\alpha$ and $\\mu$, if these are to arise from a scalar field varying over cosmological time.  In the last section we examine the cosmological evolution of the scalar field in more detail and suggest a direct test of unified scenarios by comparing WEP violation with atomic clock measurements of the constancy of fundamental parameters.  \\subsection{Consistency of ``varying constant'' observations}  Clearly, the experimental situation on possible variation of $\\alpha$ is currently unsatisfactory, while that for $\\mu$ requires further independent testing. We discuss two points: first, whether nonzero variations in both $\\alpha$ and $\\mu$ are inconsistent with other measurements of dimensionless fundamental quantities; second, whether recent results on $\\alpha$ are directly comparable.  The quantities $y\\equiv \\alpha^2 g_p$ and $x \\equiv \\alpha^2 g_p m_e/m_p$ have been measured with good precision using quasar absorption \\cite{Murphy01,Tzanavaris}. From two systems at redshifts below $0.7$ the fractional variation of $y$ is bounded at the level of $5\\times 10^{-6}$, hence is not inconsistent with the claim $\\Delta \\alpha = (-0.57 \\pm 0.11) \\times 10^{-5}$ (made for systems over a wide range in redshift) given constant $g_p$.~\\footnote{Possible variation of $g$-factors was discussed in \\cite{Dmitriev:2002kv}; any effect on $g_p$ is expected to be small.} The result for $x$ indicates a slight but not significant {\\em increase} into the past, $\\Delta x/x = (1.2 \\pm 1.0) \\times 10^{-5}$. Again given constant $g_p$, this is difficult to reconcile with a decrease in $\\alpha$ and an increase in $\\mu$ as claimed in \\cite{mu} at the level $\\Delta \\mu/\\mu = (2.4 \\pm 0.6) \\times 10^{-5}$. Hence we cannot assume that both $\\alpha$ and $\\mu$ have exactly the claimed nonzero variation. Rather we will make estimates assuming that at least one nonzero variation exists (usually $\\mu$) at the claimed level.  The most recent negative results for $\\Delta\\alpha$ are based on very few absorption systems, the improved sensitivity being ascribed to more accurate spectroscopy. The Webb group results by contrast involve a sample of about 150 systems.\\footnote{Analysis of more recent observations is ongoing (M.~Murphy, private communication).} This makes it possible to compare the estimated error in each system% statistical scatter over the sample, and carry out checks against systematic error. One outstanding question is the statistical error due to fitting the velocity profile of the absorbing system; see for example the first reference of \\cite{levshakov}. It may require a large number of systems to be analyzed to settle the question. ",
        "conclusions": "We have investigated bounds arising from violations of the Weak Equivalence Principle in scalar field theories in which the proton-electron mass ratio $\\mu$ may vary. We related such bounds to unified theories, astrophysical observations of quasar absorption spectra and atomic clock measurements of the constancy of physical parameters. Model-independent bounds on the rate of change of the scalar field and on its total variation at any given epoch are obtained from the effect of its kinetic energy on the expansion of the Universe.  For a given nonzero variation of $\\mu$, the scalar couplings to the proton or electron mass are bounded below, thus so is the resulting differential acceleration, up to a single parameter $K/c_2^2$. This parameter is a ratio of scalar coupling constants, which arises from how the nucleon and electron masses are related in the underlying particle theory. Hence such underlying theories, in particular gauge unification at high energy, could be tested with variation of $\\mu$.  Since the value of $\\alpha$ may also vary in such underlying theories, we also analyzed its variation in relation to WEP violation, extending previous works. We found that in many scenarios for generating varying $\\mu$ from unified models, the differential acceleration parameter $\\eta$ is near or above its current experimental limit. In particular, in scenarios where due to cancellations the variation of $\\mu$ is not much larger than that of $\\alpha$, the contributions to differential accelerations do not cancel and the violation of WEP is at or above the experimental limit. Thus it seems very difficult to reconcile currently claimed nonzero variations of $\\mu$ and $\\alpha$ within unified models.  Our results differ from those of \\cite{BarrowElectron}, where the cosmological variation of $\\mu$ was bounded to be many orders of magnitude below current levels of sensitivity. The main difference is that in \\cite{BarrowElectron} the source of the scalar field variation was taken to be the coupling to the local electron density, whereas we allow the field to be driven by a potential energy and in principle by couplings to other constituents of the Universe. This illustrates the point that in models where the form of the scalar Lagrangian is known or assumed, more stringent bounds may be found. Our aim was however to find bounds which are as model-independent as possible.   \\subsection*"
    },
    "0608/astro-ph0608398.txt": {
        "abstract": "The Bardeen-Petterson general relativistic effect has been suggested as the mechanism responsible for precession in some accretion disk systems. Here we examine separately four mechanisms (tidally-induced, irradiation-induced, magnetically-induced and Bardeen-Petterson-induced) that can lead to warping and precession. We use a sample of eight X-ray binaries and four Active Galactic Nuclei (AGNs) that present signatures of warping and/or precession in their accretion disks to explore the viability of the different mechanisms. For the X-ray binaries SMC\\,X-1 and 4U\\,1907+09 all four mechanisms provide precession periods compatible with those observed, while for Cyg\\,X-1 and the active galaxies Arp\\,102B and NGC\\,1068, only two mechanisms are in agreement with the observations. The irradiation-driven instability seems incapable of producing the inferred precession of the active galaxies in our sample, and the tidally-induced precession can probably be ruled out in the case of Arp\\,102B. Perhaps the best case for a Bardeen-Petterson precession can be achieved for NGC\\,1068. Our results show that given the many observational uncertainties that still exist, it is extremely difficult to confirm unambiguously that the Bardeen-Petterson effect has been observed in any of the other sources of our sample. ",
        "introduction": "The importance of accretion processes in the line formation and continuum emission in AGNs and galactic sources, such as microquasars and X-ray binaries, has long been recognized in the literature, even though the presence of an accretion disk in some of those systems is inferred only indirectly.  The continuous improvement in the capability of telescopes and detectors to obtain high-resolution images, and spectra with higher sensitivity, allowed for the direct probing of the physical characteristics of some accretion disk systems (e.g., \\citealt{ray96,jon00,jaf04,fat06}).  In contrast to the standard picture of a flat disk surrounding the accreting object, some galactic and extragalactic sources present signatures of warping and precession in their disks and/or jets (e.g., \\citealt{marg84,sil88,abra00,gal04,cap06}). The apparent lack of correlation between the orientation of the radio jets and the plane of the host galaxy's disk might also be attributed to warped disks (e.g., \\citealt{sch02}).  The Bardeen-Petterson effect \\citep{bape75}, a gravitational perturbation predicted by general relativity, is one of the physical mechanisms that has been proposed to explain warping and precession in accretion disks (e.g., \\citealt{scfe96,nepa00,fran05,kin05,cap06}). In order to explore whether this mechanism is universal, in the sense that it can be acting in both galactic and extragalactic accretion disks, we have selected eight X-ray binaries and four AGNs that present signatures of warping/precession in their disks. We have analyzed individually the predictions of three additional precession mechanisms for the same sample of sources, in order to compare these predictions with those from the Bardeen-Petterson effect. In Section 2, we describe the accretion disk model used throughout the paper, as well as constraints on the basic parameters of the accretion disk. In Section 3, we present the four warping/precession mechanisms studied in this work. A brief introduction of the sample of sources, as well as their basic parameters (inferred observationally) is given in Section 4. A comparison between the theoretical predictions and observations is carried out in Section 5, and conclusions follow. ",
        "conclusions": "In this work, we have selected eight X-ray binaries and four AGNs that present signatures of precession in their accretion disks in order to examine the compatibility between their precession periods and the predictions from four distinct warping/precession physical mechanisms: tidal torques from a companion in a binary system, irradiation- and magnetically-driven instabilities, and the Bardeen-Petterson relativistic effect.  We have assumed a power-law surface density distribution for the accretion disks, constraining their physical parameters from observational data available in the literature.  For the X-ray binaries in our sample, we found that tidal torques from a companion in a binary system provide precession timescales compatible with those inferred in SS\\,433, Her\\,X-1, LMC\\,X-4 and SMC\\,X-1, as indeed has been previously suggested in the literature \\citep{katz73,larw98}. In addition, we showed that tidal torques can also drive precession in 4U\\,1907+09, which (as far as we know) had not been proposed before this work. The mechanism can be ruled out for Cen\\,X-3 and Cyg\\,X-2. In the case of our AGN sample (assuming that they contain binary black holes), we have shown that tidal torques cannot produce the disk precession inferred in Arp\\,102B.  Although the irradiation-driven instability usually provides precession timescales that agree with those observed in our X-ray binary sample, it is unsuccessful in reproducing the reported precession periods for the four AGNs considered here. The critical radii at which those AGN accretion disks become unstable against radiation torques were found to be larger than the expected disks' outer radii (for $\\eta\\ne1$). Since the irradiation-driven instability produces prograde disk precession, the retrograde precession observed in SS\\,433, Her\\,X-1 and LMC\\,X-4 may only be generated in the presence of some external torques.  We have shown that torques due to a misaligned magnetic field (with respect to the perpendicular direction to an accretion disk) can also reproduce the observed precession periods of our sample. For the accreting neutron stars, a dipolar configuration must be a good approximation to their magnetic fields; in such a case (or for $\\chi<-1$), the magnetic-driven torquing can induce precession in the inner parts of the accretion disks at the rate inferred observationally. For the black hole X-ray binaries, magnetic configurations with $\\chi\\geq -1$ are favored. Although the geometry of the magnetic field is not usually constrained by observations in AGNs, our results indicate the existence of two distinct regimes: for accretion disks with $s\\lesssim -1$, the z-component of the magnetic field must be radially constant in order to reproduce the observed precession periods; otherwise, it is necessary to have a constant surface density ($s=0$) with a magnetic field weakening radially along the disk ($\\chi\\leq -1$). In addition, if the accretion disk of Cyg\\,X-1 is actually precessing progradely \\citep{lac06}, we can rule out magnetically-driven torques as the cause of its precession.  The Bardeen-Petterson effect produces precession timescales compatible with those observed in all sources of our sample, considering that the alignment of the accretion disk evolves on a similar timescale (as in \\citealt{scfe96}). The timescale for reaching the observed precession rate is usually shorter than the estimated lifetime of the sources. Considering the accretion disk parameters used in this work, we have found that all sources in our sample have Bardeen-Petterson radii smaller than the disk's outer radius.  NGC\\,1068, for which maser observations have been used to infer the physical properties of the accretion disk, may provide the strongest candidate for the Bardeen-Petterson effect. In this case, we have shown that the irradiation-driven instability is incompatible with the observations. Tidal torques are also highly unlikely (even if a binary black hole were to be present in this system) since it would most likely result in other changes to the disk morphology, which are not observed. Magnetically-driven torques would work only if a rather contrived field configuration exists ($B_\\mathrm{Z}$ constant along the precessing part of the disk). Our conclusions give further support to the analysis of \\citet{cap06}, who showed that the Bardeen-Petterson effect is consistent with the warping morphology suggested by \\citet{gal04}.  The Bardeen-Petterson precession is an important general-relativistic effect, and its unambiguous detection would be of great value. This work has shown, however, that given the many observational uncertainties that still exist (both in X-ray binaries and in AGNs), it is extremely difficult to rule out definitively other precession mechanisms.  It is important to emphasize that although we have analyzed separately the four precession mechanisms in this work, they are not mutually exclusive. This means that more than one mechanism might be operating in a given system.  %% If you wish to include an acknowledgments section in your paper, %% separate it off from the body of the text using the"
    },
    "0608/astro-ph0608504_arXiv.txt": {
        "abstract": "We derive an empirical effective temperature and bolometric luminosity calibration for G and K dwarfs, by applying our own implementation of the InfraRed Flux Method to multi--band photometry. Our study is based on 104 stars for which we have excellent $BV(RI)_C JHK_S$ photometry, excellent parallaxes and good metallicities.  Colours computed from the most recent synthetic libraries (ATLAS9 and MARCS) are found to be in good agreement with the empirical colours in the optical bands, but some discrepancies still remain in the infrared. Synthetic and empirical bolometric corrections also show fair agreement.  A careful comparison to temperatures, luminosities and angular diameters obtained with other methods in literature shows that systematic effects still exist in the calibrations at the level of a few percent. Our InfraRed Flux Method temperature scale is 100~K hotter than recent analogous determinations in the literature, but is in agreement with spectroscopically calibrated temperature scales and fits well the colours of the Sun.  Our angular diameters are typically 3\\% smaller when compared to other (indirect) determinations of angular diameter for such stars, but are consistent with the limb-darkening corrected predictions of the latest 3D model atmospheres and also with the results of asteroseismology.  Very tight empirical relations are derived for bolometric luminosity, effective temperature and angular diameter from photometric indices.  We find that much of the discrepancy with other temperature scales and the uncertainties in the infrared synthetic colours arise from the uncertainties in the use of Vega as the flux calibrator.  Angular diameter measurements for a well chosen set of G and K dwarfs would go a long way to addressing this problem. ",
        "introduction": "\\label{intro}  Temperatures, luminosities and radii are amongst the basic physical data against which models of stellar structure and evolution are tested.  In this paper, we address one particular area of the Hertzsprung-Russell (HR) diagram, that of the G and K dwarfs --- focusing particularly on stars for which the effects of stellar evolution have been negligible or nearly negligible during their lifetimes. Somewhat higher mass stars than those considered here have been extensively studied historically because their luminosities may be used to infer stellar ages. Lower mass stars of late-G to K spectral types, have been neglected to some extent, probably because there have been few secondary benefits in getting stellar models right for these stars, and the lack of good parallax, diameter and other data. This situation is rapidly changing: firstly, Hipparcos has provided the requisite parallax data; secondly, interferometric techniques are making the measurement of diameters for such small stars likely to be routine within a few years; thirdly, we have our own motivation in the form of an ongoing project to follow the chemical history of the Milky Way from lower mass stars, for which we can infer indirectly the amount of their constituent helium via stellar luminosity (Jimenez et al 2003). In order to achieve our goals, we need accurate and homogeneous effective temperatures and luminosities for our G+K dwarfs.  The effective temperature of a stellar surface is a measure of the total energy, integrated over all wavelengths, radiated from a unit of surface area. Since its value is fixed by the luminosity and radius, it is readily calculated for theoretical stellar models, and as one of the coordinates of the physical HR diagram, it plays a central role in discussions of stellar evolution. Most observations, however, provide spectroscopic or photometric indicators of temperature that are only indirectly related to the effective temperature.  The effective temperature of lower main sequence stars is not easy to determine and different measurement techniques are still far from satisfactory concordance (e.g. Mishenina \\& Kovtyukh 2001, Kovtyukh et al. 2003). At present, spectroscopic temperature determinations return values that are some 100~K hotter that most of the other techniques.  Even in a restricted and thoroughly studied region like that of the Solar Analogs, effective temperature determinations for the same star still differ significantly, by as much as 150~K (Soubiran \\& Triaud, 2004). Models predict values that are some 100~K hotter than those measured (Lebreton et al. 1999 and references therein).  In this work, we apply the Infrared Flux Method to multi-band $BV(RI)_{C}JHK_S$ photometry of a carefully selected sample of G and K dwarf stars.  We compare observed with synthetic broad-band colours computed from up-to-date (1D) model atmospheres. Such models are then used to estimate the missing flux needed to recover bolometric luminosities from our photometry. Other than the high quality of the observational data, the strength of this work relies on using very few basic assumptions: these are the adopted Vega absolute calibration and zero-points. This also makes clearer the evaluation of possible errors and/or biases in the results.  Both the absolute calibration and the zero-points are expected to be well known and the latest generations of model atmospheres produce realistic fluxes for a wide range of temperatures, gravities and abundances (Bessell 2005) so that the adopted model and calibration are at the best level currently available.  The paper is organized as follows: in Section 2 we describe our sample; we compare different libraries of model atmospheres with observational data in Section 3 and in Section 4 we present our implementation of the Infrared Flux Method along with the resulting temperature scale. In Section 5 we test our scale with empirical data for the Sun and solar analogs and with recent interferometric measurements of dwarf stars. The comparison with other temperature determinations is done in Section 6 and in Section 7 and 8 we give bolometric calibrations and useful tight relations between angular diameters and photometric indices. We conclude in Section 9. ",
        "conclusions": "We have used the Infrared Flux Method to deduce fundamental stellar parameters of effective temperature, bolometric magnitude and angular diameter for a sample of 104 G and K dwarfs. This semi-direct method is mostly based on empirical data obtained from our own observations or carefully selected from the literature in order to achieve the highest accuracy feasible. Our stars all have excellent $BV(RI)_C JHK_S$ photometry, excellent parallaxes and good metallicities.  Most of the bolometric flux of our stars is seen in the optical and infrared bands covered $BV(RI)_C JHK_S$. For the remaining 15 to 30\\% of the flux lying outside these bands, model atmospheres have been used. Good to very good agreement is found for the colours of the stars and the synthetically derived colours of a number of sets of model atmospheres in the literature (ATLAS9, MARCS, BaSel 3.1).  The zero-points of our temperature, luminosity and angular diameter scales (Sections \\ref{irfm}, \\ref{BolSec} and \\ref{AngCal}) depend only on our adopted absolute calibration of Vega. The advantages and disadvantages of our use of Vega are discussed in Appendix A. The main advantage is that Vega's absolute calibration has been extensively studied over a wide wavelength range via satellite measurements. This allows us to put firm constraints on the systematic uncertainties. A likely disadvantage is that Vega is known to be pole-on and a fast rotator which prevents the use of a unique Vega model atmosphere in fitting the observations.  Our calibration of effective temperature is found to be some 100~K hotter than similar determinations in the literature, while it is also found to be in close agreement with effective temperature scales based on spectroscopic methods. To reduce our temperature scale by 100 K, the required changes to Vega's 2MASS zero-points and/or absolute calibration are significantly larger than the known uncertainties in these quantities. Neither can the model atmospheres be blamed, because we find such good agreement between the model atmosphere colours and the empirical data.  We worked hard to impute inconsistency to one or other of the scales, but were unable to resolve the issue beyond any reasonable doubt. Our angular diameters are smaller by about 3\\% with respect to other indirect determinations, but they seem more in line with the predictions of asteroseismology and 3D model atmospheres. We conclude that temperatures, luminosities and angular diameters calibrations for lower main sequence G and K dwarfs retain systematics of the order of a few percent.  The high quality and homogeneity of the data produce very tight empirical colour--metallicity--temperature and angular diameter--colour calibrations. In particular, the relation between angular diameters and magnitudes in the $J$ band is remarkable and it indicates an high sensitivity of this band to angular diameters. Since many lower main sequence stars are known to host planets, our relations can be used to accurately determine the physical properties of the parent star thus allowing to effectively break the degeneracy between the properties of the stars and that of the planets when extremely high precision data are unavailable (e.g. Bakos et al. 2005, Sato et al. 2005).  Future interferometric measurements of angular diameters would go a long way to addressing the uncertainties we have shown as well as to test our findings. In particular, the shape of the visibility function in the second lobe would directly probe limb-darkening corrections (e.g. Bigot et al. 2006) for our stars. Likewise, tying the absolute calibration and the observed colours to a solar twin would probably reduce many of these uncertainties. At present good candidate solar twins are HD 146233 (Porto de Mello \\& da Silva) and HD 98618 (Mel\\'endez et al. 2006). Accurate interferometric measurements together with high precision photometry for a set of nearby solar-like stars would permit to set the absolute calibration via solar analogs directly (e.g. Campins et al. 1985). Besides, extremely high precision multiple bandpass studies of transiting extra-solar planets would provide direct angular diameter estimates for our stars, and will further test our temperature scale."
    },
    "0608/astro-ph0608218_arXiv.txt": {
        "abstract": "{\\HHth{} is a well-known Pre-Main-Sequence star in Taurus.  HST observations have revealed a flared, edge-on disk driving a highly-collimated optical jet, making this object a case study for the disk-jet-outflow paradigm.} {We searched for a molecular outflow, and attempted to better constrain the star and disk parameters.} {We obtained high angular resolution ($\\about{}1''$) observations of the dust continuum at 2.7 and 1.3~mm, and of the \\twCO{}~\\Jtwo{}, \\thCO{}~\\Jtwo{} \\&~\\Jone{}, \\CeiO{}~\\Jone{}~emissions around \\HHth{}. A standard disk model is used to fit the \\thCO{} \\Jtwo{} $uv$-plane visibilities and derive the disk properties, and the stellar mass. An \\emph{ad hoc} outflow model is used to reproduce the main properties of the \\twCO{}~\\Jtwo{} emission.} {The rotation vector of the disk points toward the North-Eastern jet. The disk rotation is Keplerian: Using a distance of 140\\pc{}, we deduce a mass of 0.45\\Msun{} for the central star. The disk outer radius is 420\\au{}. A highly asymmetric outflow originates from the inner parts of the disk. Only its North-Eastern lobe was detected: it presents to first order a conical morphology with a $30\\degr{}$ half opening angle and a constant (12\\kms{}) radial velocity field. Outflow rotation was searched for but not found. The upper limit of the outflow rotation velocity is 1\\kms{} at 200\\au{} of the jet axis.} {\\HHth{} is a low mass TTauri of spectral type around M1 and age 1 to 4~Myrs, surrounded by a medium size Keplerian disk, of mass around $4\\, 10^{-3}\\Msun$. It beautifully illustrates the jet-disk-outflow interaction, being so far the only star to display a jet and outflow connected to a well defined Keplerian disk, but reveals a surprisingly asymmetric (one-sided) outflow despite a relatively symmetric jet. Furthermore, these observations do not enable to assign the origin of the molecular outflow to entrainment by the optical jet or to a disk wind. In the latter hypothesis, the lack of rotation would imply an origin in the inner 15\\au{} of the disk.} ",
        "introduction": "\\TableObs{} %  During the last decade, the properties of the circumstellar environment of pre-main-sequence (PMS) stars have been extensively studied. Millimeter arrays, such as the IRAM Plateau de Bure interferometer, routinely provide CO line and continuum images of low-mass PMS stars which show that many of them are surrounded by large ($R_\\emr{out}\\sim 200-800\\au$) disks in Keplerian rotation \\citep[\\eg{}][]{koerner93,dutrey94,mannings97,guilloteau98,simon00}. For several sources, the angular resolution and sensitivity of the observations enable to fit disk models to the data, hence constraining physical parameters of the disk, and even the mass of the central source~\\citep[][and references therein]{dutrey06}. While mm continuum observations are sensitive to the dust thermal emission of the disk, near-infrared (NIR) or optical observations can be used to trace the stellar light scattered by particles at the disk surface \\citep[\\eg{}][]{roddier96}. If the system is edge-on, the disk itself can be seen as a dark lane on the equatorial plane (\\eg{} \\HHth{}:~\\citet{burrows96}, GM Auriga:~\\citet{stapelfeldt97}).  Most of pre-main-sequence stars also show an ejection of matter.  In several sources, optical linear jets are observed through the emission of forbidden atomic lines. Interestingly, \\citet{bacciotti03} and \\citet{coffey04} have detected possible rotation signatures in at least four jets (DG\\,Tau, TH\\,28, RW\\,Aur, LkH$\\alpha$\\,321): They have observed systematic velocity differences between opposite sides of the jet axis. Indeed, rotation of the jet is predicted by magnetocentrifugal models of accretion-ejection~\\citep{konigl00,shu00}, through the coupling of the outflowing gas to the helicoidal structure of the magnetic field. In younger, more embedded sources, large molecular outflows are observed, mainly through the emission of the rotational lines of CO and its isotopologues~\\citep[][and references therein]{arce06}. Those molecular outflows are usually identified with ambient molecular gas that has been entrained and thus put into motion by the underlying protostellar jet. In any case, it is particularly interesting to check whether a rotation component is present in molecular outflows.  In this context, one of the most interesting objects is the young PMS star \\HHth{}, located in the Taurus molecular cloud at a distance of $\\sim 140\\pc$.  \\citet{burrows96} obtained spectacular images with the Hubble Space Telescope that revealed a flared edge-on disk of diameter $\\sim 450\\au$, separated in two parts by a dust lane, and a highly-collimated jet emanating out of the central region. Since then, the \\HHth{} disk and jet have been studied in many details \\citep[\\eg{}][]{bacciotti99, stapelfeldt99a, wood00, cotera01, wood02, watson04}. \\citet{stapelfeldt99a} reported variability and asymmetry in the emission from the circumstellar disk by comparing HST optical observations performed between 1994 and 1998; the photometric variability of the star and its possible effects on the disk was investigated by~\\citet{wood00}. This source is a unique candidate to probe the complex interaction between the disk, the jet and the outflow.  However, \\HHth{} has been poorly investigated at mm wavelengths so far. This is probably due to the confusion arising from the \\HHth{} parent molecular cloud, which increases the complexity of any detailed observations using the emission of CO isotopologues. The \\twCO{} emission of the cloud is expected to be optically thick, hence hiding significantly the disk emission. BIMA observations of the \\thCO{}~\\Jone{} emission in the environment of HL~Tau~\\citep[Fig.~2 of][]{welch00} shows that \\HHth{} is located in the south edge of a shell driven by XZ~Tau. Still, the remarkable properties of \\HHth{} makes it a textbook case for the study of the disk/jet system in young PMS stars, and we have therefore decided to use the IRAM Plateau de Bure interferometer to perform high angular resolution ($\\about{}1''$) observations of the emission of the circumstellar material surrounding \\HHth{} in several CO isotopologues lines. In this paper we present the results of this study and the analysis of the thermal dust and molecular line emissions of the disk and the outflow. ",
        "conclusions": "In the \\HHth{} case, a judicious use of the complementary line emission allowed us to derive a number of very robust results: 1) The systemic velocity of \\HHth{} is 7.25\\kms{}; 2) The disk outer radius is 420\\au{}; 3) The rotation vector of the disk points toward the North-Eastern jet; 4) The disk is in Keplerian rotation; 5) The stellar mass is 0.45 \\Msun; 6) A highly asymmetric molecular outflow originates from the inner parts of the disk; 7) The outflow material is mainly located on the thin edges of a cone with an opening angle of $30\\degr$; 8) The outflow velocity is essentially radial with a magnitude of $12\\kms$; 9) No rotation of the outflow is detected.  The total disk mass is less certain: \\about{}$4\\,10^{-3}$\\Msun{}. This depends on the assumed dust emissivity. Our measurement indicate a rather low value of $\\beta$, 0.4, suggestive of significant grain growth. In addition, but with still less certainty, the disk appears cold, around 12 K at 100 AU, suggesting a significant temperature gradient between the disk plane and the surface layers. Despite this low temperature, there is no measurable depletion of CO compared to the average abundance in the Taurus cloud, in contrast with the general trend for all other circumstellar disks~\\citep[except those around HAe stars, see ][]{Pietu_etal2006}.  The direct determination of the stellar mass improves the evolutionary status of HH 30, but this remains hampered by the uncertainties on the total luminosity. The most plausible solution is that of a \\about{}4~Myr old object, with a total luminosity of order 0.2 $\\Lsun$, but a younger and brighter star cannot be excluded.  Finally, the origin of the molecular outflow remains unclear, with no conclusive arguments to distinguish between an entrainment mechanism by the optical jet or a disk wind. The one sided nature of the outflow remains an unexplained issue in both hypothesis."
    },
    "0608/astro-ph0608168_arXiv.txt": {
        "abstract": "Observations of the Milky Way by the SPI/INTEGRAL satellite have confirmed the presence of a strong 511~KeV gamma-ray line emission from the bulge, which require an intense source of positrons in the galactic center. These observations are hard to account for by conventional astrophysical scenarios, whereas other proposals, such as light DM, face stringent constraints from the diffuse gamma-ray background. Here we suggest that light superconducting strings could be the source of the observed 511~KeV emission. The associated particle physics, at the $\\sim$~1~TeV scale, is within reach of planned accelerator experiments, while the distinguishing spatial distribution, proportional to the galactic magnetic field, could be mapped by SPI or by future, more sensitive, satellite missions. ",
        "introduction": "The problem of the birth, propagation and annihilation of positrons in the Galaxy has been a major topic of astrophysical investigation, since the first detection\\cite{johnson} of the 511~KeV gamma-ray line signature. The SPI instrument on-board ESA's INTEGRAL satellite has established the presence of a diffuse source of positrons in the Galactic Center (GC)\\cite{SPI}. The observed photon flux of \\begin{equation} 9.9^{+4.7}_{-2.1} \\times 10^{-4}~{\\rm cm}^{-2} {\\rm s}^{-1} \\label{f0} \\end{equation} with a line-width of about 3~KeV is in good agreement with previous measurements\\cite{osse}.  For the spatial distribution of the 511~KeV line component, the mapping results point to an intense bulge emission, better explained by an extended distribution than by a point source. Assuming a Gaussian spatial distribution for the flux, a full-width at half maximum of $9^\\circ$ is indicated. The disk component has been either absent, or weakly detected in the initial results.  The origin of these Galactic positrons remains a mystery. Several scenarios involving astrophysical sources have been proposed, including neutron stars or black holes, massive stars, supernovae, hypernovae, gamma ray bursts or cosmic rays\\cite{astrof}. However, the fraction of positrons produced in such processes is uncertain, and it is unclear that the positrons could fill the whole bulge.  Alternatively, mechanisms associated with the Dark Matter~(DM) at the GC have been put forward. If DM is constituted by a light (~MeV) scalar, its decay or annihilation could account for the observed signal\\cite{Boehm:2003bt}.  The positrons should be injected at non-relativistic energies so that the associated bremsstrahlung emission does not violate the COMPTEL and EGRET measurements of diffuse radiation from the Galactic Center\\cite{beacom}. For the DM scenario, this implies that the DM particles should be lighter than $\\sim$~20-30~MeV. Moreover, inflight annihilation of positrons would also overproduce gamma-rays unless the positrons are injected at energies below $\\sim$~3~MeV\\cite{beacom2}, thus disfavoring some of the scenarios in\\cite{astrof} and\\cite{Boehm:2003bt}.  We will discuss here the possibility, proposed in\\cite{Ferrer:2005xv}, that a network of light superconducting strings\\cite{Witten:1984eb} occurring in particle physics just beyond the standard model could be a source of the galactic positrons. This scenario predicts a characteristic positron distribution that could be used to distinguish this source from the other possibilities.  Assuming that a tangle of superconducting strings exists in the Milky Way, then the strings are frozen in the plasma as long as the radius of curvature is larger than a certain critical length scale. If the curvature radius is smaller, the string tension wins over the plasma forces and the string moves with respect to the magnetized plasma. During the string motion, the loop will cut across the Milky Way magnetic field, generating current as given by Faraday's law of induction.  The current is composed of zero modes of charged particles, including positrons, propagating along the string. The external magnetic field shifts the modes of the charge carriers into the bulk and modifies the dispersion relation dramatically so that their energy remains below the threshold for expulsion\\cite{Ferrer:2006ne} (which for the positron zero modes is 511~KeV). An additional perturbation, like inhomogeneities in the magnetic field, string motion and curvature, or scattering by counter-propagating particles\\cite{Barr:1987ij}, ejects the zero modes at the threshold of 511~KeV. The ejected positrons will annihilate with the ambient electrons, thus emitting 511~KeV gamma rays. ",
        "conclusions": ""
    },
    "0608/astro-ph0608442_arXiv.txt": {
        "abstract": "We report on evolutionary calculations of the onset of mass transfer in AM CVn binaries, treating the donor's evolution in detail.  We show that during the early contact phase, while the mass transfer rate, $\\Mdot$, is increasing, gravity wave (GW) emission continues to drive the binary to shorter orbital period, $\\Porb$.  We argue that the phase where $\\Mdot > 0$ and $\\nudot > 0$ ($\\nu = 1/\\Porb$) can last between $10^3$ and $10^6$ yrs, significantly longer than previously estimated.  These results are applied to RX J0806+1527 ($\\Porb = 321$ s) and RX J914+2456 ($\\Porb=569$ s), both of which have measured $\\nudot > 0$.  \\emph{Thus, a $\\nudot > 0$  does not select between the unipolar inductor and accretion driven models proposed as the source of X-rays in these systems}.  For the accretion model, we predict for RX J0806 that $\\ddot{\\nu} \\approx \\ee{1.0-1.5}{-28}$ Hz s$^{-2}$ and argue that timing observations can probe $\\ddot{\\nu}$ at this level with a total $\\approx 20$ yr baseline. We also  place constraints on each system's initial parameters given current observational data. ",
        "introduction": "} The true nature of the two candidate ultracompact binary systems originally discovered with \\textit{ROSAT},  RX J0806+1527 \\citep{beuermann99} and RX J1914+2456 \\citep{haberl95}, has been a source of much controversy. Both systems are soft X-ray sources whose light curves are modulated on periods of 321 s (RX J0806, hereafter J0806) and 569 s (RX J1914, hereafter J1914) \\citep{motch96,israel99,burwitz01}. In both cases, the X-ray light is 100\\% modulated for roughly half the period \\citep{cropper98,israel99} and the optical light is modulated on the same period as the X-ray light \\citep{ramsay00,israel02,ramsay02a,ramsay02b,israel03} with little evidence for other periodicities \\citep[but see][]{ramsay06}. It is widely believed, based partially on the stability and singularity of these periods, that the modulations are on the systems' orbital periods, $\\Porb$ \\citep[although see][for an alternate interpretation]{norton04}. This would make J0806 and J1914 the two shortest period binaries known.   The X-ray production mechanism in these systems has been much debated. Two competing models have come to the fore: the unipolar inductor (UI) model  \\citep{wu02} and the direct-impact accretion model \\citep{marsh02}. Both model these systems as white dwarf-white dwarf (WD-WD) binaries with $\\Porb$ equal to the X-ray period. In the UI model, a magnetized, more massive primary WD spinning asynchronously with respect to the orbit induces an electric field in the secondary WD and drives a current between the two. X-rays are produced by resistive dissipation in the primary's atmosphere.  In the accretion model, the secondary fills it Roche Lobe ($\\RL$) and mass transfer occurs. The binary's compact geometry leads to the accretion stream directly impacting the primary and no accretion disk forms. Both models lead to a spatially small X-ray production site on the primary, explaining, in principle, the observed  X-ray modulation \\citep[see, however][]{barros05}.  One potential means of distinguishing between the two models is the \\emph{secular} time derivative of $\\Porb$, or equivalently of the orbital frequency, $\\nu = 1/\\Porb$, $\\nudot$. The UI model predicts $\\nudot > 0$. Assuming fully degenerate donors and enforcing exactly that the donor's radius, $R_2$ equals its Roche radius: $\\RL = R_2$, the accretion model would predict $\\nudot < 0$.  Determinations of $\\nudot$ from both optical \\citep{hakala03,hakala04,ramsay05} and X-ray observations \\citep{strohmayer02,strohmayer03,strohmayer04,strohmayer05} \\emph{all} show that $\\nudot > 0$ over the entire epoch of observations. On the assumption that these results reflect the \\emph{secular} $\\nudot$, this has been taken as evidence against accretion models \\citep[e.g.;][]{israel03,strohmayer02,strohmayer04,ramsay05}.  However, the $R_2 = \\RL$ constraint requires a very large mass transfer rate $\\Mdot \\approx 10^{-7} - 10^{-5} \\msun$ yr$^{-1}$ and if $\\Mdot$ is significantly below this ``equilibrium'' value, observing a $\\nudot > 0$ is possible.  The $\\Mdot$ could be lower than expected either due to a non-secular mechanism \\citep{marsh05} or during the mass transfer turn-on phase \\citep{willems05}.  However, the timescales on which the binary (re)establishes its equilibrium $\\Mdot$ are estimated to be very short: 100-1000 yrs \\citep{marsh05} or 2-20 yrs \\citep{willems05}, making observing such phases unlikely.  Here we reexamine the suggestion J0806 and J1914 are direct-impact accretors seen during the $\\Mdot$ turn-on phase in the context of stellar evolution models developed to realistically address the complete evolution of AM CVn binaries.  In \\S 2, we show that the population of AM CVns forming through a double-degenerate channel \\citep[see][]{nelemans01,deloye05} naturally produces systems whose contact phase evolution give $\\Porb$ and $\\nudot$ consistent with both J0806 and J1914. The maximum $\\Mdot$ is sufficiently high in these systems that $\\RL$ must penetrate below the donor's photosphere, naturally increasing the $\\Mdot$ growth times to $\\approx 10^2-10^6$ yrs.  We explore in \\S 3 the range of initial conditions that will lead to J0806 and J1914. We determine constraints on these initial conditions and present predictions for $\\ddot{\\nu}$ in J0806 as a diagnostic tool.  Finally, we present our discussion in \\S 4. ",
        "conclusions": "} We have shown that the measured $\\nudot$ in J0806 and J1914 are consistent with the \\emph{secular} $\\nudot$ expected for accretion driven systems leading to AM CVn binaries.  As such, the $\\nudot$ measurements \\emph{do not} provide a means of distinguishing between accretion and UI models posited for these systems.  We have shown that accretion models consistent with observations give $\\taumd \\approx 10^2-10^4$ yrs in J0806 and $\\taumd \\approx \\ee{5}{5}$ yrs in J1914. For comparison, \\citet{dallosso06} estimate that the UI timescale for J1914 is $\\approx 10^4 - \\ee{3}{5}$ yrs. We also predicted that accretion models give $\\ddot{\\nu} \\approx \\ee{1-1.5}{-28}$ Hz s$^{-2}$ for J0806; another 5-10 yrs baseline of timing J0806 will be sufficient to test this prediction. Quantitative predictions of $\\ddot{\\nu}$ for the UI model are needed to determine if $\\ddot{\\nu}$ will be a useful diagnostic tool. Our calculations also constrain each system's initial conditions and we relate the current $\\Lacc$ to $\\Porbi$ in both J0806 and J1914.  Recenlty, \\citet{dantona06} proposed that J0806 could be modeled with a degenerate He WD donor with a thick H atmosphere supporting $p-p$ burning. In their model, $\\Porb$ decreases until the H layer is removed, exposing the degenerate He WD.  Our pure He donor models shows that a thick H layer is not necessary to explain a phase where $\\nudot > 0$ during early-contact AM CVn evolution. The combination of sub-photospheric mass transfer and spread in initial donor degeneracy are sufficient to produce either J0806 and J1914.  Indeed, J1914 is naturally accomodated within the AM CVn initial conditions of our model, but not by \\citet{dantona06}, presumably becuase they assume a fully degenerate donor interior. We agree, however, that the presence of H in J0806 or J1914 is not unexpected given the uncertainty in the amount of H left on the donor after the CE phase.  Pure He models are a limiting case that could be altered to include an arbitrarily thick H layer, changing our results quantitatively. As such, a detection of H in J0806 or J1914 does not rule out accretion models nor on its own discriminate between our and the \\citet{dantona06} models.  Future work on these systems must seek to diagnose the source of their X-ray light.  An avenue for this would be consideration of the optical output of the systems where we would expect significant differences between the UI and accretion models. Assessment of potential fine-tuning problems for accretion models in J1914 is also needed.  This would be best performed in a general study addressing the detection probability for early contact AM CVns. Semi-degenerate systems are formed less frequently \\citep{deloye05}, but evolve more slowly at longer $\\Porb$, so the \\textit{a prioi} expectations are unclear.  This would also constrain theoretical expectations for the intial conditions leading to AM CVn binaries. The usefulness of such studies, however, hinges on the determination that J0806 and J1914 are indeed direct-impact accretors, not unipolar inductors.  The community should continue efforts toward distinguishing between these models for these enigmatic systems."
    },
    "0608/astro-ph0608674_arXiv.txt": {
        "abstract": "{} {This paper describes a systematic search for high-order multiplicity among wide visual Pre-Main Sequence (PMS) binaries. } {We conducted an Adaptive Optics survey of a sample of 58 PMS wide binaries from various star-forming regions, which include 52 T Tauri systems with mostly K- and M-type primaries, with the NIR instrument NACO at the VLT. } {Of these 52 systems, 7 are found to be triple (2 new) and 7 quadruple (1 new). The new close companions are most likely physically bound based on their probability of chance projection and, for some of them, on their position on a color-color diagram. The corresponding degree of multiplicity among wide binaries (number of triples and quadruples divided by the number of systems) is 26.9\\,$\\pm$7.2\\% in the projected separation range $\\sim$\\,0\\farcs07-12$^{\\prime\\prime}$, with the largest contribution from the Taurus-Auriga cloud. We also found that this degree of multiplicity is twice in Taurus compared to Ophiuchus and Chamaeleon for which the same number of sources are present in our sample. Considering a restricted sample composed of systems at distance 140-190\\,pc, the degree of multiplicity is 26.8\\,$\\pm$8.1\\%, in the separation range 10/14 AU - 1700/2300 AU (30 binaries, 5 triples, 6 quadruples). The observed frequency agrees with results from previous multiplicity surveys within the uncertainties, although a significant overabundance of quadruple systems compared to triple systems is apparent. Tentatively including the spectroscopic pairs in our restricted sample and comparing the multiplicity fractions to those measured for solar-type main-sequence stars in the solar neighborhood leads to the conclusion that both the ratio of triples to binaries and the ratio of quadruples to triples seems to be in excess among young stars. Most of the current numerical simulations of multiple star formation, and especially smoothed particles hydrodynamics simulations, over-predict the fraction of high-order multiplicity when compared to our results. The circumstellar properties around the individual components of our high-order multiple systems tend to favor mixed systems (i.e. systems including components of wTTS and cTTS type), which is in general agreement with previous studies of disks in binaries, with the exception of Taurus, where we find a preponderance of similar type of components among the multiples studied. } {} ",
        "introduction": "It has long been recognized that the formation of binary or multiple stars is an efficient way to solve one of the most prominent problems in the theory of star formation, namely the \"angular momentum problem\" (Mestel \\& Spitzer\\,\\cite{Mestel_Spitzer1956}, Larson\\,\\cite{Larson_1972}, Mouschovias\\,\\cite{Mouschovias_1977}, Simon et al.\\,\\cite{Simon_etal1995}, Bodenheimer\\,\\cite{Bodenheimer_1995}, Larson\\,\\cite{Larson_2002}). Yet, although we have a general understanding of the formation of single low-mass stars (e.g. Shu et al.\\,\\cite{Shu_etal_1987}), the theory of the origin of binary and multiple low-mass stars is not quite settled (for a review see Tohline\\,\\cite{Tohline_2002}). The favoured mechanism for the formation of binary and multiple stars is the collapse and fragmentation of molecular cloud cores, either fragmentation during dynamical collapse (e.g., Boss\\,\\cite{Boss_1986}) or from a quasistatic situation (i.e. disk fragmentation e.g., Bonnell\\,\\cite{Bonnell_1994}) or filament fragmentation (e.g. Zinnecker\\,\\cite{Zinnecker_1991}, Bonnell \\& Bastien\\,\\cite{Bonnell_Bastien1991}). Prompt fragmentation following the collision of two Jeans masses has also been suggested (Pringle\\,\\cite{Pringle_1991}). In addition, further processes likely to occur during the formation of binary and multiple stars like circumbinary disk accretion, interaction with circumstellar disks and dynamical interactions are nowadays being investigated in detail through numerical simulations (Bate et al.\\,\\cite{Bate_etal2003}, Sterzik \\& Durisen\\,\\cite{Sterzik_Durisen2003}, \\cite{Delgado_Donate_etal2004}, Goodwin et al.\\,\\cite{Goodwin_etal2004}). The scope of these studies is mainly to obtain statistical properties of the stellar systems resulting from these evolutionary calculations for direct comparison with  observations; such as, relative frequency of single, binary and multiple stars, the properties of multiple stars and brown dwarfs, and the initial mass function.  Several multiplicity surveys among T Tauri stars using high-angular resolution techniques have been carried out over the last decade, with the main outcome that binary stars are much more common in several of the young star-forming T-associations compared to the main-sequence stars in the solar-neighborhood (Leinert et al.\\,\\cite{Leinert_etal1993}, Ghez et al.\\,\\cite{Ghez_etal1993}, Ghez et al.\\,\\cite{Ghez_etal1997a}, Duch\\^{e}ne\\,\\cite{Duchene_1999}). In addition, it seems that a comparable or even higher frequency of binary and multiple stars is found at earlier evolutionary stages (e.g. Reipurth 2000, Duch\\^{e}ne et al.\\,\\cite{Duchene_etal_ppv} and references therein). Paradoxically, although studies of high-order multiple star (N$>$2) systems might provide important constraints about the process of multiple star formation (Batten\\,\\cite{Batten_1973}, Fekel\\,\\cite{Fekel_1981}, Sterzik \\& Tokovinin\\,\\cite{Sterzik_Tokovinin_2002}, Tokovinin \\& Smekhov\\,\\cite{Tokovinin_Smekhov_2002}, see also Tokovinin\\,\\cite{Tokovinin_2001}), relatively little has been done observationally in a systematic way regarding higher-order multiplicity among Pre-Main Sequence (PMS) stars. An exception is Koresko (\\cite{Koresko_2002}) who observed the close environment of 14 PMS binaries in southern star-forming regions with high-angular resolution techniques and found that half of them are in fact hierarchical triples or present resolved circumstellar structures consistent with marginally resolved companions. In the last years, an increasing number of high-angular resolution observations reports the discovery of new companions in known multiple T Tauri systems. Examples are T Tau itself (Koresko\\,\\cite{Koresko_2000}, K\\\"{o}hler et al.\\,\\cite{Koehler_etal2000a}), VW Cha (Brandeker et al. \\cite{Brandeker_etal2001}), V773 Tau (Duch\\^{e}ne et al.\\,\\cite{Duchene_etal2002}), AS 353 (Tokunaga et al.\\,\\cite{Tokunaga_etal2004}). Some of these systems present close to edge-on circumstellar disks, more easily detectable in multiple systems and with current imaging techniques like AO. The only known cases of such disks are HK Tau B (Stapelfeldt et al.\\,\\cite{Stapelfeldt_etal1998}, Koresko\\,\\cite{Koresko_1998}), HV Tau C (Monin \\& Bouvier\\,\\cite{Monin_Bouvier_2000}, Stapelfeldt et al.\\,\\cite{Stapelfeld_etal2003}), and LkH$\\alpha$ 263\\,C (Chauvin et al.\\,\\cite{Chauvin_etal2002}, Jayawardhana et al.\\,\\cite{Jayawardhana_etal2002}). Also, it has been noted that approximately half of the known T Tauri spectroscopic binaries are in fact hierarchical triples with an outlier companion (Melo\\,\\cite{Melo_2003}, Sterzik et al.\\,\\cite{Sterzik_etal2005}). All the above suggests that with the current observational techniques we are able to reveal a large enough number of high-order multiple PMS stars to allow for meaningful statistical analysis of their occurrence and properties.  We report here about a high-angular resolution survey which focuses on the N$>$2 multiplicity among T Tauri stars. We searched in a sample of 58 wide PMS binaries extracted from the list of Reipurth \\& Zinnecker (\\cite{Reipurth_Zinnecker1993}, hereafter RZ93) to discover additional close companions using the Adaptive Optics (AO) near-infrared instrument NAOS/CONICA (NACO, Rousset et al.\\,\\cite{Rousset_etal2002}) attached to the telescope UT4/Yepun of the Very Large Telescope (VLT). The rationale behind this effort was the following\\,: if one considers a wide 2\\,arcsec PMS binary (corresponding to $\\sim$\\,300\\,AU at 150\\,pc, the distance of the most prominent star forming regions), then a close ($\\sim$\\,0.1-0.2\\,arcsec) companion could still exist without violating the hierarchical triple system stability criterion (Harrington\\,\\cite{Harrington_1975}, Eggleton  \\& Kiseleva\\,\\cite{Eggleton_Kiseleva_1995}, Mardling \\& Aarseth\\,\\cite{Mardling_Aarseth_2001}).  This paper is organized as follows\\,: in Sect.\\,\\ref{sect:sample} we present the sample of wide PMS binaries and the AO near-infrared observations obtained at VLT with NACO in Sect.\\,\\ref{sect:obs_data_red}. Candidate visual companions to these binaries detected in the images are presented in Sect.\\,\\ref{sect:results} and physical companions are identified on both colors and statistical grounds. This allows us, in Sect.\\,\\ref{sect:multiplicity_stat}, to derive the frequency of triples/quadruples systems among wide young binaries and to compare this quantity to previous surveys (Sect.\\,\\ref{sect:multiplicity_stat_compa_surveys}), to recent results of numerical simulations of multiple star formation (Sect.\\,\\ref{sect:multiplicity_stat_compa_theory}), as well as to what is found among main-sequence multiples (Sect.\\,\\ref{sect:MS_mult} and \\ref{sect:total_mult}). We also discuss the properties of these triples/quadruples systems in terms of dynamical stability (Sect.\\,\\ref{sect:syst_stability}) and relative disk evolutions (Sect.\\,\\ref{sect:disk_evolution}). Finally, we summarize our main results in Sect.\\,\\ref{sect:summary}.   \\begin{table*} \\scriptsize \\caption{Observed sample of wide PMS binaries.} \\begin{center} \\renewcommand{\\arraystretch}{0.5} \\setlength\\tabcolsep{5pt} \\begin{tabular}{l@{\\hspace{1mm}}c@{\\hspace{1mm}}l@{\\hspace{2mm}}l@{\\hspace{1mm}}r@{}l@{\\hspace{3mm}}l@{\\hspace{1mm}}r@{\\hspace{2mm}}r@{\\hspace{2mm}}r @{\\hspace{4mm}}r@{\\hspace{2mm}}r@{\\hspace{1mm}}c @{\\hspace{0mm}}c@{\\hspace{2mm}}l@{\\hspace{2mm}}l@{\\hspace{0mm}}c} \\hline\\noalign{\\smallskip}  Name & HBC & Other & \\multicolumn{1}{c}{R.A.} & \\multicolumn{2}{c}{Decl.} &  Cloud & D & V & K & \\multicolumn{2}{r}{RZ93$^c$} & add. & SB & SpT1 & SpT2 & Ref  \\\\  &  & name & \\multicolumn{3}{c}{[J2000.0]} &   & [pc] & &  & $\\rho$ & PA & comp.$^b$ & & & & (SpT) \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\object{LkH$\\alpha$ 262/263} ..      &  8+9 & \t\t& 02 56 08.4 & +   & 20 03 40  &  MBM 12  \t&    275   & 14.6  &  9.5  & 4.1 & 57 & X &  & M0 & M0+M4 & 1 \\\\ \\noalign{\\smallskip} \\object{J 4872} ...............      \t\t  & \t& \t\t& 04 25 17.6 & +   & 26 17 51  &  Tau-Aur  \t&    142   & 13.0  &  8.6  & 3.5 & 231 & X &  & K9 & M1 & 2 \\\\ \\noalign{\\smallskip} \\object{FV Tau} .............        \t\t  & 386+387 & \t\t& 04 26 53.6 & +   & 26 06 55  &  Tau-Aur  &    142   & 15.4  &  7.4   & 0.7 & 289 & X &  & K5+K6 & M3+M5 &  1 \\\\ \\noalign{\\smallskip} \\object{UX Tau} .............        \t\t  & 43+42 & \t\t& 04 30 04.0 & +   & 18 13 49  &  Tau-Aur  \t&    142   & 10.7  &  7.6    & 2.7 & 270 & X &  & K4 & M2+M3 &  1 \\\\ \\noalign{\\smallskip} \\object{DK Tau} .............        \t\t  & 45 & \t\t& 04 30 44.3 & +   & 26 01 25  &  Tau-Aur &    142   & 12.6  &  7.1    & 2.4 & 119 & &  & K7/K9 & K7/M1 & 3,2 \\\\ \\noalign{\\smallskip} \\object{HK Tau} .............        \t\t  & 48 & \t\t& 04 31 50.6 & +   & 24 24 18  &  Tau-Aur  &    142   & 15.0  &  8.6    & 2.4 & 172 & &  & M1 & M2 & 4, 2  \\\\ \\noalign{\\smallskip} \\object{LkH$\\alpha$ 266} .........   \t  & 51+395 & V710 Tau & 04 31 57.8 & +   & 18 21 37  &  Tau-Aur  \t&    142   & 14.6  &  8.5    & 3.2 & 177 & &  & M0.5 & M2 & 3  \\\\ \\noalign{\\smallskip} \\object{GG Tau} .............        \t\t  & 54 & \t\t& 04 32 30.3 & +   & 17 31 41  &  Tau-Aur  &    142   & 12.3  &  7.4    & 1.4 & 185 & X &  & K7+M0.5 & M5.5+M7.5 &  1 \\\\ \\noalign{\\smallskip} \\object{UZ Tau} .............        \t\t  & 52+53\t& \t\t& 04 32 43.0 & +   & 25 52 32  &  Tau-Aur  \t&    142   & 12.9  &  7.4    & 3.6 & 275 & X & 1 & M1 & M2+M2 & 1 \\\\ \\noalign{\\smallskip} \\object{HN Tau} .............        \t\t  & 60+406 & \t\t& 04 33 39.3 & +   & 17 51 53  &  Tau-Aur  &    142   & 13.7  &  8.4    & 3.2 & 217 & &  & K5 & M4 & 3 \\\\ \\noalign{\\smallskip} \\object{IT Tau} ...............      \t\t  & \t&  Haro 6-26 & 04 34 13.9 & +   & 26 11 42  &  Tau-Aur  \t&    142   & 14.9  & 7.9    & 2.4 & 224 & &  &  K3 & M4 & 2 \\\\ \\noalign{\\smallskip} \\object{L1642-1} ............        \t\t  & 413 & \tEW Eri & 04 35 02.3 &$-$& 14 13 41  &  L1642  &    140   & 13.7  &  7.7    & 2.7 & 349 & &  &  K7 &  & 5 \\\\ \\noalign{\\smallskip} \\object{RW Aur} ............         \t\t  & 80+81 & \t\t& 05 07 49.6 & +   & 30 24 05  &  Tau-Aur  \t&    142   & 10.3  &  7.0    & 1.4 & 254 & & 2 & K1: & K5: & 4, 1  \\\\ \\noalign{\\smallskip} \\object{CO Ori} ..............       \t\t  & 84 & \t\t& 05 27 38.3 & +   & 11 25 39  &  Orion &    460   & 10.6  &  6.5    & 2.0 & 280 & &  & F8:e & & 4  \\\\ \\noalign{\\smallskip} \\object{AR Ori} .............        \t\t  & 155 & \t\t& 05 35 54.1 &$-$& 05 04 14  &  Orion  &    460   & 13.9  &  9.8    & 2.0 & 249 & &  & K7/M0 & & 4  \\\\ \\noalign{\\smallskip} \\object{LkH$\\alpha$ 336} ........    \t  & 190+516 & \t\t& 05 54 20.1 & +   & 01 42 56  &  L1622 &    460   & 14.4  &  9.2    & 5.8 & 95 & &  & K7 & M0.5 &  1 \\\\ \\noalign{\\smallskip} \\object{CGH$\\alpha$ 5/6} ........    \t  & \t& \t\t& 07 31 37.4 &$-$& 47 00 22  &  Gum Neb. \t&    450   & 14.2  &  9.1    & 11.1 & 317 & &  & K7 & K2/K5 & 6  \\\\ \\noalign{\\smallskip} \\object{PH$\\alpha$ 14} ............  \t  & 554 & \t\t& 08 08 33.8 &$-$& 36 08 10  &  Gum Neb. \t&    450   & 15.8  &10.3    & 0.6 & 164 & &  & M2: & & 4,7  \\\\ \\noalign{\\smallskip} \\object{PH$\\alpha$ 30} ............  \t  & \t& SPH 42\t& 08 12 05.6 &$-$& 35 31 45  &  Gum Neb. \t&    450   & 15.1  &12.2    & 0.7 & 345 & &  & B2e & & 8 \\\\ \\noalign{\\smallskip} \\object{vBH 16} .............        \t\t  & \t& vdBH 16 & 08 27 39.0 &$-$& 51 09 50  &  Gum Neb. \t&    450   & 15.6  &  9.3    & 4.5 & 330 & &  &  &  \\\\ \\noalign{\\smallskip} \\object{PH$\\alpha$ 51} ............. \t  & 561& \t\t& 08 15 55.3 &$-$& 35 57 58  &  Gum Neb. \t&    450   & 15.9  &11.1    & 0.6 & 185 & &  & K7/M0 &  & 8 \\\\ \\noalign{\\smallskip} \\object{HD 76534} .........          \t  & \t& He 3-225 &  08 55 08.7 &$-$& 43 28 00  & Vela R2$^d$ &   830$^d$    &   8.0  &  7.8    & 2.1 & 303 & &  & B2 & B3 & 14 \\\\ \\noalign{\\smallskip} \\object{SX Cha} .............        \t\t  & 564 & T3\t\t& 10 55 59.9 &$-$& 77 24 41  &   Cha I  \t\t&    160   & 14.7  &  8.7    & 2.2 & 288 & &  &M0.5 & & 4  \\\\ \\noalign{\\smallskip} \\object{Sz 15} .................     \t\t  & \t&  T19 & 11 05 41.5 &$-$& 77 54 44  &   Cha I  \t\t&    160   &   13.2     &10.6    & 10.6 & 28 & &  & K5/K6 & & 7 \\\\ \\noalign{\\smallskip} \\object{ESO H$\\alpha$ 281} .....      & \t& \t\t& 11 07 04.0 &$-$& 76 31 45  &   Cha I  \t\t&    160   &   13.7$^a$     &  9.7   & 1.7 & 167 & &  &  M0.5 & M4.5 & 10 \\\\ \\noalign{\\smallskip} \\object{Sz 19} ................      \t\t  & 245 & DI Cha & 11 07 20.7 &$-$& 77 38 07  &   Cha I &    160   & 10.9  &  6.2    & 4.6 & 203 & &  & G2 & & 4  \\\\ \\noalign{\\smallskip} \\object{VV Cha} ............         \t\t  & 573 & \tT27\t& 11 07 28.4 &$-$& 76 52 12  &   Cha I &    160   & 14.8  &  9.5    & 0.8 & 14 & &  & M1.5 & M3 & 10 \\\\ \\noalign{\\smallskip} \\object{VW Cha} ............         \t\t  &  575 & \tT30/T31\t& 11 08 01.8 &$-$& 77 42 29  &   Cha I  &    160   & 12.6  &  7.0    & 16.7 & 221 & X & 3 &  K5/K7 & K7 &  1 \\\\ \\noalign{\\smallskip} \\object{Glass I} ...............     \t\t  & \t& Ced 111 IRS 4\t & 11 08 15.4 &$-$& 77 33 54  &   Cha I  \t\t&    160   & 13.3  &  6.9    & 2.4 & 105 & &  &  K3 & G5e & 11 \\\\ \\noalign{\\smallskip} \\object{CoD $-$29$\\degr$8887} ... & \t& TWA 2 & 11 09 14.0 &$-$& 30 01 39  &   TW Hya \t&     50    & 11.1  &  6.7    & $<$0.8 & 40 & &  & M2e & M2 & 12 \\\\ \\noalign{\\smallskip} \\object{Sz 30} ..................    \t\t  & \t& T39 & 11 09 12.3 &$-$& 77 29 12  &   Cha I  \t\t&    160   & 13.2  &  9.0    & 1.1 & 22 & &  & M0.5 & M2 & 10 \\\\ \\noalign{\\smallskip} \\object{Hen 3-600} ..........        \t  & \t&  TWA 3\t& 11 10 28.9 &$-$& 37 32 05  &   TW Hya \t&     50    & 12.1  &  6.8    & 1.4 & 230 & & 4 & M3e & M3.5 & 12 \\\\ \\noalign{\\smallskip} \\object{Sz 41} .................     \t\t  & 588 & T51 & 11 12 24.5 &$-$& 76 37 06  &   Cha I  &    160   & 11.6  & 8.0    & 1.9 & 164 & & 5 & K3.5  &  & 9  \\\\ \\noalign{\\smallskip} \\object{CV Cha} .............        \t\t  & 247+589 & \tT52\t& 11 12 27.8 &$-$& 76 44 22  &   Cha I  &    160   & 11.0  & 6.9    & 11.4 & 105 & & 6 & G8 & G: & 4  \\\\ \\noalign{\\smallskip} \\object{Sz 48} ..................\t  \t  & \t& \t\t& 13 00 53.2 &$-$& 77 09 10\t&    Cha II\t\t&    178   & 17.4 & 9.5     & 1.5 & 232 & &  & K7/M0 & M0 & 10   \\\\ \\noalign{\\smallskip} \\object{BK Cha} ..............\t\t  & \t& \t\t& 13 07 09.3 &$-$& 77 30 24\t&    Cha II\t\t&    178   & 15.2   & 8.4     & 0.8 & 329 & &  & K5/K7 & M0.5 & 10 \\\\ \\noalign{\\smallskip} \\object{Sz 60} ..................\t\t  & \t& \t\t& 13 07 23.4 &$-$& 77 37 23\t&    Cha II\t\t&    178   &  17.1 &  9.5    & 3.4 & 280 & &  & M1 & M4 & 13 \\\\ \\noalign{\\smallskip} \\object{Sz 62} ..................     \t\t  & \t& \t\t& 13 09 50.7 &$-$& 77 57 24  &   Cha II \t\t&    178   & 15.6  & 9.1    & 1.1 & 261 & &  & M2 & M3.5 & 10  \\\\ \\noalign{\\smallskip} \\object{Herschel 4636} ....          \t  & \t& Hen 3-949 & 13 57 44.1 &$-$& 39 58 45  & NGC 5367 \t&    630   &   9.7  & 7.2     & 3.7 & 214 & &  & B3 & & 14 \\\\ \\noalign{\\smallskip} \\object{ESO H$\\alpha$ 283} ......     & \t& \t\t& 15 00 29.6 &$-$& 63 09 46  & Circinus \t&    700   &  15.5  &10.3   & 2.1 & 242 & &  &  &  \\\\ \\noalign{\\smallskip} \\object{Sz 65} .................. \t\t  & 597 & \t\t& 15 39 27.7 &$-$& 34 46 17\t&    Lup I &    190   & 12.7  &  8.0    & 6.4 & 98 & &  & M0 & & 15  \\\\ \\noalign{\\smallskip} \\object{Sz 68} ..................\t\t  & 248 & \tHT Lup & 15 45 12.9 &$-$& 34 17 31\t&    Lup I &    190   & 10.4  &  6.5    & 2.6 & 295 & X &  & K2 & M6 &  1  \\\\ \\noalign{\\smallskip} \\object{HO Lup} ..............\t\t  & 612 & \t\t& 16 07 00.6 &$-$& 39 02 19\t&    Lup III &    190   & 13.0  &  8.6    & 1.5 & 36 & &  & K7/M0 & M2 & 10  \\\\ \\noalign{\\smallskip} \\object{Sz 101} ................\t\t  & \t& \t\t& 16 08 28.4 &$-$& 39 05 32\t&    Lup III\t&    190   & 15.5  &  9.4    & 0.8 & 305 & &  & M2.5 & M3.5 & 10  \\\\ \\noalign{\\smallskip} \\object{Sz 108} ................\t\t  & 620 & \t\t& 16 08 42.7 &$-$& 39 06 18\t&    Lup III &    190   & 13.1  &  8.8    & 4.2 & 25 & &  & M1 & & 15  \\\\ \\noalign{\\smallskip} \\object{Sz 120} ................\t\t  & \t    & HD 144965 & 16 10 10.6 &$-$& 40 07 44  &    Lup III\t&    190   &   7.1\t  &  6.2    & 2.6 & 142 & &  & B5  & & 15  \\\\ \\noalign{\\smallskip} \\object{WSB 3} ...............\t\t  & \t& \t\t& 16 18 49.5 &$-$& 26 32 53\t&    Oph\t&    160   & 15.0 &  9.3    & 0.6 & 162 & &  & M3  & & 16  \\\\ \\noalign{\\smallskip} \\object{WSB 11} .............\t\t  & \t& \t\t& 16 21 57.3 &$-$& 22 38 16\t&    Oph\t&    160   & 18.5$^a$  &10.1   & 0.5 & ... & &  &  &  \\\\ \\noalign{\\smallskip} \\object{WSB 20} .............\t\t  & 257 & \tHaro 1-4\t& 16 25 10.5 &$-$& 23 19 14\t&    Oph\t&    160   & 13.4  &  7.5    &$ <$1.0 & 23 & &  & K6/K7 & & 4  \\\\ \\noalign{\\smallskip} \\object{WSB 28} .............\t\t  & \t& ISO-Oph 27 & 16 26 20.7 &$-$& 24 08 48\t&    Oph\t&    160   &  14.7$^a$ &  9.5    & 5.1 & 358 & &  & M3 & M7 & 17 \\\\ \\noalign{\\smallskip} \\object{SR 24} .................\t \t  & 262 & \tWSB 42    & 16 26 58.8 &$-$& 24 45 37\t&    Oph \t&    160   &  15.9  &  7.1    & 5.2 & 348 & X &  & K2: & M0.5 &   1  \\\\ \\noalign{\\smallskip} \\object{Elias 2-30} ..........\t\t  & \t&  SR 21\t\t& 16 27 10.2 &$-$& 24 19 16\t&    Oph\t&    160   & 14.1  &  6.7    & 6.4 & 175 & &  & G2.5 & M4 & 17  \\\\ \\noalign{\\smallskip} \\object{WSB 46} .............\t\t  & 641+640 & \tROXs 20\t& 16 27 15.1 &$-$& 24 51 39\t&    Oph\t&    160   &   14.9 &  9.4    & 10.3 & 301& &  & M0/M3  & M1/M4 &18 \\\\ \\noalign{\\smallskip} \\object{Haro 1-14c} .........\t  \t  & 644+267 & \t\t& 16 31 04.4 &$-$& 24 04 32\t&    Oph\t&    160   & 12.3  &  7.8    & 12.9 & 122 & & 7 & K3 & M0 & 4 \\\\ \\noalign{\\smallskip} \\object{ROXs 43} .............\t\t  & \t& NTTS 162819-2423 & 16 31 20.1 &$-$& 24 30 05\t&    Oph\t&    160   & 10.6  &  6.7    & 4.3 & 13 &  X & 8 & G0 & K3 &  1 \\\\ \\noalign{\\smallskip} \\object{Elias 2-49} ...........\t\t  & \t& HD 150193 & 16 40 17.9 &$-$& 23 53 45  &    Oph\t&    160   &   8.9\t  &  5.5    & 1.1 & 227 & & & A1/A3Ve &  &  19 \\\\ \\noalign{\\smallskip} \\object{HBC 652} ............\t\t  & 652 & V2507 Oph & 16 48 18.0 &$-$& 14 11 15\t &    L162\t\t&    160   & 13.5  &  7.5    & 8.7 & 322 & &  & K4/K5 & & 4  \\\\ \\noalign{\\smallskip} \\object{LkH$\\alpha$ 346} ........... & 275 & \tTH$\\alpha$ 27-4 & 17 11 03.9 &$-$& 27 22 57\t &    B59\t\t&    160   & 14.3 &   8.1   & 5.0 & 329 & &  & K7 & M5 &  1  \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} \\label{Tab:sample} \\begin{minipage}[position]{18cm} Note\\,: HBC = catalog entry number in Herbig \\& Bell\\,\\cite{Herbig_Bell_1988}.\\\\ $^a$\\,: R magnitude. \\\\ $^b$\\,: Additional visual companion(s) discovered since RZ93 (see references in Table\\,\\ref{Tab:systems_parameters} or in the text). \\\\ $^c$\\,: Astrometry of the wide binary as quoted in RZ93 with $\\rho$ in arcsec and PA in degree (only one entry for the possible triples quoted in RZ93). \\\\ $^d$\\,: see Herbst\\,\\cite{Herbst_1975} and van den Ancker et al.\\,\\cite{Van_den_Ancker_etal1998}. \\\\ Spectral type references\\,: (1) see Table\\,\\ref{Tab:systems_accretion}. (2) Duch\\^{e}ne\\,\\cite{Duchene_etal1999}. (3) Hartigan et al.\\,\\cite{Hartigan_etal1994}. (4) Herbig \\& Bell\\,\\cite{Herbig_Bell_1988}. (5) Sandell et al.\\,\\cite{Sandell_etal_1987}. (6) Reipurth \\& Pettersson\\,\\cite{Reipurth_Pettersson1993}. (7) Kim et al.\\,\\cite{Kim_etal2005}. (8) Pettersson\\,\\cite{Pettersson_1987}. (9) Luhman\\,\\cite{Luhman_2004}. (10) Brandner \\& Zinnecker\\,\\cite{Brandner_Zinnecker_1997}. (11) Feigelson \\& Kriss\\,\\cite{Feigelson_Kriss_1989}. (12) Torres et al.\\,\\cite{Torres_etal2003}. (13) Hughes \\& Hartigan\\,\\cite{Hughes_Hartigan_1992}. (14) Viera et al.\\,\\cite{Viera_etal2003}. (15) Hughes et al.\\,\\cite{Hughes_etal1994}. (16) Meyer et al.\\,\\cite{Meyer_etal1993}. (17) Prato et al.\\,\\cite{Prato_etal2003}. (18) Luhman \\& Rieke\\,\\cite{Luhman_Rieke_1999}. (19) Gray \\& Corbally\\,\\cite{Gray_Corbally_1998}. \\\\ Distance references\\,: Chamaeleon (I\\,: 160\\,pc $\\pm$ 15\\,pc, II\\,: 178\\,pc $\\pm$ 18\\,pc, Whittet et al.\\,\\cite{Whittet_etal1997}), Gum Nebula (450\\,pc, Kim et al.\\,\\cite{Kim_etal2005}), Lupus (190\\,pc $\\pm$ 27\\,pc, Wichmann et al.\\,\\cite{Wichmann_etal1998}), Ophiuchus (160\\,pc, Whittet\\,\\cite{Whittet_1974}), Taurus-Auriga (142\\,pc $\\pm$ 14\\,pc, Wichmann et al.\\,\\cite{Wichmann_etal1998}), TW Hya (50\\,pc $\\pm$ 12\\,pc, Mamajek\\,\\cite{Mamajek_2005}), MBM12 (275\\,pc, Luhman\\,\\cite{Luhman_2001}), L1642, aka MBM20 (112-160\\,pc, Hearty et al.\\cite{Hearty_etal2000}, 140\\,pc adopted), NGC\\,5367, Circinus, L162 and B59 (values reported in RZ93 adopted). \\\\ SB notes\\,: (1) UZ Tau A\\,: 19.1-days single-lined SB (SB1) (Mathieu et al.\\,\\cite{Mathieu_etal1996}), converted to double-lined SB (SB2) by Prato et al.\\,\\cite{Prato_etal2002}. (2) RW Aur A\\,: suspected SB1 with P=2.77-days (Gahm et al.\\,\\cite{Gahm_etal1999}, Petrov et al.\\,\\cite{Petrov_etal2001}, see also Alencar et al.\\,\\cite{Alencar_etal2005}). (3) VW Cha A\\,: possible SB2 (Melo\\,\\cite{Melo_2003}). (4) TWA 3A\\,: suspected SB2 (Torres et al.\\,\\cite{Torres_etal2003}). (5) Sz 41 A\\,: probable SB (Reipurth et al.\\,\\cite{Reipurth_etal2002}) with period of about 125 days (best candidate out of five radial-velocity variables). (6) CV Cha\\,: possibly SB, variable radial-velocity (Reipurth et al.\\,\\cite{Reipurth_etal2002}). (7) Haro 1-14c\\,: 591-days SB1 (Reipurth et al.\\,\\cite{Reipurth_etal2002}), converted to SB2 by Simon \\& Prato\\,\\cite{Simon_Prato_2004}. (8) ROXs 43 A\\,: 89-days SB1 (Mathieu et al.\\,\\cite{Mathieu_etal1989}). \\end{minipage} \\end{table*}   \\begin{table*} \\scriptsize \\caption{Observation Log.} \\begin{center} \\renewcommand{\\arraystretch}{0.5} \\setlength\\tabcolsep{5pt} \\begin{tabular}{l@{\\hspace{2mm}}l@{\\hspace{2mm}}r@{\\hspace{4mm}}r@{$\\times$}l@{\\hspace{1mm}}r@{$\\times$}l@{\\hspace{1mm}}r@{$\\times$}lr@{\\hspace{1mm}}r@{\\hspace{2mm}}c@{\\hspace{2mm}}r@{\\hspace{1mm}}r@{\\hspace{1mm}}r@{\\hspace{4mm}}c@{\\hspace{2mm}}cc@{\\hspace{1mm}}r@{\\hspace{1mm}}r} \\hline\\noalign{\\smallskip} Name & Obs. date & Cam &  \\multicolumn{6}{c}{Int. time [s]} && WFS-Dic. & seeing$^a$ & \\multicolumn{3}{c}{FWHM [mas]}  && Airmass  & \\multicolumn{3}{c}{Strehl ratio [\\%]} \\\\ \\cline{4-9}  \\cline{13-15}  \\cline{18-20}\\\\ &  &  & \\multicolumn{2}{c}{[FeII]} & \\multicolumn{2}{c}{H$_2$} & \\multicolumn{2}{c}{Br$\\gamma$} && & [arcsec]  & [FeII]  & H$_2$ & Br$\\gamma$ && & [FeII]  & H$_2$ & Br$\\gamma$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} LkH$\\alpha$ 262/263 ..      & 2002 Nov 14 & S27 & 15 & 2.7 & 15 & 1.5 & 15 & 1.4 && VIS-VIS  & 1.0 &  129 &  89    &   88   && 1.43 & 4 & 14 & 12 \\\\ \\noalign{\\smallskip} J 4872 ...............      \t\t  & 2002 Nov 13 & S13  \t& 30 & 1.1 & 30 & 1.6 & 30 & 1.6 && VIS-VIS  & 0.8 &   72   &  73   &   72   && 1.62 & 18 & 42 & 40  \\\\ \\noalign{\\smallskip} FV Tau$^b$ .............             & 2002 Nov 13 & S27  \t& 9 & 47 & 9 & 19 & 9 & 19 && IR-N20C80 & 0.8 &  121  &  89   &  81   && 1.59 &  7  & 19 & 13  \\\\ \\noalign{\\smallskip} UX Tau .............        \t\t  & 2002 Oct 22 & S13  \t& 60 & 0.6 & 60 & 0.6 & 60 & 0.6 && VIS-VIS & 1.1 &    60  &  68   &  68   && 1.40 &  22  & 38  &  47  \\\\ \\noalign{\\smallskip} DK Tau .............        \t\t  & 2002 Nov 13 & S13  \t& 30 & 1.4 & 30 &1.0 & 30 & 0.9 && IR-N20C80 & 0.9 &   64   &  70   &  68    && 1.59 & 15   & 29  & 32  \\\\ \\noalign{\\smallskip} HK Tau .............        \t\t  & 2003 Feb 17 & S13  \t& 12 & 5.5 & 12 & 2.9 & 12 & 2.9 && IR-N20C80 & 1.0 &   80   &  73   &  73    && 1.75 & 7   & 21  & 26 \\\\ \\noalign{\\smallskip} LkH$\\alpha$ 266 .........   \t  & 2002 Oct 22 & S27  \t& 105 & 0.3 & 105 &\t0.3 & 105 & 0.3 && VIS-VIS & 1.7 &   90   &  82   &  93    && 1.44 & 5   & 12  & 8 \\\\ \\noalign{\\smallskip} GG Tau .............        \t\t  & 2002 Oct 22 & S27  \t& 60 & 1.0 & 30 & 1.0 & 30 & 1.0 && VIS-VIS & 1.7 &   67   &  70   &  87   && 1.49 & 12   &  31 & 19     \\\\ \\noalign{\\smallskip} UZ Tau .............        \t\t  & 2002 Nov 14 & S13  \t& 45 & 0.6 & 45 & 0.5 & 45 & 0.5 && VIS-VIS & 1.0 &   84   &  77   &  77  \t&& 1.68 & 8  & 17  &  27       \\\\ \\noalign{\\smallskip} HN Tau .............        \t\t  & 2003 Feb 18 & S13 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}& 15 & 1.5 && VIS-VIS & 0.6 &   ...     & ...      &  72   && 1.59 & ...   & ...  & 29 \\\\ \\noalign{\\smallskip} IT Tau ...............      \t\t  & 2003 Feb 19 & S13  \t& 15 & 3.8 & 15 & 1.7 & 15 & 1.7 && IR-N20C80 & 1.1 &  104  & 77    &  75   && 1.84 & 2   & 17  & 21 \\\\ \\noalign{\\smallskip} L1642-1 ............        \t\t  & 2003 Feb 20 & S13  \t& 60 & 0.9 & 60 & 0.6 & 60 & 0.6 && VIS-VIS & 1.1 &    78  &  72   & 77  \t&& 1.50 & 7   & 23  & 22  \\\\ \\noalign{\\smallskip} RW Aur ............         \t\t  & 2002 Nov 17 & S13  \t& 45 & 0.7 & 45 & 0.5 & 45 & 0.5 && IR-N20C80 & 0.8 &    66  &  70   & 70   \t&& 1.79 & 15   & 34  & 36 \\\\ \\noalign{\\smallskip} CO Ori ..............       \t\t  & 2002 Nov 13 & S13  \t& 30 & 2.0 & 30 & 1.0 & 30 & 2.0 && VIS-VIS & 1.0 &    68  &  71    & 73   && 1.53 & 12   & 23  & 28 \\\\ \\noalign{\\smallskip} AR Ori .............        \t\t  & 2003 Feb 19 & S13  \t& 3 & 20 & 3 & 10 & 3 & 10 && VIS-VIS & 1.1 &   151& 122   & 91    && 1.21 & 3   & 8  & 11 \\\\ \\noalign{\\smallskip} LkH$\\alpha$ 336 ........    \t  & 2003 Feb 19 & S27  \t& 6 & 10 & 6 & 9.0 & 6 & 8.5 && VIS-VIS & 1.4 &     89 &   87   & 90  \t&& 1.20 & 6   & 12  & 12 \\\\ \\noalign{\\smallskip} CGH$\\alpha$ 5/6 ........    \t  & 2002 Dec 23 &S27  \t& 36 & 1.3 & 36 & 1.0 & 36 & 1.0 && VIS-VIS & 0.7 &     85 &   84   & 79  \t&& 1.08 &  6  & 13  & 21 \\\\ \\noalign{\\smallskip} PH$\\alpha$ 14 ............  \t  & 2002 Dec 23 &S27  \t& 15 & 4.0 & 15 & 2.0 & 15 & 2.0 && VIS-VIS & 0.7 &     86 &   80   & 78    && 1.02 & 8   & 16  & 24 \\\\ \\noalign{\\smallskip} PH$\\alpha$ 30 ............  \t  & 2003 Jan 17 &S27  \t& 3 & 28 & 3 & 21 & 3 & 20 && VIS-VIS &  ... &   104 &   84   & 83  \t&& 1.02 &  9  &  27 & 32      \\\\ \\noalign{\\smallskip} vBH 16 .............        \t\t  & 2003 Jan 27 &S13  \t& 15 & 5.5 & 15 & 4.5 & 15 & 4.5 && IR-N20C80 & 0.9 &   122 &  112  & 97   && 1.15 & 4   & 10  & 13 \\\\ \\noalign{\\smallskip} PH$\\alpha$ 51 ............. \t  & 2003 Jan 23 &S13  \t& 3 & 85 & 3 & 31 & 3 & 30 && VIS-VIS & 1.0 &   135 &  107  & 88    && 1.02 & 3   & 7  & 11 \\\\ \\noalign{\\smallskip} HD 76534 .........          \t  & 2003 Jan 20 &S13  \t& 60 & 0.6 & 60 & 0.6 & 60 & 0.6 && VIS-VIS & 1.6 &     69 &    80   & 72    && 1.18 & 12   & 18  & 26 \\\\ \\noalign{\\smallskip} SX Cha .............        \t\t  & 2003 Jan 22 &S13  \t&12 & 8.5 & 12 & 2.5 & 12 & 2.5 && IR-N20C80 & 0.9 &     93 &    82   & 76    && 1.66 & 7   & 16  & 21 \\\\ \\noalign{\\smallskip} Sz 15 .................     \t\t  & 2003 Jan 22 &S13  \t& 12 & 3.9 & 12 & 4.0 & 12 & 3.9 && VIS-VIS & 0.7 &     85 &    85   & 83    && 1.67 & 7   & 17  & 16 \\\\ \\noalign{\\smallskip} ESO H$\\alpha$ 281 .....      & 2003 Jan 22\t &S13  \t& 3 & 51 & 3 & 34 & 3 & 33 && VIS-VIS & 0.8 &     99 &    83   & 87    && 1.62 & 5   & 16  & 16\t \\\\ \\noalign{\\smallskip} Sz 19$^b$ ................      \t  & 2003 Jan 20 &S13  \t& 6 & 22 & 6 & 16 & 6 & 16 && VIS-VIS & 1.5 &     69 &    76   & 76    && 1.66 & 12   & 22  & 23 \\\\ \\noalign{\\smallskip} VV Cha ............         \t\t  & 2003 Feb 19 &S13  \t& 3 & 18 & 3 & 8.0 & 3 & 7.5 && IR-N20C80 & 0.8 &   173 &  108   & 100   && 1.66 & 3   & 9  & 9 \\\\ \\noalign{\\smallskip} VW Cha ............         \t\t  & 2003 Feb 20 &S27  \t& 75 & 0.8 & 75 & 0.3 & 75 & 0.3 && IR-N20C80 & 1.2 &     77 &     76   &   72  && 1.68 &  11  & 25  &  33    \\\\ \\noalign{\\smallskip} Glass I ...............     \t\t  & 2003 Feb 20 &S13  \t& 12 & 5.0 & 12 & 3.5 & 12 & 3.0 && IR-N20C80 & 1.3 &     88 &     79   &   82  && 1.67 & 5   & 21  & 24 \\\\ \\noalign{\\smallskip} CoD $-$29$\\degr$8887 ... & 2003 Jan 20  &S13  \t& 90 & 0.3 & 90 & 0.3 & 90 & 0.4 && VIS-VIS  & 1.1 &     62 &     69   &   69  && 1.01 & 18   & 31  & 34 \\\\ \\noalign{\\smallskip} Sz 30$^c$ ..................   \t& 2003 Jan 20 &S13 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}& 3 & 20 && IR-N20C80 & 1.5 &  ...     &   ...       & 184  && 1.67 &  ...  &  ... &  7    \\\\ \\noalign{\\smallskip} Hen 3-600 ..........        \t  & 2003 Feb 20 &S13  \t& 60 & 0.4 & 60 & 0.4 & 60 & 0.4 && IR-N20C80 & 1.5 &  74    &   91     &   72    && 1.03 & 10   & 19  & 26 \\\\ \\noalign{\\smallskip} Sz 41 .................     \t\t  & 2003 Jan 22 &S27  \t& 75 & 0.8 & 75 & 0.5 & 75 & 0.5 && VIS-VIS & 0.9 &  70    &   74     &   78  && 1.64 &  11  & 27  &  30   \\\\ \\noalign{\\smallskip} CV Cha$^b$ .............        \t  & 2003 Jan 20 &S27  \t&  6 & 16 & 6 & 11 & 6 & 11 && VIS-VIS & 1.3 & 110   &   93     &   75    && 1.65 & 5   & 15  & 22 \\\\ \\noalign{\\smallskip} Sz 48 ..................\t  \t  & 2004 Apr 06\t & S13 & 4 & 48 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& IR-K & 0.6 & 184   &   ...   &   ...      && 1.66 & 2   & ...  & ...      \\\\ \\noalign{\\smallskip} BK Cha ..............\t\t  & 2004 Apr 03\t & S13 & 4 & 30 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& IR-K & 0.5\t & 78    &   ...   &   ...      && 1.66 & 7   & ...  & ...           \\\\ \\noalign{\\smallskip} Sz 60 ..................\t\t  & 2004 Apr 05\t & S13 & 4 & 60 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& IR-K & 0.5\t & 115   &   ...   &   ...      && 1.66 & 5   & ...  & ...         \\\\ \\noalign{\\smallskip} Sz 62 ..................     \t\t  & 2003 Jan 22 & S27 & 15 & 3.0 & 15 & 1.6 & 15 & 1.6 && VIS-VIS & 1.0 & 150   &  105    &   91 && 1.68 & 3   & 7  & 11     \\\\ \\noalign{\\smallskip} Herschel 4636 ....          \t  & 2003 Jan 20 & S13 & 9 & 3.0 & 9 & 4.5 & 9 & 4.5 && VIS-VIS & 1.0 & 70   &  112    & 104&& 1.08 & 15   & 17  & 17 \\\\ \\noalign{\\smallskip} ESO H$\\alpha$ 283 ......     & 2003 Mar 26 & S13 & 3 & 46 & 3 & 35 & 3 & 33 && VIS-VIS & 0.7 & 76   &    74    &   75  && 1.36 & 13   & 35  & 40     \\\\ \\noalign{\\smallskip} Sz 65 .................. \t\t  & 2004 Apr 06\t & S13 & 36 & 2.2 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& VIS-VIS  & 0.8\t & 60   &   ...   &   ... \t&& 1.36 & 20   & ...  & ...    \\\\ \\noalign{\\smallskip} Sz 68 ..................\t\t  & 2004 Apr 06\t & S13 & 88 & 0.6 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& VIS-VIS & 0.6\t & 61   &   ...   &   ...    && 1.34 &  27  & ...  & ...          \\\\ \\noalign{\\smallskip} HO Lup ..............\t\t  & 2004 Apr 06\t & S13 & 4 & 13.8 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& IR-K & 0.7 & 80   &   ...   &   ...      && 1.37 & 7   & ...  & ...  \\\\ \\noalign{\\smallskip} Sz 101 ................\t\t  & 2004 Apr 10\t & S13 & 4 & 18 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& IR-K & 0.7 & 79   &   ...   &   ...      && 1.03 & 10   & ...  & ...  \\\\ \\noalign{\\smallskip} Sz 108 ................\t\t  & 2004 Apr 10 & S13 & 8 & 8.0 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& IR-K & 0.6\t& 76   &   ...   &   ...      && 1.03 & 13   & ...  & ... \\\\ \\noalign{\\smallskip} Sz 120 ................\t\t  & 2004 Apr 10\t & S13 & 180 & 0.6 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& VIS-VIS & 0.7 & 70   &   ...   &   ...        && 1.04 & 18   & ...  & ...  \\\\ \\noalign{\\smallskip} WSB 3 ...............\t\t  & 2004 May 01 & S13 & 4 & 16.8 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& IR-K & 0.9 & 133    &   ...   &   ...         && 1.10 & 4   & ...  & ... \\\\ \\noalign{\\smallskip} WSB 11 .............\t\t  & 2004 Jun 18 & S13 & 4 & 56 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& IR-K & 1.1 & 93   &   ...   &   ...          && 1.04 & 2   & ...  & ...     \\\\ \\noalign{\\smallskip} WSB 20 .............\t\t  & 2004 May 01 & S13 & 12 & 5.0 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& IR-K & 0.8 & 75   &   ...   &   ...          && 1.06 & 12   & ...  & ... \\\\ \\noalign{\\smallskip} WSB 28 .............\t\t  & 2004 May 01 & S13 & 4 & 24 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& IR-K & 0.8 & 117   &   ...   &   ...          && 1.04 & 8   & ...  & ...     \\\\ \\noalign{\\smallskip} SR 24 .................\t \t  & 2004 May 01 & S13 & 16 & 4.0 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& IR-K & 1.1 & 72   &   ...   &   ...     && 1.03 & 20   & ...  & ...     \\\\ \\noalign{\\smallskip} Elias 2-30 ..........\t\t  & 2004 May 01 & S13 & 40 & 2.0 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& IR-K & 1.0 & 72   &   ...   &   ...      && 1.02 & 13   & ...  & ...\\\\ \\noalign{\\smallskip} WSB 46 .............\t\t  & 2004 May 01 & S27 & 4 & 13.6 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& IR-K & 0.8 & 148   &   ...   &   ...       && 1.01 & 3   & ...  & ... \\\\ \\noalign{\\smallskip} Haro 1-14c .........\t  \t  & 2004 May 01 & S27 & 52 & 1.0 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& VIS-VIS & 1.1 & 82   &   ...   &   ...        && 1.07 & 7   & ...  & ... \\\\ \\noalign{\\smallskip} ROXs 43 .............\t\t  & 2004 May 04 & S13 & 100 & 0.6 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& VIS-VIS & 0.5 & 80   &   ...   &   ...      && 1.06 &  36  & ...  & ...          \\\\ \\noalign{\\smallskip} Elias 2-49 ...........\t\t  & 2004 May 01 & S13 & 56 & 1.2 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& VIS-VIS & 0.9 & 77   &   ...   &   ...            && 1.03 & 11   & ...  & ... \\\\ \\noalign{\\smallskip} HBC 652 ............\t\t  & 2004 May 01 & S27 & 100 & 0.5 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& IR-K & 0.9 & 84   &   ...   &   ...            && 1.03 & 9   & ...  & ...  \\\\ \\noalign{\\smallskip} LkH$\\alpha$ 346 ........... \t  & 2004 May 01 & S27 & 16 & 4.0 &\\multicolumn{2}{c}{...}&\\multicolumn{2}{c}{...}&& IR-K & 0.9 & 104    &   ...   &   ...            && 1.02 &  16  & ...  & ...       \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{center} \\label{Tab:obs_log} \\begin{minipage}[position]{17cm} $^a$\\,: The seeing value quoted is a median value of all frames at 550\\,nm at zenith.    \\\\ $^b$\\,: Observed with a neutral density filter.  \\\\ $^c$\\,: Only 2 acquisition frames in Br$\\gamma$ with very poor Strehl ratio. \\\\ \\end{minipage} \\end{table*} ",
        "conclusions": "\\label{sect:Discussion}  \\subsection{Multiplicity as a function of mass} \\label{sect:mass_depend_mult} A breakdown of the distance-limited sample by primary mass, i.e. by spectral type of the primary, leads to MF/wB=38\\,$\\pm$14\\%  and CSF/wB=1.1\\,$\\pm$0.2 for K types (13 binaries, 2 triples, 6 quadruples), MF/wB=13\\,$\\pm$9\\%  and CSF/wB=0.3\\,$\\pm$0.1 for M types  (13 binaries, 2 triples). The trend is consistent with the known decrease of multiplicity with decreasing primary mass (Sterzik \\& Durisen \\cite{Sterzik_Durisen2003}).   \\subsection{Comparison between clouds} \\label{sect:clouds_comparison} The question that arises naturally is the one about the multiplicity frequency in different clouds, although the latter is of lower statistical significance than that of the total sample. The cloud with the highest value is Tau-Aur (5 triples-quadruples/11 wide binaries, MF/wB=45\\,$\\pm$20\\%, CSF/wB=1.3\\,$\\pm$0.3), followed by Cha\\,I (2/10, MF/wB=20\\,$\\pm$14\\%, CSF/wB=0.5\\,$\\pm$0.2) and Ophiuchus (2/10, MF/wB=20\\,$\\pm$14\\%, CSF/wB=0.4\\,$\\pm$0.2). Some star formation models suggest that cloud intrinsic properties might imprint different binary or multiplicity properties (Sterzik \\& Durisen \\cite{Sterzik_Durisen1998}, Sterzik et al. \\cite{Sterzik_etal2003}), but no model is able to quantify the differences observed.    \\begin{table*} \\scriptsize \\caption{Comparison of the multiplicity frequency per binary (MF/wB and MF/B) and the companion star frequency per binary (CSF/wB and CSF/B) of our work with those derived from previous multiplicity surveys among T Tauri stars in a similar separation range $\\sim$14-1700 AU$^a$, and with recent numerical simulations (bottom part). Also indicated are the number of wide (separation $>$\\,0$\\farcs$5) binaries (wB), binaries (B), triples (T), quadruples (Q) and higher multiples ($>$4).} \\begin{center} \\renewcommand{\\arraystretch}{0.5} \\setlength\\tabcolsep{7pt} \\begin{tabular}{llcccccr@{\\,$\\pm$\\,}lr@{\\,$\\pm$\\,}lr@{\\,$\\pm$\\,}lr@{\\,$\\pm$\\,}l} \\hline\\noalign{\\smallskip} Reference                      & cloud   \t& wB & B & T & Q & $>$4 &  \\multicolumn{2}{c}{MF/wB} \t&  \\multicolumn{2}{c}{CSF/wB} &  \\multicolumn{2}{c}{MF/B} \t&  \\multicolumn{2}{c}{CSF/B}\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\noalign{\\smallskip} This work           \t\t\t&   several   \t\t\t\t&  30 & 42.3 & 5 & 6 & - &  26.8 & 8.1\\% &   0.68 &\t0.13  &    20.6   &  6.2\\%  &   0.53   &   0.10 \\\\ \\noalign{\\smallskip} \\noalign{\\smallskip} \\noalign{\\smallskip} \\noalign{\\smallskip} \\multicolumn{15}{c}{Other Multiplicity Surveys}\t \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} L93           \t\t&   Tau-Aur    \t\t\t\t&  24 & 39 & 3 & 2 & - &   17.2 & 7.7\\% &   0.41\t&\t0.12  &    11.4   &  5.1\\%  &   0.27   &   0.08 \\\\ KL98  \t\t&   Tau-Aur    \t\t\t\t&  21 & 30 & 6 & 1 & - &   25.0 & 9.4\\% &   0.54\t&\t0.14  &    18.9   &  7.2\\%  &   0.41   &   0.11  \\\\ G97            \t&   Cha/Lup/CrA    \t\t\t&  19 & 21 & 2 & -  & - &   9.5   & 6.7\\% &   0.19\t&\t0.10  &    8.7     &  6.1\\%  &   0.17   &   0.09 \\\\  K00     \t\t&   Sco-Cen\t\t   \t\t& 24  & 37 & 6 & - & - &  20.0\t&\t8.2\\% \t\t&   0.40\t&\t0.12  &    14.0   &  5.7\\%  &   0.28   &   0.08\\\\ K01    \t\t&   Cha\t\t\t\t\t&  5   & 16 & 1 & - & - &  16.7  \t&\t16.7\\%\t\t&   0.33\t&\t0.24 &    5.9   &  5.9\\%  &   0.12   &   0.08\\\\ R05    \t\t&   Oph\t\t\t\t\t&  22 & 45 & 5 & - & - &  18.5  \t&\t8.3\\%\t\t&   0.37\t&\t0.12 &    10.0   &  4.5\\%  &   0.20   &   0.06\\\\ \\noalign{\\smallskip} Tot(Other surveys)\t\t&\t\t\t\t&  115 & 188 & 23 & 3 & - &  18.4  \t&\t3.6\\%\t&   0.39\t&\t0.05 &    12.2   &  2.4\\%  &   0.26   &   0.03 \\\\ Tot(All surveys)\t\t\t&\t\t\t\t&  132 & 217.3 & 23 & 7 & - &  18.5  \t&\t3.4\\%\t&   0.41\t&\t0.05 &    12.1   &  2.2\\%  &   0.27   &   0.03 \\\\ \\noalign{\\smallskip} \\noalign{\\smallskip} \\noalign{\\smallskip} \\multicolumn{15}{c}{Numerical Simulations}\t\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} SD03\t&   \t& -  & 0.21 & 0.08 & 0.03 &  & \\multicolumn{2}{c}{...} &  \\multicolumn{2}{c}{...} &   \\multicolumn{2}{c}{34\\%}\t&   \\multicolumn{2}{c}{0.78} \\\\ D03\t\t&   \t& -  & 0.154 & 0.108 & 0.045 &  & \\multicolumn{2}{c}{...} &  \\multicolumn{2}{c}{...} &  \\multicolumn{2}{c}{49.8\\%}&   \\multicolumn{2}{c}{1.14} \\\\ D04\t\t&   \t& -  & 11 & 2 & 2 & 3 & \\multicolumn{2}{c}{...} &  \\multicolumn{2}{c}{...} &  38.9  \t&\t14.7\\%\t\t&   1.28\t&\t0.27 \\\\ G04\t\t&   \t& -  &  7  & 6  & 2 & 1 & \\multicolumn{2}{c}{...} &  \\multicolumn{2}{c}{...} &  56.3  \t&\t18.8\\%\t\t&   1.38\t&\t0.29 \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{center} \\label{Tab:mult_other_surveys} \\begin{minipage}[position]{17cm} $^a$\\,: This range is assuming that all systems are at 140\\,pc. Additionally the upper limit is actually set by confusion with background/foreground stars, and varies from survey to survey. The surveys of Leinert et al. (\\cite{Leinert_etal1993}), K\\\"{o}hler \\& Leinert (\\cite{Koehler_Leinert1998}), and Ghez et al. (\\cite{Ghez_etal1997a}) have set their upper limit to 12$\\arcsec$ ($\\sim$\\,1700\\,AU at 140\\,pc), while the surveys of K\\\"{o}hler et al. (\\cite{Koehler_etal2000b}), K\\\"{o}hler (\\cite{Koehler_2001}) and Ratzka et al. (\\cite{Ratzka_etal2005}) have upper limits of 6$\\arcsec$.\\\\ References\\,: L93\\,: Leinert et al. \\cite{Leinert_etal1993}, KL98\\,: K\\\"{o}hler \\& Leinert \\cite{Koehler_Leinert1998}, G97\\,:  Ghez et al. \\cite{Ghez_etal1997a}, K00\\,: K\\\"{o}hler et al. \\cite{Koehler_etal2000b}, K01\\,: K\\\"{o}hler \\cite{Koehler_2001}, R05\\,: Ratzka et al. \\cite{Ratzka_etal2005}, SD03\\,: Sterzik \\& Durisen \\cite{Sterzik_Durisen2003}, D03\\,: Delgado-Donate et al. \\cite{Delgado_Donate_etal2003}, D04\\,: Delgado-Donate et al. \\cite{Delgado_Donate_etal2004}, G04\\,: Goodwin et al. \\cite{Goodwin_etal2004}. \\end{minipage} \\end{table*}   \\subsection{Comparison with previous multiplicity surveys} \\label{sect:multiplicity_stat_compa_surveys} We compare our result with the proportion of triples/quadruples found in previous NIR multiplicity surveys which have focussed on the brightest sources (generally K\\,$\\la$\\,10) with broadly similar separation range, sensitivity\\footnote{However, our AO survey is probably slightly more sensitive than previous speckle surveys.}, and statistical uncertainties. These are the studies by Leinert et al. (\\cite{Leinert_etal1993}) and K\\\"{o}hler \\& Leinert (\\cite{Koehler_Leinert1998}) in Tau-Aur, Ghez et al. (\\cite{Ghez_etal1997a}) in Chamaeleon, Lupus and CrA, K\\\"{o}hler et al. (\\cite{Koehler_etal2000b}) in the Scorpius-Centaurus OB association, K\\\"{o}hler (\\cite{Koehler_2001}) in Chamaeleon, and  Ratzka et al. (\\cite{Ratzka_etal2005}) in Ophiuchus. We primarily base this comparison on the multiplicity frequency per wide binary (MF/wB), as defined above (Eq.\\ref{MF/wB}), and the companion star frequency per wide binary (CSF/wB, Eq.\\ref{CSF/wB}). Since those surveys used samples drawn from lists of apparently single stars, and not wide binaries like our, they include all resolved binaries with separations down to the resolution limit. They therefore give multiplicity fractions per binary, i.e. MF/B and CSF/B, where B includes binaries with separations $\\la$0$\\farcs$5. We thus have to count in these surveys only binaries with separations larger than 0$\\farcs$5 (i.e. $\\sim$70 AU at 140\\,pc). This will give us an estimate of MF/wB and CSF/wB for these surveys, and allow a direct comparison with our own survey results. Note that in all these surveys we consider a companion number uncorrected for either completeness or chance projection, as these corrections are difficult to assess properly and moreover roughly compensating each other. While completeness depends on the sensitivity of these surveys which is rather uniform, chance projection is most of the time treated by assuming a 'safe' maximum separation which depends on the background/foreground surface density of stars. On our side, we would have to slightly restrict the separation range to the resolution achieved by those surveys (i.e. typically $\\sim$0$\\farcs$1-12$^{\\prime\\prime}$ which is $\\sim$14-1700\\,AU at 140\\,pc), but this is not significant since only one companion was detected below this limit (SR\\,24\\,C at separation 0$\\farcs$08). Table\\,\\ref{Tab:mult_other_surveys} summarizes the result of this comparison. It shows that in first order our newly derived multiplicity agrees with the previous surveys, within the uncertainties.  In a second step, we will attempt to estimate the fraction of binaries which were not included in our survey, i.e. those with a separation range $\\sim$0$\\farcs$1-0$\\farcs$5. This will allow us to convert our MF/wB to MF/B (as well as CSF/wB to CSF/B) and to perform a direct comparison with the results from the other surveys. For that purpose, we derived the fraction of binaries of those surveys with separations between $\\sim$0$\\farcs$1 and 0$\\farcs$5. We only consider those surveys for which the maximum separation was set to 12$\\arcsec$, and find a fraction of close binaries $\\alpha$=28.9\\,$\\pm$5.7\\%. The difference with the fraction derived from all surveys, including those with upper separations of 6$\\arcsec$, is not statistically different since in that case $\\alpha$=38.8\\,$\\pm$4.5\\%. The corrected number of binaries with separations $>$0$\\farcs$1 for our survey is then B=wB/(1-$\\alpha$)=30/(1-0.29)=42.3. Thus, MF/B = 20.6\\,$\\pm$6.2\\% and CSF/B = 0.53\\,$\\pm$0.10 leading to the same conclusion as above.  Noticeably, the CSF/B (and CSF/wB) is significantly higher in our survey than for all the other surveys. If we average the result for all the other surveys the difference is at the 2$\\sigma$ level. This overabundance of companions is a consequence of a higher fraction of quadruples over triples in our survey with respect to all previous surveys. The fraction of systems with multiplicity higher or equal than four to the number of systems with multiplicity higher or equal than three (called $f_4$ in the Sect.\\,\\ref{sect:MS_mult}) is 55.0\\,$\\pm$22.2\\% for our survey, while it is only 13.0\\,$\\pm$7.5\\% on average for all other surveys.    \\subsection{Comparison with theory} \\label{sect:multiplicity_stat_compa_theory} There is a probable overabundance of high-order multiples produced by the current simulations of star formation with respect to current observations. Direct comparison is not possible since theoretical multiplicity frequencies include binaries with separations down to typically $\\sim$\\,3-5\\,AU. Wider high-order companions with separations $\\ga$\\,2000\\,AU are quite rare, so that the upper bound of the separation range should be less of a concern. We will assume here that the corrections to be applied in order to obtain MF/B in the same separation range as the one probed here are minor, i.e. we will neglect the correction for systems in the separation range $\\sim$\\,3-5\\,AU-10\\,AU. In some cases, this might not be the case. For example in the simulations of Goodwin et al. (\\cite{Goodwin_etal2004}) the fraction of close binaries with separations $\\la$\\,10\\,AU would be up to half of all binaries formed (S. Goodwin 2006, private communication). We therefore have to keep in mind that the theoretically derived MF/B should be considered as lower-limits. In the following, we summarize the theoretical studies used for the comparison.  Sterzik \\& Durisen (\\cite{Sterzik_Durisen2003}) performed few-body cluster decay simulations. Although that study neglected the effect of remnant molecular gas and disk accretion and treated only the process of dynamical evolution of young small N-body clusters, it yields highly significant and robust statistics since a large number of realizations (10 000) has been computed. A degree of multiplicity of MF/B=34\\% (CSF/B = 0.78) was found. Delgado-Donate et al. (\\cite{Delgado_Donate_etal2003}) modeled the dynamical decay of a large number (a hundred) of small-N (N=5) star-forming clusters including the effects of competitive accretion and dynamical evolution through Smoothed Particles Hydrodynamics (SPH) simulations with a $\\sim$1AU spatial resolution, and found a rather high multiplicity frequency close to MF/B=50\\%. A similar high frequency of multiple systems was the outcome of two other recent and more sophisticated SPH simulations. Delgado-Donate et al. (\\cite{Delgado_Donate_etal2004}) simulated the fragmentation of 10 small-scale turbulent molecular clouds and their subsequent dynamical evolution, including this time the effect of accretion disks into the evolution of multiples. Goodwin et al. (\\cite{Goodwin_etal2004}) followed the collapse and fragmentation of 20 dense star-forming cores with a low-level of turbulence. In both cases a high frequency of high-order multiples was obtained (Table\\,\\ref{Tab:mult_other_surveys}).    \\subsection{Comparison with Main-Sequence multiplicity} \\label{sect:MS_mult}  Most of our targets will evolve towards G and K-types on the Main-Sequence, therefore we will compare our results with the field star surveys focussing on those spectral types. Duquennoy \\& Mayor (\\cite{Duquennoy_Mayor1991}, hereafter DM91) gives an estimate of the high-order multiplicity among solar-type Main-Sequence dwarfs. However selection effects may apply, as already realized in DM91, and more recent surveys and compilations (e.g. Tokovinin \\& Smekhov \\cite{Tokovinin_Smekhov_2002}) argue for a larger fraction of high-order multiples among MS dwarfs. In the following, we will compare the multiplicity fractions from the MSC catalog (Tokovinin \\cite{Tokovinin_1997}) with ours. MSC gives their results using a {\\it multiplicity fraction $f_n$} defined as (Batten\\,\\cite{Batten_1973})\\,:  \\begin{equation} f_n = \\frac{N_n}{N_{n-1}}. \\label{f_n} \\end{equation}  \\noindent where $N_{n}$ is the number of systems of multiplicity {\\it at least} $n$. Batten (\\cite{Batten_1973}) estimated $f_n$ to be of the order of 25\\%, independently of the order of multiplicity. Tokovinin (\\cite{Tokovinin_2006}) estimated from MSC a fraction of triple systems (or higher) relative to binaries (or higher) $f_3$=19\\,$\\pm$4\\% and a fraction of (at least) quadruples to (at least) triples $f_4$=22\\,$\\pm$2\\%. A similar value of $f_3$=20\\,$\\pm$6\\% was found by Tokovinin \\& Smekhov (\\cite{Tokovinin_Smekhov_2002}) in a sample of dwarfs with spectral types later than F5 extracted from MSC. This latter value is very similar to the former one, suggesting that the high-mass evolved stars included in MSC should have little effects on the multiplicity fractions. In comparison, a value of $f_3$=11\\,$\\pm$4\\% can be estimated from the sample of DM91. In our case, it happens that $f_3$=MF/B, so that $f_3$=20.6\\,$\\pm$6.2\\% and $f_4$=55.0\\,$\\pm$22.2\\%. While the fraction of triples to binaries is in agreement with what is found for the main-sequence, there is a marginally significant overabundance of quadruples with respect to triples in our survey when compared to MSC values. A clear caveat of this comparison is that multiplicity fractions for main-sequence stars are derived for much larger linear separation range than for our survey. We attempt to address this issue in the next section.     \\subsection{Total multiplicity} \\label{sect:total_mult}  We searched in the literature for spectroscopic binaries (SBs) included in our sample. Systematic searches for spectrsocopic companions of T Tauri stars comprise the radial-velocity surveys of Mathieu et al. (\\cite{Mathieu_etal1989}), Reipurth et al. (\\cite{Reipurth_etal2002}), Torres et al. (\\cite{Torres_etal2003}), and Melo (\\cite{Melo_2003}). According to them, there are 8 SBs, of which 5 need confirmation, in our sample (see Table\\,\\ref{Tab:sample}). Since not all the components of our sample have been searched for spectroscopic companions, this number is likely to be a lower limit. We can therefore tentatively estimate the increase of multiplicity that would result if we could include SBs by taking into account all those potential 8 SBs. This would lead to multiplicities MF/wB=39.0\\,$\\pm$9.8\\%  and CSF/wB=1.02\\,$\\pm$0.16 for companions with separation down to a few tens of AU (25 binaries, 8 triples, 6 quadruples, 2 quintuples)\\,\\footnote{This includes the 0$\\farcs$016 companion to ROXs\\,43\\,C (Simon et al.\\,\\cite{Simon_etal1995}).}. This also corresponds to multiplicity fractions (Eq.\\,\\ref{f_n}) $f_3$=39\\,$\\pm$10\\%, which is an upper limit here (see below), $f_4$=50\\,$\\pm$18\\%, and $f_5$=25\\,$\\pm$18\\%, i.e. the ratio of quadruples to triples is still higher than that of the main-sequence ($f_5$ is within the uncertainties of that of MSC which is 20\\,$\\pm$4\\%).  We try to better estimate $f_3$ by tentatively correcting for the binaries with separations lower than 0$\\farcs$5 as before, but not for those systems with separation between a few tenths of AU to $\\sim$\\,10\\,AU, i.e. almost pure SBs, as these systems might not be so ubiquitous (Sterzik et al.\\,\\cite{Sterzik_etal2005}, Tokovinin et al.\\,\\cite{Tokovinin_etal2006}). This gives MF/B=$f_3$=31.3\\,$\\pm$7.8\\%  and CSF/B=0.82\\,$\\pm$0.13. Therfore, the ratio of triples to binaries appears also higher for young systems than for the main-sequence, although at a marginally significant statistical level and with the caution that this result might be in part biased by the assumption we made about the paucity of tight binaries. Noticeably, a relatively good match is found with the numerical models of dynamical evolution of young small N-body clusters (see Sect.\\,\\ref{sect:multiplicity_stat_compa_theory} and Table\\,\\ref{Tab:mult_other_surveys}).    \\subsection{Hierarchical configurations} \\label{sect:hierarchy}  Among the quadruple systems, three systems are composed of two pairs separated by a distance larger in projection than the pair separations, and four systems are composed of a pair with a wide companion and an even wider companion. An almost equal number of systems between these two hierarchical configurations is also what is found among Main-Sequence systems in the MSC catalog (Tokovinin\\,\\cite{Tokovinin_2001}).   \\begin{table} \\scriptsize \\caption{Systems dynamical stability. Application of the stability criteria from Eggleton  \\& Kiseleva (1995) (their Y$^{min}_{0}$, cf Eq. 2) and from Mardling \\& Aarseth (2001) ($R^{min}_{crit} = \\frac{R^{out}_{p}}{a_{in}}$ their Eq. 90) for a given partition (1+2)+3, q$_{in}$=$m_1/m_2$ and q$_{out}$=$(m_1+m_2)/m_3$. R$_{mes}$ is the ratio of the outer to the inner projected separation of the triple partition considered$^a$.} \\begin{center} \\renewcommand{\\arraystretch}{0.7} \\setlength\\tabcolsep{7pt} \\begin{tabular}{lcccccc} \\hline\\noalign{\\smallskip} System                       \t\t& Partition$^b$ &\tq$_{in}$ \t&q$_{out}$$^c$&\tY$^{min}_{0}$$^c$\t&\tR$^{min}_{crit}$$^c$&\tR$_{mes}$\t\\\\ & \t\t\t&\t\t\t& \t\t\t&\t\t\t\t\t&\t\t\t\t\t&\t\t\t\t\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} VW Cha...........            \t&   (B+C)+A\t&\t1.14\t\t&\t0.75\t\t&\t\t6.16\t\t\t&\t\t3.92\t\t\t&\t\t5.8\t   \t \\\\ CGH$\\alpha$ 5/6........      \t&   (A+B)+C\t&\t1.84\t\t&\t0.93\t\t&\t\t5.89\t\t\t&\t\t3.75\t\t\t&\t\t23.4\t   \t \\\\ Sz 30.................       \t\t&   (A+B)+C\t&\t1.60\t\t&\t1.34\t\t&\t\t5.45\t\t\t&\t\t3.49\t\t\t&\t\t3.1\t\t  \\\\ UZ Tau.............          \t\t&   (B+C)+A\t&\t1.08\t\t&\t0.60\t\t&\t\t5.25\t\t\t&\t\t3.38\t\t\t&\t\t9.8\t   \t \\\\ J 4872...............        \t\t&   (A+B)+CD\t&\t1.10\t\t&\t1.39\t\t&\t\t5.43\t\t\t&\t\t3.48\t\t\t&\t\t19.3\t   \t \\\\ UX Tau.............          \t\t&   (A+D)+BC\t&\t6.81\t\t&\t2.40\t\t&\t\t4.81\t\t\t&\t\t3.22\t\t\t&\t\t2.2\t  \t \\\\ LkH$\\alpha$ 336..........    \t&   (A+C)+B\t&\t2.83\t\t&\t1.33\t\t&\t\t5.46\t\t\t&\t\t3.5\t\t\t&\t\t1.41\t\t  \\\\ FV Tau.............          \t\t&    (C+D)+AB\t&\t3.15\t\t&\t0.27\t\t&\t\t7.87\t\t\t&\t\t5.20\t\t\t&\t\t17.4\t  \t \\\\ GG Tau.............          \t\t&    (C+D)+AB\t&\t2.85\t\t&\t0.11\t\t&\t\t9.80\t\t\t&\t\t...$^d$\t\t&\t\t6.92\t  \t \\\\ LkH$\\alpha$ 262/263...    &    (A+B)+C\t&\t0.95\t\t&\t1.17\t\t&\t\t5.62\t\t\t&\t\t3.58\t\t\t&\t\t9.9\t   \t \\\\ LkH$\\alpha$ 346 ......... \t&    (A+B)+CD\t&\t3.11\t\t&\t0.42\t\t&\t\t7.08\t\t\t&\t\t4.57\t\t\t&\t\t5.3\t      \t \\\\ ROXs 43 ............    \t\t&    (A+B)+C\t&\t1.13\t\t&\t1.16\t\t&\t\t5.63\t\t\t&\t\t3.59\t\t\t&\t\t13.4\t        \t\\\\ SR 24 .................    \t\t&    (B+C)+A\t&\t1.79\t\t&\t0.63\t\t&\t\t6.40\t\t\t&\t\t4.09\t\t\t&\t\t62.6\t   \t \\\\ Sz 68 ..................    \t\t&    (A+B)+C\t&\t3.77\t\t&\t25.3\t\t&\t\t3.26\t\t\t&\t\t2.84\t\t\t&\t\t22.3\t       \t\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{center} \\label{Tab:syst_stability} \\begin{minipage}[position]{9cm} $^a$\\,:  In case of ambiguity the lowest value is adopted. \\\\ $^b$\\,: When several configurations exist, the more unstable is reported. \\\\ $^c$\\,:   When several mass estimates exist, these columns report the values corresponding to the greater values of Y$^{min}_{0}$ and R$^{min}_{crit}$.\\\\ $^d$\\,: Not applicable since $m_3/(m_1+m_2)>5$. \\end{minipage} \\end{table}  \\subsection{System dynamical stability} \\label{sect:syst_stability}  In order to quantitatively investigate the system dynamical stability, we applied to our sample the semi-empirical criteria developed by Eggleton  \\& Kiseleva (1995) and Mardling \\& Aarseth (2001). The result is shown in Table\\,\\ref{Tab:syst_stability} where R$_{mes}$ should be at least larger than Y$^{min}_{0}$ or R$^{min}_{crit}$ to ensure stability. Assumptions are that the periastron of the outer orbit as well as the semi-major axis and the apastron of the inner orbit are approximated by their projected separations, i.e. the excentricity of the outer orbit is assumed to be zero and the system is supposed to be seen face-on. The component masses used are taken from the literature when available and quoted in Table\\,\\ref{Tab:systems_accretion}, or roughly estimated by us using flux ratios and the assumption of coevality (Sect.\\,\\ref{sect:comp_Spt_mass_age}). Under the above assumptions, 3 systems (LkH$\\alpha$\\,336, Sz\\,30 and UX\\,Tau) appear to be unstable by both criteria while three others (LkH$\\alpha$ 346, VW\\,Cha and GG\\,Tau) appear unstable by the Eggleton  \\& Kiseleva (1995) criterion only. In that respect, it is interesting to point out that recently, Beust \\& Dutrey (\\cite{Beust_Dutrey_2006}) have shown through a detailed analysis of the dynamics of the GG\\,Tau system that only some special orbital configuration would allow GG\\,Tau CD to be stable, although this result depends strongly on the adopted size of the circumbinary disk around GG\\,Tau\\,AB.  Even without the values reported in Table\\,\\ref{Tab:syst_stability}, one can immediately note from Fig.\\,\\ref{fig:triples_quadruples} that the three systems LkH$\\alpha$\\,336, Sz\\,30 and UX\\,Tau are apparently non-hierarchical. This corresponds to a frequency 3/14=21$\\pm$12\\%. Note that K\\\"{o}hler \\& Leinert (\\cite{Koehler_Leinert1998}) found 1 apparent non-hierarchical triple among 7 triples, while K\\\"{o}hler et al (\\cite{Koehler_etal2000b}) and Ratzka et al. (\\cite{Ratzka_etal2005}) found1/6 and 1/5, respectively. This yields a fraction of apparent non-hierarchical triples among triples of $\\sim$15-20\\%, and in combination with our survey, a total fraction of 6/32=19$\\pm$8\\%.   Ambartsumian (\\cite{Ambartsumian_1954}) provided a detailed statistical analysis of how many Trapezium systems are in fact pseudo-Trapezium systems, i.e. projected hierarchical systems. With the assumption of circular orbits, the probability for a hierarchical triple system to be seen in a pseudo-Trapezium configuration is (his Eq. 10)\\,:  \\begin{equation} U = \\frac{lg2-\\frac{1}{2}-\\frac{1}{24}\\frac{k_0^2}{k_1^2}}{lgk_1-lgk_0}. \\label{Amba_proba} \\end{equation}  \\noindent where $k_0 < AC/AB < k_1$, AC and AB are the maximum and minimum projected separations, respectively, $k_0$ and $k_1$ their ratios in the case of a Trapezium and hierarchical configuration, respectively. In our multiple system sample, we have $k_0 \\sim$\\,2 and $k_1\\sim$\\,10, thus U=0.05. Following Ambartsumian (\\cite{Ambartsumian_1954}), we include the quadruple systems by considering an equal number of the two hierarchical configurations discussed in Sect.\\,\\ref{sect:hierarchy} and a fraction of quadruples among systems with multiplicity higher than 2 of one half. This leads to a probability P=0.07 that apparent non-hierarchical systems are in fact projected hierarchical systems and this is in statistical agreement with the value we found.  However, if we suppose that the relative orbital inclinations could be rather small, as the study of Sterzik \\& Tokovinin (\\cite{Sterzik_Tokovinin_2002}) among field triple systems seems to indicate, then the fraction of projected hierarchical systems among apparent non-hierarchical systems would be much smaller and 1 or 2 of our 3 systems could be really non-hierarchical. These systems would then be either very young, or in re-arrangement after a recent close triple approach following or not the ejection of a companion. In all cases, these systems are not expected to be long-term stable and may later decay. For example, Delgado-Donate et al. (\\cite{Delgado_Donate_etal2003}) found that the non-hierarchical systems formed in their simulations decay in less than $\\sim$0.5\\,Myr. UX\\,Tau is the only systems with an estimated age for A, D and the combined BC system (see Table\\,\\ref{Tab:systems_accretion}), and it seems not consistent with a very young ($\\la$\\,1\\,Myr) system. On the other hand, this could indeed be the case for LkH$\\alpha$\\,336.       \\subsection{Disk evolution} \\label{sect:disk_evolution}   \\begin{table*} \\scriptsize \\caption{Known T Tauri types (CTTS or WTTS) for all single components or otherwise pairs of the triple/quadruple systems. Also included are spectral types, A$_V$, masses, and ages of the components from the literature, as well as our own estimates for the new and poorly studied components when possible (see Sect.\\,\\ref{sect:comp_Spt_mass_age}).} \\begin{center} \\renewcommand{\\arraystretch}{0.7} \\setlength\\tabcolsep{7pt} \\begin{tabular}{l@{\\hspace{2mm}}c@{\\hspace{3mm}}l@{\\hspace{0.5mm}}rlccr@{\\hspace{2mm}}rl@{\\hspace{2mm}}c@{\\hspace{3mm}}c @{\\hspace{1mm}}r@{\\,$\\pm$\\,}l@{\\hspace{1mm}}r@{\\,$\\pm$\\,}l@{\\hspace{1mm}}r} \\hline\\noalign{\\smallskip} System   &   Comp.  \t& other name \t&    HBC \t&   SpT & A$_V$ &   Ref  & EW(H${_\\alpha}$)  & $\\Delta$(K$-$L) & Type & Ref  \t\t     & Ref  \t    &  \\multicolumn{2}{c}{Mass}  &  \\multicolumn{2}{c}{Age}   & Ref   \\\\ &    or pair   &           \t\t&         \t&\t      &\t            &  (SpT) &    [\\AA]    \t\t\t&  [mag]   \t\t      &           & (H${_\\alpha}$)  & (K$-$L)&  \\multicolumn{2}{c}{[M$_{\\sun}$]}  &  \\multicolumn{2}{c}{[Myr]} &   \t\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} CGH$\\alpha$ 5/6 ........ &    A-B \t& CGH$\\alpha$ 6\t&\t&  K7        \t& ... & 24  \t&      5.0  &  ...  \t&   W &  24   &    & \\multicolumn{2}{c}{...} &  \\multicolumn{2}{c}{...} &  \\\\ &     A \t& \t\t\t \t&   \t&  K7  \t&  & 24  \t&  ...\t       &  ...   \t&  ...   &  \t    &     &\\multicolumn{2}{c}{0.79\\,$^{+0.09}_{-0.03}$}&\\multicolumn{2}{c}{$<$\\,1}& 1\\\\ &     B \t& \t\t\t \t&   \t&  M0-4  \t&  & 1  \t&  ...\t       &  ...   \t&  ...   &  \t    &     &\\multicolumn{2}{c}{0.43\\,$^{+0.26}_{-0.14}$}&\\multicolumn{2}{c}{$<$\\,1}& 1\\\\ &     C \t& CGH$\\alpha$ 5 \t&   \t&  K2-5  \t& ... & 24  \t&  126.9  &  ...   \t&   C  &  24   &    &\\multicolumn{2}{c}{1.31\\,$^{+0.09}_{-0.3}$}&\\multicolumn{2}{c}{5\\,$^{+8}_{-4}$}& 1\\\\ \\noalign{\\smallskip} VW Cha ........... &     A              \t& \t& 575    & K5-7   & 1.37 &  2  \t&      7.9-116  &  ...\t&   C \t&   2-3  &    &\\multicolumn{2}{c}{1.00}&\\multicolumn{2}{c}{0.4} & 21 \\\\ &     B-C           \t&\t&            &   K7     & 1.23 &  2  \t&      0.48   & \t...\t&   W   \t&   2     &     &\\multicolumn{2}{c}{0.40-0.35} & \\multicolumn{2}{c}{0.4}& 21 \\\\ &     D \t\t&\tSz 23  &&   M2.5 &  6.4  &  8  \t&      ...       \t & \t...\t&   C$^a$ &           &     &\\multicolumn{2}{c}{0.20} & \\multicolumn{2}{c}{0.3} & 8 \\\\ \\noalign{\\smallskip} Sz 30 ................. &     A        &\t&   \t     &  M0.5  \t& 0.58  &  2  \t\t&    11.0      & \t...\t&   C    &   2   &    &\\multicolumn{2}{c}{0.48}&\\multicolumn{2}{c}{2} & 2 \\\\ &     B        &\t&   \t     &  M2     \t& 0.19  &  2  \t\t&      2.6      &       ...\t&   W   &   2   &    &\\multicolumn{2}{c}{0.30}&\\multicolumn{2}{c}{2} & 2 \\\\ &     C$^{c1}$        &\t&   &  M3-5 \t&   & 1   \t&      ...       \t &       ...\t&   ...    &        &    &\\multicolumn{2}{c}{0.23\\,$^{+0.14}_{-0.10}$}&\\multicolumn{2}{c}{4\\,$^{+6}_{-3}$}&1\\\\ &     C$^{c2}$        &\t&   &  K8-M3 \t&   & 1   \t&      ...       \t &       ...\t&   ...    &        &    &\\multicolumn{2}{c}{0.58\\,$^{+0.23}_{-0.18}$}&\\multicolumn{2}{c}{17\\,$^{+47}_{-13}$}&1\\\\ \\noalign{\\smallskip} LkH$\\alpha$ 336 .......... &     A     &LkH$\\alpha$ 336\t&  190   &  K7    & 1.03 &  4,9  &      24.0-35.0   & \t...\t&   C   \t&    4,9   &   &\\multicolumn{2}{c}{0.51} &\\multicolumn{2}{c}{0.5} &  9 \\\\ &     B     &LkH$\\alpha$ 336/c&  516  &  M0.5     & 2.62 &  4,9  &        5.1-8.0   & \t...\t&   W   \t&    4,9   &   &\\multicolumn{2}{c}{0.52} &\\multicolumn{2}{c}{2.2} &  9  \\\\ &     C$^{c1}$     &\t&  \t    &  M5-6\t & &  1  &        ...   & \t...\t&   ...   \t&         &   \t  &\\multicolumn{2}{c}{0.18\\,$^{+0.10}_{-0.05}$}&\\multicolumn{2}{c}{$<$\\,1}&1 \\\\ &     C$^{c2}$     &\t&  \t    &  K3-9\t & &  1  &        ...   & \t...\t&   ...   \t&         &   \t  &\\multicolumn{2}{c}{1.18\\,$^{+0.22}_{-0.42}$}&\\multicolumn{2}{c}{4\\,$^{+5}_{-3}$}&1 \\\\ \\noalign{\\smallskip} J 4872 ............... &     A-B              &\t&   \t     &  K9   & ... &  5  \t\t&        0.8   & \t...\t&   W   &    5  &   &\\multicolumn{2}{c}{...} &\\multicolumn{2}{c}{...} &   \\\\ &     C-D             &\t&   \t     &  M1  & ... &  5  \t\t&        4.2   & \t...\t&   W   &    5  &   &\\multicolumn{2}{c}{...} &\\multicolumn{2}{c}{...} &  \\\\ &     A             \t  &\t&   \t     &  K9  & &  5  \t\t&        ...\t & \t...\t&   ...\t   &        &   &\\multicolumn{2}{c}{0.78\\,$^{+0.03}_{-0.06}$}&\\multicolumn{2}{c}{10\\,$^{+15}_{-6}$}&1\\\\ &     B             \t  &\t&   \t     &  K5-M2 & &  1  \t\t&        ...\t & \t...\t&   ...\t   &        &   &\\multicolumn{2}{c}{0.71\\,$^{+0.31}_{-0.23}$}&\\multicolumn{2}{c}{10\\,$^{+15}_{-6}$}&1\\\\ &     C             \t  &\t&   \t     &  M1  \t    & &  5  \t\t&        ...\t & \t...\t&   ...\t   &        &   &\\multicolumn{2}{c}{0.59\\,$^{+0.14}_{-0.11}$}&\\multicolumn{2}{c}{17\\,$^{+18}_{-11}$}&1\\\\ &     D             \t  &\t&   \t     &  M0-3    & &  1  \t\t&        ... \t & \t...\t&   ...\t   &        &   &\\multicolumn{2}{c}{0.48\\,$^{+0.23}_{-0.13}$}&\\multicolumn{2}{c}{17\\,$^{+18}_{-11}$}&1\\\\ \\noalign{\\smallskip} UX Tau ............. &     A                  &\tUX Tau E\t\t& 43      &  K4-5  &0.26&  5,6  \t&     3.9-9.5    \t\t&     0.65\t&   C  \t&  4-6  \t\t&  6  \t& 1.09 & 1.4 & 4.5 & 1.5 & 6 \\\\ &     B-C              &\tUX Tau W\t\t& 42      &  M2  \t  &0.26&  5  \t&     4.0-4.5    \t\t&     -0.08\t&   W   \t& 4-6   \t\t&  6  \t&  0.52 & 1.4 & 5.9 & 1.9 & 6  \\\\ &     D                  &\t\t\t\t&            &  M3-5  &0.57&  5,6  \t&       8.5         \t\t&     0.15\t&   W   \t&   6     \t\t&  6 \t&  0.16 & 1.3 & 1.6 & 4.1 & 6 \\\\ \\noalign{\\smallskip} UZ Tau ............. &     B-C              &\tUZ Tau W\t\t& 53      &  M3  \t  & &  4,9  \t&      80.0    \t\t&       ...\t&         &   7     \t&         &  \\multicolumn{2}{c}{...}  &\\multicolumn{2}{c}{...} &   \\\\ &     B                  &\t\t\t\t&            &  M2  \t  & 0.55 &  6,20     \t&      54.0       \t\t&       0.68\t&   C   &   20    \t&   6   & 0.56 &1.3 & 2.6 & 1.8 & 6 \\\\ &     C                  &\t\t\t\t&            &  M2-3  & 1.75 &  6,20     \t&      97.0     \t\t&       0.38\t&   C   &   20    \t&   6   & 0.52 & 1.3 & 3.2 & 1.8 & 6 \\\\ &     A                  &\tUZ Tau E\t\t&  52     &  M1  \t   & 0.33 &  4,9,6  \t&      74-82       \t\t&       ...\t&   C   &   6,7 \t&    \t  & 0.65 & 1.2 & 5.1 & 2.3 & 6 \\\\ \\noalign{\\smallskip} FV Tau ............. &     A     \t            &\t\t\t\t&  386  & K5   \t    &5.40&  4,9,20 \t&       6.0-15.0    \t&    0.21\t&   C    \t&   6,20    &  6  & 1.12 & 1.3 & 3.7 & 1.5 & 6 \\\\ &     B\t            &\t\t\t\t&           & K6-M2   &5.40&  6,20   \t&    41.0-63.0   \t\t&    0.94\t&   C    \t&   6,20    &  6  & 0.89 & 1.3 & 6.3 & 1.7 & 6 \\\\ &     C\t            &\tFV Tau/c\t\t&  387  & M2.5-3  &3.25&  4,9,20 \t&    17.0-21.0     \t&    0.18\t&   W$^d$&   6,20    &  6   & 0.41 & 1.3 & 1.8 & 2.3 & 6 \\\\ &     D\t            &\tFV Tau/c\t\t&           & M3.5-5  &7.00&  6,20  \t&  224.0-800       \t&    0.86\t&   C$^e$ &   6,20    &  6  & 0.13 & 1.2 & 19 & 2.8 & 6 \\\\ \\noalign{\\smallskip} GG Tau ............. &     A &\tGG Tau Aa\t& 54     &  K7-M0      &0.30& 10,20   \t&   $ >$42-57.0    \t&    0.71\t&  C  &   6,20 &   6   & 0.76 & 0.09 & 1.0 & 2.2 & 6 \\\\ &     B &\tGG Tau Ab\t&           &  M0.5-2     &0.45& 10,20   \t&    16.0-21.0  \t\t&    0.62\t&  C  &   6,20 &   6   & 0.68 & 0.02 & 1.1 & 1.8 & 6 \\\\ &     C &\tGG Tau Ba\t&           &  M5.5        &0.00&  11  \t&    21.0-22.0 \t\t&    0.52\t&  C  &   6,5   &   6    & 0.12 & 0.03 & 1.8 & 2.5 & 6 \\\\ &     D &\tGG Tau Bb\t&   \t     &  M7.5  \t&0.00&  11  \t&    32.0-19.0  \t\t&    0.43\t&  C  &   6,5   &   6    & 0.042 & 0.019 &\\multicolumn{2}{c}{$<$1} & 6 \\\\ \\noalign{\\smallskip} LkH$\\alpha$ 262/263... &    A-B && 9  \t     & M3-4 &0.00& 13,14,23 & 30.0 &  0.32 &  :C  & 4,14 & 13  & \\multicolumn{2}{l}{$\\sim$0.45-0.45} &\\multicolumn{2}{c}{$\\sim$2}& 12 \\\\ &     A  &\t&\t     & M3\t&&  13,14,23 & ...      & ...\t&   ...  &    &    & \\multicolumn{2}{c}{0.40\\,$^{+0.15}_{-0.13}$} &\\multicolumn{2}{c}{2\\,$^{+4}_{-1}$}& 1 \\\\ &     B  &\t&\t     & M3\t&&  13,14,23 & ...      & ...\t&   ...  &    &    & \\multicolumn{2}{c}{0.42\\,$^{+0.17}_{-0.12}$} &\\multicolumn{2}{c}{2\\,$^{+4}_{-1}$}& 1 \\\\ &     C  &\t&\t     & M0\t&& 12 \t\t   & ...      & ...\t&   C  &    &    & \\multicolumn{2}{l}{$\\sim$0.7} & \\multicolumn{2}{c}{$\\sim$2} & 12 \\\\ &     D  & LkH$\\alpha$ 262 & 8  & M0   &1.17& 4,13,14 \t&     31.0   &  0.62 &   C & 4,14 & 13 &  \\multicolumn{2}{c}{0.69\\,$^{+0.08}_{-0.09}$}& \\multicolumn{2}{c}{2\\,$^{+4}_{-1}$}& 1 \\\\ \\noalign{\\smallskip}  LkH$\\alpha$ 346 ........... &     A-B-C \t  & LkH$\\alpha$ 346 SE &   275    & M5  & 0.65 & 4  &  18.0  &  ... &  W  & 4  &  & \\multicolumn{2}{c}{...} & \\multicolumn{2}{c}{...} & \\\\ &     A \t  \t  &  \t\t\t\t&\t&  M5    &&     1\t&      ...   \t   &  ...  &  ...  &  &  &  \\multicolumn{2}{c}{0.28\\,$^{+0.16}_{-0.12}$} &  \\multicolumn{2}{c}{$<$\\,1} & 1  \\\\ &     B \t  \t  &  \t\t\t\t&\t&  M6-7 &&     1\t&      ...\t   &  ...  &  ...  &   &  &  \\multicolumn{2}{c}{0.09\\,$^{+0.03}_{-0.02}$} & \\multicolumn{2}{c}{$<$\\,1} & 1  \\\\ &     C \t  \t  &  \t\t\t\t&\t&  M6-7 &&     1\t&      ...\t   &  ...  &  ...  &   &   &  \\multicolumn{2}{c}{0.09\\,$^{+0.03}_{-0.01}$} & \\multicolumn{2}{c}{$<$\\,1} & 1  \\\\ &     D\t\t  & LkH$\\alpha$ 346 NW &   275    & K7 & 1.16 & 4 & 28.0      &  ...  &  C  & 4 &   & \\multicolumn{2}{c}{0.8\\,$^{+0.16}_{-0.04}$} & \\multicolumn{2}{c}{2\\,$^{+5}_{-1}$} & 1  \\\\ \\noalign{\\smallskip}  ROXs 43 .............. &     A-B   & NTTS 162819-2423S &  &  G0         &2.3& 18,22       &    0.4-5.6 &  0.56-0.98 &  C   & 18 & 18-25  &  \\multicolumn{2}{l}{$\\sim$1.5$^b$} & \\multicolumn{2}{c}{$\\sim$7$^b$} & 22 \\\\ &     B       & \t&       \t\t         &  K5-M2\t&&  1         &   ... &  ...  &  ...    &       &        & \\multicolumn{2}{c}{0.71\\,$^{+0.25}_{-0.21}$} &  \\multicolumn{2}{c}{$\\sim$7$^b$} & 1 \\\\ &     C       & NTTS 162819-2423N & &  K3-K5\t&1.3&  18,22 &  0.9\t&  0.09-0.35   &  W  &  18 & 18-25 &  \\multicolumn{2}{c}{$\\sim$1.3$^b$} & \\multicolumn{2}{c}{$\\sim$3$^b$} & 22 \\\\ \\noalign{\\smallskip} SR 24 ................. &     A      &  SR 24S  &  262  &  K2\t   &4.49& 4    & 76.0  &    0.99   &  C &  4 &  17  & \\multicolumn{2}{c}{$>$ 1.4}& \\multicolumn{2}{c}{1\\,$^{+7}$} & 1\\\\ &     B-C  &  SR 24N &  262  &  M0.5\t   &3.07& 4    & 24.0  &    0.69   &  C &  4 &  17  &  \\multicolumn{2}{c}{ ...} & \\multicolumn{2}{c}{ ...}  \\\\ &     B &   &      \t\t\t    &  K4-M4 && 1    &   ...     &    1.38   &  C &     &   25  &\\multicolumn{2}{c}{0.61\\,$^{+0.6}_{-0.27}$} & \\multicolumn{2}{c}{1\\,$^{+7}$} & 1 \\\\ &     C &   &   \t\t\t    &  K7-M5 && 1    &   ...     &    1.07   &  C &      &  25  &\\multicolumn{2}{c}{0.34\\,$^{+0.46}_{-0.18}$} & \\multicolumn{2}{c}{1\\,$^{+7}$} & 1 \\\\ \\noalign{\\smallskip} Sz 68 .................. &     A-B  &   &  248 &   K2\t\t&4.00&  3,16,19   &  6.8 &    0.26    &  C   & 3  &   16  & 2.5 & 0.1 & \\multicolumn{2}{c}{$<$ 0.1} & 16  \\\\ &     B      &   &          &   K5-M1\t&&  1 \t            &   ...   &    ...\t     &  ...   &     &          &\\multicolumn{2}{c}{1.10\\,$^{+0.40}_{-0.37}$}& \\multicolumn{2}{c}{$<$ 0.1} & 1  \\\\ &     C      &   &          &   M6\t\t&0.00&  16\t\t   &    ...  &    0.12    &  W  &     &   16  & \\multicolumn{2}{c}{0.1\\,$^{+0.1}_{-0.02}$}& 0.2 & 0.1 & 16  \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{center} \\label{Tab:systems_accretion} \\begin{minipage}[position]{18cm} Note\\,: HBC = catalog entry number in Herbig \\& Bell\\,\\cite{Herbig_Bell_1988}. \\\\ $^a$\\,: The T Tauri type classification of Sz 23 is from G\\`{o}mez \\& Mardones\\,\\cite{Gomez_Mardones_2003}.  \\\\ $^b$\\,: Estimated by eye from the plot of Bouvier \\& Appenzeller\\,\\cite{Bouvier_Appenzeller_1992}.\\\\ $^{c1}$\\,: Coeval with component A. $^{c2}$\\,: Coeval with component B.\\\\ $^d$\\,: Passive disk (McCabe et al.\\,\\cite{McCabe_etal2006}). \\\\ $^e$\\,; IRC candidate (Woitas et al.\\,\\cite{Woitas_etal2001}, see also McCabe et al.\\,\\cite{McCabe_etal2006}). \\\\ References\\,: (1) This work. (2) Brandner \\& Zinnecker\\,\\cite{Brandner_Zinnecker_1997}. (3) Appenzeller et al.\\,\\cite{Appenzeller_1983} - Rydgren\\,\\cite{Rydgren_1980}. (4) Cohen \\& Kuhi\\,\\cite{Cohen_Kuhi_1979}. (5) Duch\\^{e}ne et al.\\,\\cite{Duchene_etal1999}. (6) White \\& Ghez\\,\\cite{White_Ghez_2001}. (7) Cohen \\& Kuhi\\,\\cite{Cohen_Kuhi_1979} - Edwards et al.\\,\\cite{Edwards_1987}. (8) G\\`{o}mez \\& Mardones\\,\\cite{Gomez_Mardones_2003}. (9) Hartigan et al.\\,\\cite{Hartigan_etal1994}. (10) White et al.\\,\\cite{White_etal1999}. (11) Luhman\\,\\cite{Luhman_1999}. (12) Jayawardhana et al.\\,\\cite{Jayawardhana_etal2002}. (13) Jayawardhana et al.\\,\\cite{Jayawardhana_etal2001}. (14) Herbst et al.\\,\\cite{Herbst_etal2004}. (15) Perryman et al. 1997. (16) Prato et al.\\,\\cite{Prato_etal2003}. (17) Rydgren 1976. (18) Walter et al. 1994. (19) Hughes et al. 1994. (20) Hartigan \\& Kenyon\\,\\cite{Hartigan_Kenyon2003}. (21) Brandeker et al.\\,\\cite{Brandeker_etal2001}. (22) Bouvier \\& Appenzeller\\,\\cite{Bouvier_Appenzeller_1992}. (23) Luhman\\,\\cite{Luhman_2001}. (24) Reipurth \\& Pettersson\\,\\cite{Reipurth_Pettersson1993}. (25) McCabe et al.\\,\\cite{McCabe_etal2006}.    \\end{minipage} \\end{table*}   The study of disk evolution in young multiple system allows to investigate samples of stars of various masses that are {\\it a priori} coeval. It is however necessary to first understand the influence of multiplicity on disk evolution, i.e. in which conditions the disks can be considered to evolve in isolation in multiple systems. This has been the subject of several early studies among T Tauri binaries (see Monin et al.\\,\\cite{Monin_etal_2006_PPV} and references therein for a recent review). Only a few attempts to investigate disk evolution of higher order multiple systems have been achieved recently (White et al.\\,\\cite{White_etal2002}). This is however of particular interest since both close and wider pairs are usually present in such hierarchical systems and all components of a system are supposed to be coeval. In the following, we analyze the frequency of disks in our sample of triples/quadruples through the use of various disk and/or accretion diagnostics and compare to what is known for isolated binaries.  We performed a compilation of T Tauri types for both individual components and close pairs of the triple/quadruple systems from the available literature (Table\\,\\ref{Tab:systems_accretion}). Two criteria were used in order to assign these types\\,: first, the equivalent width of the H$\\alpha$ line with a threshold between WTTS and CTTS at about 10\\,\\AA, depending on the spectral type (Martin\\,\\cite{Martin_1998}, Hartigan \\& Kenyon\\,\\cite{Hartigan_Kenyon2003}); second, the $\\Delta$(K-L) color excess which was shown to correlate with both strong [OI] and H$\\alpha$ emission lines (Edwards et al.\\,\\cite{Edwards_etal1993}). In the latter case $\\Delta$(K-L)$>$0.4 indicates the presence of an optically thick accretion disk (Edwards et al.\\,\\cite{Edwards_etal1993}). The $\\Delta$(K-L) color excess was computed from K-L measurements and A$_V$ available in literature (for which the reference is reported in Table\\,\\ref{Tab:systems_accretion}), spectral types either from the literature or estimated by us and intrinsic colors of Kenyon \\& Hartmann (\\cite{Kenyon_Hartmann_1995}). For some systems individual component T Tauri types are missing and only the global type of the usually tightest pair is available. However, one can consider that a tight pair presenting spectral WTTS properties is in fact composed of two WTTS, while no conclusion is possible in the case of a tight pair for which the composite spectrum shows CTTS properties. We recently obtained spatially-resolved spectroscopy for some of these tight systems and the results will be presented elsewhere.  While the number of systems in Table\\,\\ref{Tab:systems_accretion} does not allow us to draw any firm conclusions, some trends can nevertheless be identified. Among the 9 systems with CTTS/WTTS information known and/or deduced for each component, one is presumably composed of only WTTS (J\\,4872). Since this system does not contain information about differential disk evolution, we will focus on the remaining 8 systems (CG\\,H$\\alpha$\\,5/6, VW\\,Cha, UX\\,Tau, UZ\\,Tau, FV\\,Tau, GG\\,Tau, LkH$\\alpha$\\,346, SR\\,24). Half of these\\,\\footnote{FV\\,Tau, as discussed below in the text, is not a mixed system.} are systems of mixed type (i.e. at least one component with a different type). The fraction of mixed systems is even higher if one includes the 4 systems that are already mixed even without the full knowledge of all components' T Tauri types (i.e. Sz\\,30, LkH$\\alpha$\\,336, ROXs\\,43, Sz\\,68). We end up with a fraction of mixed systems of $\\sim$\\,65\\,\\% (8/12), similar or even higher to what is found among binaries ($\\sim$\\,40\\,\\%, Monin et al.\\,\\cite{Monin_etal_2006_PPV}). The much higher fraction of Taurus systems among non-mixed systems (3/4, compared to 1/8 for mixed systems) supports the suggestion of a regional dependance of the fraction of mixed systems, or at least a lower fraction in Taurus as in the case of binaries (Monin et al.\\,\\cite{Monin_etal_2006_PPV}). On the other hand, we do not find any difference in the distribution of estimated age in mixed- and non-mixed systems which could explain the existence of these two populations, i.e. the age of non-mixed systems is not significantly younger than that of mixed systems.  Disk diagnostics in all close pairs (separation $\\la$\\,100\\,AU) are consistent with what is expected from tidal truncation theory (Armitage et al.\\,\\cite{Armitage_1999}). All close CTTS pairs (GG\\,Tau\\,AB, UZ\\,Tau\\,BC, SR\\,24\\,BC, FV\\,Tau\\,AB) exhibit estimated mass ratios q$\\ga$\\,0.5. SR\\,24\\,BC with q$\\sim$\\,0.5 has the lowest value, and this particular case may need disk replenishment from a common envelope with material falling preferentially onto the secondary's disk (Bate \\& Bonnell\\,\\cite{Bate_Bonnell1997}) in order to prevent the more truncated of the secondary disk to be dissipated prior that of the primary. In general, all four close CTTS pairs may well need disk replenishment in order for the highly truncated disks to sustain active accretion. While both GG\\,Tau\\,AB and SR\\,24\\,BC are known to be surrounded by circumbinary disks for replenishment (Dutrey et al.\\,\\cite{Dutrey_etal1994}, Andrews \\& Williams\\,\\cite{Andrews_Williams2005}), the cases of FV\\,Tau\\,AB and UZ\\,Tau\\,BC seem less obvious since no circumbinary disk has been detected so far. A massive circumbinary disk is known to be surrounding the spectroscopic pair UZ\\,Tau\\,A located some 500\\,AU from UZ\\,Tau\\,BC but it is very unlikely responsible for the resplenishment of the BC pair. It is interesting to note that only gas emission in the CO line was detected around the pair SR\\,24\\,BC (Andrews \\& Williams\\,\\cite{Andrews_Williams2005}) which suggests that other circumbinary disks or common reservoir could have been missed by previous continuum millimeter surveys. The case of the FV\\,Tau\\,CD pair deserves special comment since it seems to be, together with VW\\,Cha\\,A-BC, a close pair of mixed type. While the latter is consistent with exhausted highly truncated secondary disks, FV\\,Tau\\,CD shows a WTTS-CTTS type despite an estimated q$\\sim$\\,0.25 which would imply a more tidally truncated secondary disk. However, one should note that this apparent contradiction can be solved if one considers that component C is in fact a passive disk (i.e. a non-accreting disk, e.g. McCabe et al.\\,\\cite{McCabe_etal2006}) and component D an IRC candidate (Woitas et al.\\,\\cite{Woitas_etal2001}, see also McCabe et al.\\,\\cite{McCabe_etal2006}). Therefore, while the origin of the IRC phenomenon is still under debate, we can speculate that the latter is likely to be more massive, i.e. the primary. This implies that the close pair FV\\,Tau\\,CD which has two disks would have a star mass ratio presumably closer to unity, hence in agreement with synchronized disk evolution.  The fact that 4 mixed systems include both close {WTTS} pairs and well-separated components with disks seems to suggest that the highly truncated disks of these close pairs were not fed by a circumbinary envelope and hence became rapidily exhausted. Alternatively a reservoir could have originally existed in such multiple systems and be subsequently disrupted by dynamical interactions with the other component(s) of these systems.          In this paper, we reported on our survey for high-order multiplicity among wide visual Pre-Main Sequence (PMS) binaries conducted with NACO at the VLT. The main conclusions of our study are summarized as follows\\,:  \\noindent (1) Among the 58 PMS wide binaries surveyed, which comprises 52 T Tauri systems from various star-forming regions, we found 7 triple systems (2 new) and 7 quadruple systems (1 new). The new close companions are most likely physically bound based on their probability of chance projection and on their position on a color-color diagram. Some systems might still be not physical and future spectroscopy and common-proper motion measurement will be able to give an answer soon. The corresponding degree of multiplicity among wide binaries is MF/wB=26.9\\,$\\pm$7.2\\% in the projected separation range 0\\farcs07-12$^{\\prime\\prime}$, with the largest contribution from the Taurus cloud. Considering a restricted sample composed of systems at distance 140-190\\,pc, we obtained only a slightly different value of MF/wB=26.8\\,$\\pm$8.1\\%, in the separation range 10/14 AU - 1700/2300 AU (30 binaries, 5 triples, 6 quadruples).  \\noindent (2) A general trend of decreasing multiplicity with lower primary mass is found, consistent with the known decrease of multiplicity with primary mass mainly found among binaries (Sterzik \\& Durisen \\cite{Sterzik_Durisen2003}).  \\noindent (3) Comparison with previous multiplicity surveys focusing on the brightest sources with similar resolution and (perhaps slightly lower) sensitivity shows that there is a good agreement between our newly derived multiplicity frequency per wide binary and those derived from these surveys, in the common separation range $\\sim$14-1700 AU, although a significant overabundance of quadruple systems compared to triple systems is apparent. Considering that 4 out of the 6 quadruples of our distance-limited sample are from Taurus and that sources of that region have a frequency of quadruples to triples $f_4$=39\\,$\\pm$17\\% when combining the results from previous surveys with ours, indicates that this excess of quadruples is mainly due to Taurus.  \\noindent (4) Tentatively including the spectroscopic pairs to our restricted sample and comparing the multiplicity fractions to those measured from solar-type main-sequence stars in the solar neighborhood leads to the conclusion that both the ratio of triples to binaries ($f_3$) and the ratio of quadruples to triples ($f_4$) seems to be in excess among young stars. While the former is not statistically significant, and depends on the assumption of paucity of spectroscopic binaries with no third component, the latter may be statistically significant.  \\noindent (5) Our multiplicity frequency per binary seems lower than current predictions from numerical simulations of multiple star formation, especially SPH simulations. Noticeably, however, a relatively good match is found with the numerical models of the dynamical evolution of young small N-body clusters.  \\noindent (6) Some systems might not be dynamically stable according to the criteria of Eggleton  \\& Kiseleva (\\cite{Eggleton_Kiseleva_1995}) and Mardling \\& Aarseth (\\cite{Mardling_Aarseth_2001}). However, the three apparent non-hierarchical triple systems could be real if one assumes some correlation in the orbital inclinations.  In one case, the age of the system could be young enough to explain this configuration.  \\noindent (7) We performed a compilation of T Tauri types for both individual components and pairs of the triple/quadruple systems from the available literature in order to study the relative disk evolutions in these systems. With the caution of small number statistics we identified the following trends. The fraction of mixed systems is high, up to $\\sim$\\,65\\,\\%, similar to what is found for binaries. Likewise, the fraction of mixed systems in Taurus is low. Disk replenishment by a circumbinary envelope seems necessary in order to explain the actively accreting close pairs, while the coexistence of close disk-less pairs with other accreting components in some systems argue in favor of no replenishment or a disrupted envelope.  As for what is found for binaries, the disks seems to evolve in isolation at large separations and be affected by tidal truncation at small separations.  \\appendix"
    },
    "0608/astro-ph0608397_arXiv.txt": {
        "abstract": "The Period - metallicity - K band luminosity (PL$_K$) relation for RR Lyrae stars in 15 Galactic globular clusters and in the LMC globular cluster Reticulum has been derived. It is based on accurate near infrared (K) photometry combined with 2MASS and other literature data. The PL$_K$ relation has been calibrated and compared with the previous empirical and theoretical determinations in literature. The zero point of the absolute calibration has been obtained from the K magnitude of RR Lyr whose distance modulus has been measured via trigonometric parallax with HST. Using this relation we obtain a distance modulus to the LMC of $(m-M)_0 = 18.54 \\pm 0.15$ mag, in good agreement with recent determinations based on the analysis of Cepheid variable stars. ",
        "introduction": "RR Lyrae, as well as classical Cepheids, are considered standard candles for estimating stellar distances in the Milky Way and to Local Group galaxies. They are produced in old stellar populations, and hence they provide important information for an understanding of the age, structure and formation of their parent stellar systems. As typical Population II stars, RR Lyrae are abundant in globular clusters and have been the subject of a huge number of studies for more than a century.  Although the pulsation theory explains quite well the connection between most of the involved physical quantities (van Albada \\& Baker 1971; Caputo, Marconi \\& Santolamazza 1998), the dependence of intrinsic luminosity on metal content has been the subject of debate for nearly three decades (see Smith 1995 for a review). Only recently the slope of the luminosity-metallicity relation M$_V$-[Fe/H] seems to be converging towards a value $\\sim$0.20-0.23 that appears to be ``universal'' as supported by the most accurate studies of field RR Lyrae stars in the Milky Way (Fernley et al. 1998; Chaboyer 1999) and in the Large Magellanic Cloud (LMC, Gratton et al. 2004), and globular clusters in M31 (Rich et al. 2005).  Near infrared observations of variable stars present several advantages over optical investigations: a smaller dependence on interstellar extinction and metallicity, a smaller pulsational amplitude and more symmetrical light curves, and hence good mean magnitudes. In the last 20 years the calibration of the RR Lyrae period-IR luminosity relation (PL$_K$) has been the subject of several empirical and theoretical investigations (see Sect. \\ref{method}), but there still remains some degree of uncertainty on the dependence of $M_{K}$ on period and metallicity. To solve the problem of the dependence of $M_{K}$(RR) on these physical quantities, a large sample of RR Lyrae stars is needed, spanning a wide range of [Fe/H], for which accurate K and [Fe/H] measurements are available.  In this paper we present an accurate analysis of the infrared photometric properties of 538 RR Lyrae variables (376 RRab and 162 RRc) in 16 globular clusters (GC) with $-2.15<[Fe/H]<-0.9$. This more than doubles the number of stars used in the previous largest study of this type (Nemec, Linnell Nemec \\& Lutz 1994). By means of such a large database we calibrated the PL$_K$ relation constraining its dependence on period and metallicity on a strictly observational basis.  In \\S 2 we describe the sample of RR Lyrae stars and the used photometric IR datasets. \\S 3 is devoted to the description of the adopted method to calibrate the PL$_K$ relation. In \\S 4 we use the derived PL$_K$ relation to estimate the distance to the calibrator GCs and to a sample of field RR Lyrae stars in the LMC, and compare the results with the previous determinations in literature. Finally, we summarize our results in \\S 5. ",
        "conclusions": "From an accurate analysis of 538 RR Lyrae variables in 16 GCs using infrared (K-band) photometry we derive a PL$_K$ relation based on purely observational constraints. The derived dependences of the K magnitude on period and metallicity are in good agreement with those estimated by previous empirical studies. We confirm that the metallicity coefficient is about 2-3 times smaller than that predicted by theoretical models, as it was found in all previous empirical analyses. The zero point of the calibration has been tied to the trigonometric parallax of RR Lyr measured with HST by Benedict et al. (2002). This calibration has been used to derive the distances to the 16 calibrator GCs considered in this analysis. As a further check, this relation has been applied to RR Lyrae stars in a few central fields in the LMC yielding a distance modulus $(m-M)_{0}=18.54\\pm0.15$, in good agreement with the most recent determinations based on Cepheid variables (Persson et al., 2004; Gieren et al., 2005)."
    },
    "0608/astro-ph0608438_arXiv.txt": {
        "abstract": "We study the brightness variations of galactic red supergiant stars using long-term visual light curves collected by the American Association of Variable Star Observers (AAVSO) over the last century. The full sample contains 48 red semiregular or irregular variable stars, with a mean time-span of observations of 61 years. We determine periods and period variability from analyses of power density spectra and time-frequency distributions. We find two significant periods in 18 stars. Most of these periods fall into two distinct groups, ranging from a few hundred to a few thousand days. Theoretical models imply fundamental, first and possibly second overtone mode pulsations for the shorter periods. Periods greater than 1000 days form a parallel period-luminosity relation that is similar to the Long Secondary Periods of the Asymptotic Giant Branch stars. A number of individual power spectra shows a single mode resolved into multiple peaks under a Lorentzian envelope, which we interpret as evidence for stochastic oscillations, presumably caused by the interplay of convection and pulsations. We find a strong $1/f$ noise component in the power spectra that is remarkably similar in almost all stars of the sample. This behaviour fits the picture of irregular photometric variability caused by large convection cells, analogous to the granulation background seen in the Sun. ",
        "introduction": "Red supergiants (RSGs) are evolved, moderately massive (10--30 M$_\\odot$) He-burning stars. As such, they are key agents of nucleosynthesis and chemical evolution of the Galaxy, and have also long been known for their slow optical variations. This variability is usually attributed to radial pulsations (Stothers 1969, Wood, Bessell \\& Fox 1983, Heger et al.\\ 1997, Guo \\& Li 2002), although irregular variability caused by huge convection cells was also suggested from theory (Schwarzschild 1975, Antia et al.\\ 1984) and observations (e.g. Tuthill, Haniff \\& Baldwin 1997). The presence of oscillations, in principle, offers the possibility of asteroseismology, i.e. measuring frequencies of stars and then identifying the modes of pulsations via comparing the observations to theoretical models. However, the time scales involved in RSGs makes this approach very difficult. Model calculations predict pulsational instability for fundamental and low-order overtone modes (Stothers 1972, Wood, Bessell \\& Fox 1983, Heger et al.\\ 1997, Guo \\& Li 2002), with fundamental periods ranging from 150 to 4,000 days. Consequently, any meaningful period determination may require decades of observations. Here we analyse a homogeneous sample of RSG light curves with a typical time-span of over 22,000 days, which is the longest available observational data for these objects.  Significant interest in RSG pulsations was driven by the discovery of their distinct period-luminosity (P--L) relation. Although the relation is not as tight as for Cepheids or red giant stars, various authors pointed out that because of their great intrinsic brightnesses, RSGs might be useful as  extragalactic distance indicators (e.g. Glass 1979, Feast et al.\\ 1980, Wood \\& Bessell 1985). A recent example is  that of Pierce et al.\\ (2000), who re-calibrated the RSG P--L relation in the  near-infrared and then measured the distance to M101 from 42 RSGs in that galaxy  (Jurcevic et al. 2000). In this respect, galactic RSGs that are members of young O--B associations are particularly important because they can serve as local calibrators of the P--L relation -- provided that reliable periods can be measured. Establishing this zero point is one of the by-products of this paper.  Until recently, there was a serious discrepancy between evolutionary models and the empirical Hertzsprung--Russell diagram of RSGs. In particular, stellar evolutionary models seemed unable to produce RSGs that are as cool or as luminous as observed (Massey 2003, Massey \\& Olsen 2003). A solution was found by Levesque et al.\\ (2005), who derived a new (hotter) temperature scale from optical spectrophotometry for 74 Galactic RSGs, so that most of the discrepancy was removed. In addition, Massey et al. (2005) argued that circumstellar dust around RSGs may account for many magnitudes of extra extinction compared to other stars in the same O--B associations, so that determining fundamental  physical parameters can be quite difficult even for the brightest RSGs. The problem is very well illustrated, for instance, by the peculiar RSG VY~CMa, for which previous radius estimates of up to 2,800 R$_\\odot$ (Smith et al. 2001) were  scaled down to 600 R$_\\odot$ (Massey, Levesque \\& Plez 2006), due to the hotter temperature scale alone. These wildly different parameter values could be tested if pulsation mode identification was made possible, which is another motivation for this work.  In the General Catalogue of Variable Stars (GCVS, Kholopov et al.\\ 1985-1988) there are two categories of variable RSGs: SRc and Lc. SRc type stars are semiregular late-type supergiants with amplitudes of about 1 mag and periods from 30 days  to several thousand days. In contrast, Lc type stars are irregular variable supergiants having visual amplitudes of about 1 mag. As the distinction between {\\it semiregular} and {\\it irregular} behaviour is not well defined in the literature, stars in both categories can possibly reveal important information. So far, the most detailed study of bright galactic RSG variables as a group was that by  Stothers \\& Leung (1971), who discussed periods, luminosities and masses for  22 stars. Their periods were either taken from papers published in the 1950s or (for 8 stars) determined by the authors using visual data. For some stars, such as  $\\alpha$~Ori, $\\alpha$~Her and $\\mu$~Cep, one can find more recent attempts to derive period(s) using various sources of photometric or radial velocity data  (e.g. Mantegazza 1982, Smith, Patten \\& Goldberg 1989, Percy et al.\\ 1996, Brelstaff et al.\\ 1997,  Rinehart et al.\\ 2000). However, the typical time-span of analysed observations rarely exceeded 10--15 years and usually with many gaps, thus the derived periods (or mean cycle lengths) did not always agree.  The American Association of Variable Star Observers (AAVSO) has been collecting visual observations of variable stars for almost a century. The AAVSO database is by far the biggest one of its kind, containing over 12 million individual brightness estimates for about 6,000 variable stars. Of these, roughly 50 stars are reasonably well-observed bright galactic RSGs of the GCVS types SRc and Lc. Using their observational records in the AAVSO database, we have carried out a period and light curve analysis of this group of stars. The main aim was to derive periods for all variables. However, it turned out that in some cases the light curves had long enough time-span to study the long-term behaviour of their brightness fluctuations. Therefore, in addition to determining the dominant time-scales, we also discuss the nature and extent of the irregularities.  The paper is organised as follows. In Sect. 2 we describe the sample selection, data processing and analysis. Sect. 3 contains the results, the newly determined periods; Sect. 4 discusses the multiperiodic nature and a comparison with pulsation model predictions. In Sect. 5 we present evidence for stochastic oscillations in some of the stars and the presence of a strong $1/f$ noise. Sect. 6 briefly summarizes the main findings of this study. ",
        "conclusions": "Red supergiants as variable stars have long been known for their semiregular brightness fluctuations. Using the extremely valuable database of visual observations of the AAVSO, we were able to study the main characteristics of their regular and irregular variations. From a detailed analysis of power spectra and time-frequency distributions, this paper discussed the properties of pulsations and their physical implications.  The sample contains several types of light variations. A few stars (S~Per, VX~Sgr, AH~Sco) have very large peak-to-peak amplitudes that may reach up to 4 magnitudes in the visual. These objects show both the most coherent periodic signals and the highest level of $1/f$ noise in the power spectrum, which might be a sign of a ``superwind'' phase of RSGs that precedes the supernova explosion. We argue that multiple periods found by other studies were sometimes artifacts caused by the extreme noise levels of the systems. A more common type of variability is characterized by two distinct periods, one of a few hundred days and one of a few thousand days. The archetypes of these stars are $\\alpha$~Ori and TV~Gem, both having a period around 400 days and another around 2000 days. The shorter periods can be identified with the radial fundamental or low-order overtone modes of pulsation, while the longer one is very similar to the Long Secondary Periods of AGB stars, whose origin is not known yet, but could be due to binarity, magnetic activity or non-radial $g$-modes (Wood et al.\\ 2004). Besides fundamental pulsation, we also see evidence for first and possibly second overtone modes. Finally, in a few stars we cannot infer any periodicity at all (e.g. W~Cep, AD~Per) and these are the truly irregular variables dominated by the $1/f$ noise in the power spectrum.  The period jitter of the pulsation modes produces in the power spectrum a well-defined Lorentzian envelope. Interpreting this as evidence for stochastically excited and damped oscillations, we measured the mode-lifetime (or damping rate) in red supergiants for the first time. In most stars it is several times longer than the period of pulsations. Since the damping rate depends on the stellar structure and convection properties in a complex way, with many weakly constrained theoretical parameters (e.g. Balmforth 1992) and, moreover, currently there are no theoretical calculations directly applicable to red supergiants, it is impossible to qualify the agreement with theoretical expectations. However, the strong $1/f$ noise component that seems to be ubiquitous in the whole sample, strongly favours the Schwarzschild (1975) mechanism of producing random brightness variations with huge convection cells, analogous to the granulation background seen in the Sun.  Finally, stochastic oscillations discussed in this paper may offer an explanation for the seemingly random behaviour of less luminous red giants stars. For example, random cycle-to-cycle fluctuations of the periods of Mira stars (e.g. Eddington \\& Plakidis 1929, Percy \\& Colivas 1999) are likely to be analogous to the period jitter we found here. It would be helpful to have theoretical models that are specifically designed for these phenomena, because the extensive records of homogeneous visual observations allow us measuring noise properties of red giant and supergiant stars quite accurately. In principle, internal physics could be probed through these accurate measurements provided that realistic models are calculated."
    },
    "0608/astro-ph0608262_arXiv.txt": {
        "abstract": "Three binaries are now established sources of emission at very high energies ($>$10$^{11}$~eV). They are composed of a massive star and a compact object. The emission can be due to the interaction of the relativistic wind from a young ms pulsar with the stellar wind of the companion \\citep{maraschi}, by which rotation-power ends up as non-thermal flux. Variations at VHE energies are explained as due to $\\gamma\\gamma$ absorption and/or changes in shock location along the orbit. Resolved radio emission is due to cooling particles trailing the pulsar. ",
        "introduction": "Thanks to the new generation of Cherenkov telescopes HESS and MAGIC, the past two years have seen a spectacular increase in the number of known sources emitting photons with energies greater than a TeV (see the contributions by Cui, Santangelo and Teshima in this volume). There are now about 40 TeV sources, mostly concentrated in the Galactic plane. Three of the Galactic sources have been associated with binary stars on the basis of their positional coincidence, variability and spectra: \\psrb\\ \\citep{psrb}, \\ls\\ \\citep{lspaper} and \\lsi\\ \\citep{lsipaper}.  \\subsection{The binaries} Those three {\\em gamma-ray binaries} are composed of a massive star (O or Be type) and a compact object, neutron star or black hole, in an eccentric orbit. The orbital periods are $\\approx 4$~days (\\ls), 26~days (\\lsi) and 1237~days (\\psrb). Given the small mass functions, radial velocity measurements have been unable to pin down the nature of the compact object in \\ls\\ and \\lsi. In both cases, a black hole would imply seeing the systems at $i\\approx 30^{\\rm o}$, while a neutron star would imply a more edge-on $i\\approx 60^{\\rm o}$. Regardless, the presence of a luminous early-type star, providing a copious source of seed photons for inverse Compton scattering, is probably instrumental to very high energy (VHE) $\\gamma$-ray emission.  With distances ranging from 1.5-3~kpc, the TeV emission is at a level $\\sim 10^{33-34}$~erg~s$^{-1}$, comparable to the X-ray emission from these sources. Remarkably, all have been detected in the radio, a feature otherwise shared by only a handful of high-mass X-ray binaries. Bar from periodic radio outbursts in \\lsi\\ and \\psrb, very likely related to passage of the compact object through the dense equatorial wind of the B{\\em e} companion (\\ls\\ has a O9V companion -- and no such outburst), the emission from the systems is surprisingly steady up to timescales of years. This contrasts with typical HMXBs. Their spectral energy distributions superpose well with each other, with a rising spectral luminosity from radio to 0.1-1~MeV, a flat spectrum up to $1-10$~GeV and a drop at TeV energies.  \\subsection{What powers the VHE emission?} The high-energy emission could be powered by accretion {\\em via} capture of the intense stellar wind from the companion, but the Bondi-Hoyle rates are low and observational signatures of accretion are lacking. Still, microquasar-type models have been proposed in which the $\\gamma$-rays are emitted in a relativistic jet (see contribution by Paredes in this volume). Resolved compact radio emission in \\ls\\ and \\lsi\\ on scales of 10-100 mas and interpreted as jet emission, constitutes a strong motivation.  On the other hand, the spectral and temporal similarities between the three binaries detected in VHE $\\gamma$-rays hint at a common scenario. This is naturally provided for by \\psrb. In this system, the relativistic wind emitted by the young 48~ms pulsar is contained by the stellar wind. Particle acceleration at the termination shock leads to X-ray synchrotron emission and $\\gamma$-ray inverse Compton emission with the star photons. The measured pulsar spindown power is a modest $8\\cdot 10^{35}$~erg~s$^{-1}$, implying rotation-power is converted into radiation with a high efficiency.  The X-ray spectra have photon index $\\Gamma\\approx 1.5$, harder than those of Crab-like high-luminosity pulsars but comparable to those of pulsar wind nebula (PWN) powered by young pulsars with $\\dot{E}\\sim10^{36}$~erg~s$^{-1}$ \\citep{gotthelf}. X-ray pulses might be detectable underneath the PWN emission if scattering is negligible. However, pulsed radio emission would be strongly absorbed by the dense stellar wind in \\ls\\ and \\lsi, with free-free opacities $\\tau\\ga 10^3$. The pulsed radio emission does vanish in \\psrb\\ around periastron, when the pulsar probes the densest region of the Be wind. At that point, the orbital separation is comparable to the maximum separation in \\lsi, and much larger than in the very compact \\ls.  The advantages of the PWN scenario, discussed below and in \\citet{dubusp}, are that it readily explains the spectrum, level and stability of the emission without requiring additional assumptions from what is already known from plerions, and that it offers a very fruitful, common framework for interpretation. ",
        "conclusions": "Emission from a compact plerion provides a common, coherent framework to understand $\\gamma$-ray binaries. These sources are the short-lived progenitors of accretion-powered HMXBs, which have slowly-rotating spundown pulsars. About 30 such sources are expected in our Galaxy\\footnote{consistent with seeing 3 within 3~kpc} given the present-day HMXB population \\citep{meurs89}.  High-energy emission occurs on small scales. Strong variations in the location of the pulsar wind termination shock (\\lsi, \\psrb) and $\\gamma\\gamma$ absorption of VHE photons on starlight (\\ls) can cause periodic changes at X-rays energies and above. Particles leave the vicinity of the pulsar as they cool, forming a cometary tail. Periodic changes in the aspect of this tail can be resolved in radio on scales $\\sigma c P_{\\rm orb}$, providing a test for the PWN scenario. The radio nebula is steady on larger scales.  Compact plerions enable access to pulsar wind physics on hitherto inaccessible scales. Detailed understanding of the VHE emission, notably in \\ls, can provide important clues as to the shock geometry, magnetic field intensity and acceleration processes around young rotation-powered pulsars."
    },
    "0608/astro-ph0608276_arXiv.txt": {
        "abstract": "We discuss the motivation for high accuracy relativistic gravitational experiments in the Solar System and complementary cosmological tests. We focus our attention on the issue of distinguishing a generic scalar-theory of gravity as the underlying physical theory from the usual general relativistic picture, where one expects the presence of fundamental scalar fields associated, for instance, to inflation, dark matter and dark energy. ",
        "introduction": "\\label{sec:intro}  Present day experimental evidence indicates that gravitational physics is in agreement with Einstein's theory of General Relativity to considerable accuracy; however, there are a number of reasons, theoretical and experimental, to question the theory as the ultimate description of gravity.  On the theoretical side, difficulties arise from various corners, most stemming from the strong gravitational field regime, associated with the existence of spacetime singularities and the difficulty to describe the physics of very strong gravitational fields. Quantization of gravity is a possible way to overcome these obstacles, however, despite the success of modern gauge field theories in describing the electromagnetic, weak, and strong interactions, the path to describe gravity at the quantum level is still to be found. Indeed, our two foundational theories, Quantum Mechanics and General Relativity, are not compatible with each other. Furthermore, in fundamental theories that attempt to include gravity, new long-range forces often arise in addition to the Newtonian inverse-square law. Even at the classical level, and assuming the validity of the Equivalence Principle, Einstein's theory does not provide the most general way to establish the spacetime metric. There are also important reasons to consider additional fields, especially scalar fields. Although the latter appear in unification theories, their inclusion predicts a non-Einsteinian behavior of gravitating systems. These deviations from General Relativity include violations of the Equivalence Principle, modification of large-scale gravitational phenomena, and variation of the fundamental ``constants''. These predictions motivate new searches for very small deviations of relativistic gravity from General Relativity and drive the need for further gravitational experiments in space. These include laser astrometric measurements\\cite{solvang_lator04,Lator01,Texas@Stanford_lator,ESTEC_lator}, high-resolution lunar laser ranging (LLR)\\cite{Murphy_etal_2002} and long range tracking of spacecraft using the formation flight concept, as proposed\\cite{Pioneer} to test the Pioneer anomaly\\cite{Anderson02}. A broader discussion on the motivations to perform fundamental physics experiments in space can be found elsewhere\\cite{Matos04}.  On the experimental front, recent cosmological observations does lead one to conclude that our understanding of the origin and evolution of the Universe based on General Relativity requires that most of the energy content of the Universe resides in the presently unknown dark matter and dark energy components that may permeate much, if not all spacetime. Indeed, recent Cosmic Microwave Background Radiation (CMBR) WMAP three year data\\cite{WMAP3} indicates that our Universe is well described, within the framework of General Relativity, by a flat Robertson-Walker metric, meaning that the energy density of the Universe is fairly close to the critical one, $\\rho_c \\equiv 3H_0^2/8 \\pi G \\simeq 10^{-29} g/cm^3$, where $H_0 \\simeq 73~km~s^{-1}Mpc^{-1}$ is the Hubble expansion parameter at present. Moreover, CMBR, Supernova and large scale structure data are consistent with each other if, in the cosmic budget of energy, dark energy corresponds to about $73\\%$ of the critical density, while dark matter to about $23\\%$ and baryonic matter, the matter that we are made of, to only about $4\\%$. Furthermore, it is generally believed that the ultimate theory that will reconcile Quantum Mechanics and General Relativity will also allow for addressing the cosmological questions related with the origin and destiny of the Universe.  It is our opinion that the crystallization of these fundamental questions is well timed with recent progress in high-precision measurement technologies for physics experiments in space. This puts us in position to realistically address crucial questions, such as the nature of dark energy and dark matter, the existence of intermediate range forces and the ultimate nature of gravity. Furthermore, given the ever increasing practical significance of General Relativity, for spacecraft navigation, time transfer, clock synchronization, weight and length standards, it is just natural to expect that the theory will be regularly tested with ever increasing accuracy. Thus, it seems legitimate to speculate that the present state of physics represents a unique confluence of important challenges in high energy physics and cosmology together with technological advances and access to space, a conjunction that is likely to yield major discoveries.  In what follows we shall address the key issue of distinguishing a generic scalar-theory of gravity, as the underlying fundamental physical theory, from the usual general relativistic picture, where one expects the presence of fundamental scalar fields associated to inflation, dark matter and dark energy. In order to concretely discuss the matter we will consider a fairly general scalar-tensor theory of gravity as an example, and indicate how its main features can be extracted from high-resolution measurements of the parametrized post-Newtonian (PPN) parameters $\\beta$ and $\\gamma$. As is well known, scalar-tensor theories of gravity mimic a plethora of unification models. For instance, the  graviton-dilaton system in string/M-theory can be viewed as an specific scalar-tensor theory of gravity.  Of course, one should to bear in mind that current experimental data shows an impressive agreement with General Relativity\\cite{Will05,BPT06}. Indeed, most stringent bounds arise from the Cassini's 2003 radiometric experiment\\cite{Bertotti}:  \\beq \\gamma - 1 = (2.1 \\pm 2.5) \\times 10^{-5}~~, \\label{eq:gammast3} \\eeq and  \\beq \\beta - 1 = (1.2 \\pm 1.1) \\times 10^{-4} \\label{eq:betast3} \\eeq that arises from limits on the Strong Equivalence violation parameter, $\\eta \\equiv 4\\beta-\\gamma-3$, that are found to be $\\eta=(4.4\\pm4.5)\\times 10^{-4}$, as inferred from LLR measurements\\cite{Williams_Turyshev_Boggs_2004}.  As already mentioned, in cosmology, General Relativity allows for detailed predictions of the nucleosynthesis yields and of the properties of the CMBR, provided one admits the presence of fundamental scalar fields, the inflaton, the quintessence scalar field\\cite{Copeland06} or the generalized Chaplygin gas model underlying scalar field, complex\\cite{Bento02} or real\\cite{Bertolami04a}, to account for the late accelerated expansion of the Universe, and in the case of some candidates for dark matter, self-interacting\\cite{Bento01} or not\\cite{Nunes99}. It is worth remarking the generalized Chaplygin gas model corresponds to a unified model of dark energy and dark matter, based on the equation of state $p = -A/\\rho^{\\alpha}$, where, $p$ is the isotropic pressure, $\\rho$ is the energy density, and $A$ and $\\alpha$ are positive phenomenological constants. Its agreement with observational data has been extensively studied: CMBR\\cite{Bento02}, supernova\\cite{Bertolami04a,Bento05}, gravitational lensing\\cite{Silva03}, gamma-ray bursts\\cite{BertSilva06} and cosmic topology\\cite{Bento06}. A fully consistent picture for structure formation in the context of the generalized Chaplygin gas model remains still an open question\\cite{Bento04}. Another interesting cosmological issue concerns the resemblance of inflation and the late accelerated expansion of the Universe, which has lead to proposals where the inflaton and the quintessence scalar field are related\\cite{Peebles}.  A scalar field with a suitable potential can be also the way to explain the Pioneer anomaly\\cite{Bertolami04b}. It is interesting to point out that scalar fields can affect stellar dynamics and hence, specific measurements of, for instance, the central temperature of stars and their luminosity can allow for setting bounds on scalar field models\\cite{Bertolami05}. ",
        "conclusions": "Let us now review the main points of our discussion. It seems evident that resolving the dichotomy dark energy - dark matter versus gravity will require a concerted effort and a whole new program of dedicated experiments in space.  It is an exciting prospect that dark matter can be directly detected in underground experiments or in the forthcoming generation of colliders. Even though it does not seem feasible to directly test the properties of dark energy, it is not impossible that indirect evidence can be found in laboratory. The most bold proposal suggests the existence of a cutoff frequency of the noise spectrum in Josephson junctions\\cite{Beck}, while a more conventional approach is to investigate the effect that dark energy may have, for instance, on the variation of the electromagnetic coupling\\cite{Olive}. It follows that the characterization of dark energy and dark matter will most likely be achieved via cosmological observations, most of them to be carried out by space-borne experiments. These encompass a large array of phenomena such as supernovae, gamma-ray bursts, gravitational lensing, cosmic shear, etc. The result of these observations will also provide increasingly detailed information on the adequacy of General Relativity at cosmological scales. It is quite exciting that existing supernova data\\cite{Riess2004} together with latest CMBR data\\cite{WMAP3} and the recently discovered baryon acoustic oscillations\\cite{Eisenstein} are sufficiently constraining to virtually rule out\\cite{Maartens}, for instance, most of the braneworld inspired gravity models put forward to account for the accelerated expansion of the Universe. The prospect of testing some of these models through the study of the orbital motion of planets in the Solar System has also been recently discussed\\cite{Iorio}.  We have seen how stands the situation in what concerns scalar-tensor gravity models. Relevant results are expect within a decade from the observation of binary pulsar systems. To further test General Relativity and examine the implications of its contending theories or extensions (scalar-tensor theories, braneworld models, string inspired models, etc.) a new program of gravity experiments in space is clearly needed. We have discussed how LLR can be used to improve the knowledge of relativistic gravity and pointed out how the LATOR mission and a mission to test the Pioneer anomaly can play a key role in the search for evidence of a remnant scalar field in the Solar System, to identify new forces with ranges of a few decades of AU and, of course, to resolve the Pioneer anomaly puzzle. It is relevant to point out that the latter type of mission, besides its technological appeal, can also to used to gather information about the vicinity of the Solar System as well as to set relevant upper bounds on environmental parameters such as the density of interplanetary dust and dark matter\\cite{BertVieira}.  \\vspace{0.4cm}  \\noindent {\\bf Acknowledgments}  \\vspace{0.2cm}  \\noindent It is a pleasure to thank the members of the Pioneer Science Team for the countless discussions on the questions related with this contribution. I am particularly in debt to Jorge P\\'aramos, Slava Turyshev, Serge Reynaud, Clovis de Matos, Pierre Toubul, Ulrich Johann and Claus L\\\"ammerzahl for their insights and suggestions.   \\vspace{0.2cm}"
    },
    "0608/astro-ph0608040_arXiv.txt": {
        "abstract": "The Kennicutt--Schmidt law relates the face-on star formation rate (SFR) per unit area with the face-on gaseous column density in nearby galaxies.  Applying this relation to damped {\\lya} absorption systems (DLAs) of neutral hydrogen column density {\\nh} $> 1.6\\times 10^{21}$ cm$^{-2}$ leads to an estimate that three percent of the sky should be covered with extended sources brighter than $\\mu_{V}\\approx 28.4$ mag arcsec$^{-2}$, if DLAs at redshift $z$=[2.5,3.5] undergo {\\em in situ} star formation.  We test this hypothesis by searching the Hubble Ultra Deep Field (UDF) F606W image for low surface-brightness features of angular sizes, ranging between {\\thetdla}=0.25\\arcsec\\ and 4.0\\arcsec.  At $z=2.5-3.5$, the observed F606W fluxes correspond to roughly rest-frame 1500 \\AA\\ and the angular sizes correspond to predicted disk diameters of $d_{\\rm dla}=2-31$ kpc.  After convolving the F606W image with smoothing kernels of angular diameters {\\thetkern}={\\thetdla}, we find the number of detected objects to decrease rapidly to zero at {\\thetkern}$>$1 {\\arcsec}. Our search yields  upper limits on the comoving SFR densities that are between factors of 30 and 100 lower than predictions, suggesting a reduction by more than a factor of 10 in star formation efficiency at $z\\sim 3$.  We consider several mechanisms that could reduce star formation efficiency at high redshift. We find that the cosmological increase with redshift of the critical surface density for the Toomre instability may be sufficient to suppress star formation to the levels implied by the UDF observations. However, the uncertainties are such that Toomre instabilities may still exist. In that case star formation at column densities less than 10$^{22}$ cm$^{-2}$ may be suppressed by the low molecular content of the DLA gas. The upper limits on {\\em in situ} star formation reduce the predicted metallicities at $z\\sim 3$ to be significantly lower than observed, and reduce the heat input in the gas to be substantially lower than the inferred cooling rates.  In contrast, the radiative output from compact Lyman Break Galaxies (LBGs) with $R$ $<$ 27 is sufficient to balance the comoving cooling rate.  This leads us to posit that a significant fraction of the DLA population are hosts to more compact regions of active star formation, which may be the sources of metal enrichment for these DLAs. Such regions are likely to be LBGs. ",
        "introduction": "A principal goal of galaxy formation theory is to understand how stars form from gas.  This condensation process is fundamental for determining the star-formation history of a galaxy, the spatial distributions of its stellar populations, and its chemical evolution. Yet it is still not understood.  On small scales, the physical processes are complex and difficult to calculate (e.g. Shu {\\etal} 1987; Krumholz \\& McKee 2005), while on cosmological scales, realistic simulations of star formation are in addition constrained by limits on numerical resolution (e.g. Cen {\\etal} 2003; Nagamine {\\etal} 2004a, 2004b).  On the other hand considerable progress has been made in the observational sector. Recent multi-color imaging and spectroscopic surveys have succeeded in tracing starburst galaxies out to redshifts as large as six (Giavalisco {\\etal} 2004; Bouwens {\\etal} 2004). The majority of galaxies found in this way are the Lyman Break Galaxies (i.e. LBGs; e.g. Steidel {\\etal} 2003).  These are compact (half-light {\\em diameters} $\\sim$ 4 kpc) star-forming galaxies with mean star formation rates SFR $\\sim$ 40 {\\smpy}, after extinction corrections are applied (Shapley {\\etal} 2003). The LBGs contribute a SFR per unit comoving volume {\\rhodotz} (Steidel {\\etal} 1999; Giavalisco {\\etal} 2004) that in the redshift interval $z$ =[6,2] would consume a mass per unit comoving volume of cold neutral gas equivalent to about 10 $\\%$ of the mass content of visible stars in modern galaxies, i.e., 0.1{$\\Omega_{*}(z=0)$}. As a result, reservoirs of neutral gas at $z \\ \\ge$ 2 may be required to fuel these star-forming objects.  The purpose of this paper is to search for star formation in spatially extended regions that serve as such neutral-gas reservoirs. We focus on the damped {\\lya} systems (hereafter DLAs), the population of quasar absorption systems with {\\nh} {$\\ge$} {\\thresh}, where {\\nh} denotes observed H I column density (for a review see Wolfe, Gawiser, \\& Prochaska 2005 [hereafter WGP05]).  At $z$ $\\sim$ 3 this column-density threshold guarantees gas neutrality in most cases, which distinguishes DLAs from the {\\lya} forest and all other classes of absorption systems in which {\\nh} $<$ {\\thresh} and the gas is more than 50 $\\%$ ionized. The neutrality of the gas takes on added significance when it is realized that DLAs (1) dominate the neutral-gas content of the Universe at $z$ =[0,5] and (2) at $z$ $\\approx$ 3.5 contain sufficient gas to account for 0.5{$\\Omega_{*}(z=0)$}. Furthermore surveys for DLAs reveal a large area-covering factor for neutral gas with {\\nh} {$\\ge$} {\\thresh} (Wolfe {\\etal} 1995; Prochaska, Herbert-Fort, \\& Wolfe 2005 [PHW05]).  In the redshift interval $z$=[2.5,3.5], a redshift range in which the SFR history of galaxies is well determined, DLAs cover one third of the sky, which leads to a striking",
        "conclusions": "We searched for {\\em in situ} star formation in the neutral gas comprising DLAs by looking for low surface-brightness emission from spatially extended objects in the Hubble Ultra Deep Field F606W image.  The search was designed to detect objects in the redshift interval $z=[2.5,3.5]$ with linear {\\em diameters} ranging between 1.9 kpc and 31 kpc, which encompasses most model predictions published so far.  The search was optimized using a smoothing kernel to match the predicted size of the DLAs, ranging between an angular diameter of {\\thetdla}= 0.25{\\arcsec} and 4.0{\\arcsec}. The results are summarized as follows:  {$\\bullet$} After eliminating compact objects with high surface brightness, we found the number of extended sources to decrease rapidly with increasing {\\thetkern}: the implication is that extended sources of low surface-brightness emission rarely occur in the UDF (see Fig.~{\\ref{fig:dlanumvthet}}). Because photometric redshifts are unavailable, these findings lead to upper limits on the  comoving density of objects within the UDF with $z$=[2.5,3.5].    {$\\bullet$} The upper limits on comoving density lead to upper limits on comoving SFR density, which are between factors of 30 to 100 lower than predicted by the Kennicutt-Schmidt law.  As a result, the projected SFRs per unit area must be at least a factor 10 lower than the rates predicted by the distribution of neutral-gas column densities in DLAs. The upper limits are  quite general as they apply to objects with angular diameters encompassing the sizes of DLAs in {\\lcdm} and all other models suggested so far.  {$\\bullet$} We suggest that lower SFR efficiencies may partly result from an increase with redshift of the critical surface density for Toomre instability. Such an increase is due to cosmological evolution alone. This effect causes the DLA gas to be Toomre stable, which could remove the discrepancy between the comoving SFR density predicted by the Kennicutt-Schmidt law and our empirical upper limits. However, the uncertainties are such that gravitational instabilities may be present in DLAs. In that case we suggest that the low molecular content of the gravitationally bound clouds with $N$ $<$ 10$^{22}$ cm$^{-2}$  prevents further cooling and thus inhibits star formation. The low molecular content of the gas may be related to the minimum column density required for molecule formation (Krumholz \\& McKee 2005; Blitz \\& Rosolowsky 2006) since molecule formation in objects with low dust content and high FUV radiation fields such as found in DLAs could require values of {\\nh} exceeding 10$^{22}$ cm$^{-2}$ for shielding against photo-dissociating radiation. The low area covering factor of such high column densities may be the reason for the low molecular content measured in DLAs.  {$\\bullet$} The upper limits on comoving SFR density  reduce the rate of metal production predicted in DLAs. The resultant upper limits on metallicity at $z$ $\\sim$ 3 are lower than the observed metallicities by a factor of three. While this removes the previous problem of metal overproduction (by a factor of ten)  in DLAs, it leads to a new problem of metal underproduction.  {$\\bullet$} The upper limits on comoving SFR density result in a comoving grain photoelectric heating rate that is significantly lower than the comoving cooling rate inferred from [C II] 158 {\\micron} emission rates for about half the DLA sample. Therefore, external sources of heating are required.  We argue that the observed comoving SFRs of LBGs with $R$ $<$ 27 and $z$=[2.5,3.5] indicate a heat input sufficient to balance cooling.  {$\\bullet$} We suggest a scenario in which half the DLA population with significant [C II] 158 {\\micron} cooling is powered by centrally located, compact LBGs. We speculate that DLA metal enrichment may  be due to  P-cygni outflow from the metal-enriched LBGs. The half of the DLA population without significant [C II] 158 {\\micron} cooling is presumably heated by background radiation alone. Without an LBG, the source of metals for this DLA population is unclear.    Two  possible caveats come to mind. First, since our limits do not apply to emission from DLAs with {\\thetdla} $<$ 0.25 {\\arcsec}, it is possible that for most DLAs at $z$ $\\sim$ 3, $d_{dla}$ is less than 1.9 kpc and the SFRs follow the Kennicutt-Schmidt relation. However, as discussed in $\\S$ 3.1 all models suggested so far predict that the bulk of DLAs will have diameters larger than 1.9 kpc.  The second caveat stems from the possibility that {\\em in situ} star formation at the Kennicutt-Schmidt rate occurs in DLA gas associated with high surface-brightness objects, since such objects were automatically excluded from the survey. We tested this hypothesis by searching for low surface-brightness emission surrounding high surface-brightness objects. To enhance the sensitivity of our search we smoothed the regions in the F606W image surrounding selected targets using kernels with 0.5 {\\arcsec} $<$ {\\thetkern} $<$ 2.0 {\\arcsec}. In all cases the convolved images failed to show evidence for outlying emission with fluxes comparable to the central object. But it is difficult to draw conclusions from these data since the redshifts of all but two galaxies, at $z$ = 2.593 (Szokoly et al.\\ 2004) and $z=3.797$ (Vanzella et al.\\ 2006), respectively, are unknown. We found these two galaxies to consist of a compact core (FWHM = 0.3 {\\arcsec}) surrounded by diffusion emission, which exhibited maximum contrast with respect to sky when  {\\thetkern} = 0.5 {\\arcsec}. In the case of the $z$=2.593 galaxy the magnitude of the diffuse emission, $V$=30, while the central core  has $V$ = 25; i.e., the regions surrounding this star-forming galaxy contributes less than 1 $\\%$ of the total SFR. Because {\\rhodot} due to compact star forming galaxies is comparable to {\\rhodot} predicted by the Kennicutt-Schmidt relation for DLAs, we would conclude that neutral gas surrounding star-forming objects does not form stars according to the Kennicutt-Schmidt relation if this galaxy is typical of its class.          Finally we compare our empirical limits with previous analyses that apply the local Kennicutt--Schmidt law for evaluating {\\rhodot} in DLAs. First, Lanzetta {\\etal} (2002) measured the frequency distribution of projected SFR per unit area, $h$({\\ps}), for galaxies in the HDF in the redshift interval $z$=[0,10]. Their measurements at $z$ $\\sim$ 3 were sensitive to {\\ps} $\\ge$ 10$^{-1.5}$ {\\smpykpc}. To determine $h$({\\ps}) at lower values of {\\ps}, they applied the Kennicutt--Schmidt law to the H I column-density distribution function {\\fNX} by using the relation $h$({\\ps})$d${\\ps}=($H_{0}/c$){\\fNX}$dN$. They then used Eq.~{\\ref{eq:rhostardotN}} in the spherically symmetric limit to determine {\\rhodot}. These authors concluded that at $z\\approx 3$ the bulk of SFR density resides in regions of 0.1 {\\msolar} yr$^{-1}$ kpc$^{-2}$, where the majority of LBG's are found.  They argued that DLAs of {\\nh} $<5\\times 10^{21}$ cm$^{-2}$ contribute an  insignificant fraction of the global SFR density. This is inconsistent with the current form of {\\fNX} which when extrapolated to {\\nh} $>>$ 2{$\\times$}10$^{22}$ cm$^{-2}$, to include HSB objects such as LBGs, predicts that more than 50 $\\%$ of {\\rhodot} is contributed by DLAs with {\\nh} $<5\\times 10^{21}$ cm$^{-2}$, even when $K$=$K_{Kenn}$/2.8 as suggested by Lanzetta {\\etal} (2002). The discrepancy stems from the different expressions used to compute {\\fNX}: Lanzetta {\\etal} (2002) used  the then available expression for {\\fNX} with slope, {$\\alpha_3$}=$-$1.4,  that is flatter than $\\alpha_{3}$=$-$2.0 derived for the more accurate SDSS function (PHW05) used here. Because we have shown that the DLA  contribution to {\\rhodot} cannot be this large, the newer data do not support the suggestion of Lanzetta {\\etal} (2002) that SFRs predicted by application of the Kennicutt-Schmidt law to DLAs overlap  those derived directly from optical emission. On the other hand, their analysis is consistent with our interpretation that LBGs provide the bulk of the star formation at $z$ $\\sim$ 3. Second, Hopkins {\\etal} (2005) combined an updated version of {\\fNX} with the Kennicutt-Schmidt relation to predict {\\rhodot} in the spherically symmetric limit for the redshift interval $z$=[2.5,3.5]. These authors concluded that DLAs contribute roughly 50\\% of the observed SFR density at $z\\sim 3$. Again, the predicted {\\rhodot} is significantly higher than our empirical upper limits, and thus their models are ruled out by the UDF data. The new upper limits on {\\rhodot} are also inconsistent with the predictions of the uniform-disk model of WPG03 and WGP03. Their models do not make use of the Kennicutt- Schmidt law, but rather use the strength of {\\ciis} absorption to infer the level of {\\em in situ} star formation throughout DLAs modeled as uniform disks. Because the predicted average projected SFR, $<${\\ps}$>$=10$^{-2}$ {\\smpykpc}, is above our detection threshold, the factor of 100 disagreement between the predicted {\\rhodot} and our upper limits rules out this model. On the other hand, the new upper limits are consistent with the ``bulge model'' suggested by WGP03. In this model, star formation does not occur in the gas detected in absorption, but rather arises in a centrally located compact region, which heats the surrounding disk gas by emitting FUV radiation. The rough agreement between the comoving heating rate supplied by LBGs and the comoving cooling rate of DLAs suggests the bulges are in fact LBGs."
    },
    "0608/astro-ph0608089_arXiv.txt": {
        "abstract": "We report the discovery with the RXTE/PCA of twin kHz QPOs in the peculiar X-ray binary Circinus X-1. At eleven different epochs we observed two simultaneous kHz QPOs with centroid frequencies of up to $\\sim$225 and $\\sim$500~Hz and significances of up to 6.3 and 5.5$\\sigma$ for the lower and upper kHz QPO, respectively.  The frequency range over which the twin kHz QPOs are seen is for the most part well below that of other sources, and extends down to centroid frequencies of $\\sim$56~Hz and $\\sim$230~Hz respectively for the lower and simultaneously observed upper kHz QPO.  The strongly variable QPO frequencies and their observed correlations clearly indicate that the twin peaks are the kHz QPOs such as typically seen from low-magnetic field neutron stars, and not black-hole high-frequency QPOs.  This confirms that Cir~X-1 is a neutron star, as suspected since the detection of Type-I X-ray bursts from the field of the source 20 years ago.  The kHz QPO peak separation varies over the largest range yet seen, $\\sim$175--340~Hz, and {\\it increases} as a function of kHz QPO frequency.  This is contrary to what has been observed in other sources but in good agreement with predictions of the relativistic precession model and Alfv\\'en wave models at these low QPO frequencies.  Current beat-frequency models would require further modification in order to accommodate this. A total of 67 observations showed only a lower kHz QPO; this QPO can be followed down all the way to a centroid frequency of $\\sim$12~Hz ($Q\\sim$0.5). Two observations showed only an upper kHz QPO near 450~Hz. In addition, a strong low-frequency quasi-periodic oscillation (LF QPO) is seen in the 1--30~Hz range, as well as two components above and below this QPO.  The frequency-frequency correlations between the kHz QPOs and the LF QPO are in good agreement with those found previously in Z sources, confirming that Cir~X-1 may be a peculiar Z source, and extend these correlations to frequencies a factor 1.5--2.3 lower.  We suggest that the low frequency range over which the kHz QPOs occur in Cir~X-1 and to a lesser extent in (other) Z sources, might be due to a stronger radial accretion flow relative to the disk flow than in other kHz QPO sources, possibly related to the nature of the companion star. ",
        "introduction": "Low-mass X-ray binaries (LMXBs) often exhibit quasi-periodic oscillations (QPOs) in their X-ray flux (\\citealt{review} for a review).  The properties of these QPOs vary as a function of the location of the source in the track it traces out in X-ray color-color and hardness-intensity diagrams (CD and HIDs).  Among the low magnetic-field neutron-star systems, systematic differences in the CD/HID tracks and in the associated rapid X-ray variability allow to distinguish two subclasses: Z and atoll sources, named after the shapes of their tracks \\citep{hk89}. Both in Z sources and in atoll sources, the QPO frequency tends to increase in the same direction as the mass accretion rate is inferred to be increasing, at least on short time scales (see \\citealt{parlines}).  In about 20 LMXBs containing a low-magnetic field neutron star, two kHz QPOs have been seen simultaneously (twin kHz QPOs) with a centroid frequency separation $\\Delta\\nu$ of a few hundred~Hz that decreases typically by a few tens of~Hz when the QPO frequencies themselves increase by hundreds of~Hz (roughly, from 200 to 900~Hz and from 500 to 1200~Hz, respectively).  If the spin frequency of the neutron star, $\\nu_s$, is known, then $\\Delta\\nu$ has always been found to be either roughly equal to $\\nu_s$, or to $\\nu_s/2$, depending on the source. The frequencies of the twin kHz QPOs (called the lower and the upper kHz QPO, or $L_\\ell$ and $L_u$, in order of increasing frequency) are tightly correlated to each other, following a relation that is the same across all sources \\citep{psalt98,pbvdk,steve}.  Both kHz QPO frequencies also (e.g., \\citealt{pbvdk}) correlate with the frequencies of other variability components such as, in atoll sources, the low-frequency quasi-periodic oscillation (LF QPO; a sharp QPO peak in the 1--60~Hz range designated $L_{LF}$ -- cf. \\citealt{steve} for some of the terminology) and the break component (a broad noise component with a frequency in the 0.01--20~Hz range designated $L_b$).  The components most similar to these in the Z sources (and suspected to be physically the same) are the horizontal branch oscillation (HBO) and the low frequency noise (LFN), but the correlations of the frequencies of these components with those of the kHz QPOs differ somewhat \\citep{straaten}.  Twin kHz QPOs are a neutron-star signature.  In some LMXBs containing a black hole, high-frequency quasi-periodic oscillations (HF QPOs) have been observed. These are different from kHz QPOs: they tend to recur near fixed frequencies that are different in each source, and when more than one HF QPO occurs in the same source the QPOs usually appear at different times and at frequencies often roughly consistent with a fixed 2:3 ratio \\citep{stroh,mccrem}.  Numerous models have been proposed to explain the nature of the kHz QPOs and describe their correlations.  Some refer to a blob orbiting inside the accretion disk, either self-luminous or scattering the X-rays, or producing a hot spot on the neutron-star surface by the infall of matter from the blob while it is orbiting at the inner edge of the disk.  These models produce a rotating pattern of radiation causing a so-called 'beaming' modulation of the observed flux.  The orbital frequency of a particle in an equatorial circular orbit around a point mass with mass $M$ and angular momentum $J$ (approximating the spacetime metric around the neutron star by Kerr spacetime) is $\\nu_{\\phi}=\\sqrt{GM/(4\\pi^2r^3)} / (1+j(r_g/r)^{3/2})$, where $r_g\\equiv GM/c^2$ and $j\\equiv Jc/GM^2$. The corresponding radial epicyclic frequency in a slightly eccentric orbit and for negligible frame dragging ($j\\approx0$) is $\\nu_r\\approx\\nu_{\\phi}\\sqrt{1-6r_g/r}$.  In the relativistic precession model \\citep{stella} as well as in the sonic-point spin-resonance model \\citep{cole} the upper kHz QPO $L_u$ is identified with $\\nu_{\\phi}$. The lower kHz QPO $L_\\ell$ is interpreted in the first case as due to the periastron precession frequency $\\nu_{\\phi}-\\nu_r$, and in the second case as a beat between $\\nu_{\\phi}$ at the sonic radius and the spin frequency (at approximately $\\nu_{\\phi}-\\nu_{spin}$ or $\\nu_{\\phi}-\\nu_{spin}/2$ depending on conditions at the spin-resonance radius where the beat emission is generated). \\citet{23,1235} suggest that QPOs occurring in neutron stars are similar to those seen in black holes and result from resonance frequencies in the disk, but in neutron stars are tuned away from the fixed frequencies seen in black holes by additional physical effects (but see also \\citealt{distribution}).  Models related to orbital motion are in general more predictive of the frequencies of the QPOs than of their amplitudes and coherences, as constructing a model for the modulation mechanism or describing the damping of the oscillations is more difficult than explaining these frequencies as a (combination of) general-relativistic orbital and epicyclic frequencies.  A maximum frequency according to such models would be the orbital frequency at the innermost (marginally) stable circular orbit (ISCO) of the compact object $\\nu_{ISCO}\\approx (c^3/2\\pi6^{3/2}GM)(1+0.75j)$ for Kerr spacetime or, under certain conditions, the orbital frequency at the marginally bound orbital radius.  Both these frequencies scale inversely with mass.  For an extensive review on timing features of LMXBs and QPO models see \\citet{review}.  Circinus X-1 is a galactic X-ray source located in the galactic plane at a distance that has variously been reported to lie in the range 4--12 kpc \\citep{austr,glass,jonknele,iaria}.  The system is thought to have a 17-d highly eccentric ($e\\sim0.8$) orbit leading to periodically enhanced mass transfer at periastron \\citep{murdin}.  From optical and infrared observations a sub-giant companion star of 3--5~$M_{\\sun}$ was suggested \\citep{mass}. More recently, \\citet{jonker06} found evidence for supergiant characteristics in the optical spectrum of the system which might indicate an even higher companion mass. So, Cir~X-1 is formally a high-mass X-ray binary.  However, its X-ray phenomenology corresponds to that of a low magnetic-field compact object such as more commonly found in LMXBs.  The strong radio emission \\citep[see, e.g.,][]{radio}, highly relativistic jets \\citep{jets}, hard X-ray emission \\citep{iaria01,ding} and very strong X-ray variability \\citep{jones, samimi, atoll}, resemble properties generally more typical of black holes.  However, type-I X-ray bursts seen in 1986 from the field of the source and most likely originating in Cir~X-1 itself \\citep{flashes} classified it as a likely neutron star, with later spectroscopic analysis \\citep{tueb,gierl} supporting this at least at a phenomenological level.  No bursts were reported from the source since 1986.  Low-frequency noise, a clear low-frequency QPO $L_{LF}$ and a broad feature at 100--200~Hz were seen with EXOSAT \\citep{tennant87}. At low flux levels Cir~X-1 exhibited atoll-source behavior \\citep{atoll} including strong band-limited noise, while subsequent observations at higher flux on occasion, and in addition to a wide variety of non-Z-like patterns, showed a Z-source-like track in the HID \\citep{shirey}.  On the horizontal branch of this track, \\citet{shireyI,shireyII} found the LF QPO to move between 1.3~Hz and 32 Hz, in strong correlation with the break frequency at lower frequencies and the frequency of the broad high-frequency feature, now seen to range from 20--100 Hz. \\citet{pbvdk} interpreted this broad peak as a low-frequency, low-coherence version of the lower kHz QPO $L_\\ell$ as seen in Z and atoll sources.  They showed that its frequency fits on the relation seen in those sources between the frequencies of $L_\\ell$ and $L_{LF}$, extending it down to an $L_\\ell$ centroid frequency of $\\sim$20~Hz and, interestingly, providing a link with black-hole phenomenology.  No twin kHz QPOs were reported from Cir~X-1 up to now.  In this paper, we concentrate on the high-frequency part of the power spectra of Cir~X-1 and perform an extensive survey for kHz QPOs. We find that twin kHz QPOs do in fact occur in Cir~X-1.  We describe our methodology in Section~\\ref{method}, present our results in Section~\\ref{results} and comment on consequences for phenomenological systematics and theoretical models in Section~\\ref{dis}. ",
        "conclusions": "\\label{dis} We have discovered twin kHz QPOs in Cir~X-1 in 11 separate RXTE/PCA observations.  Among these 11 observations the centroid frequency of the upper kHz QPO ranged from 229$\\pm$18 to 505$\\pm$51~Hz and that of the lower kHz QPO from 56.1$\\pm$1.3 to 226$\\pm$18~Hz.  The lowest-frequency twin kHz QPOs observed so far, in the accreting millisecond pulsar XTE~J1807$-$294, had $\\sim$127 and $\\sim$353~Hz centroid frequencies \\citep{linares}; our Cir~X-1 data extend to well below this. A total of 80 observations showed at least one kHz QPO, with centroid frequencies extending down to as low as 12.5~Hz.  The centroid frequency separation of the twin kHz QPOs varies by a factor of two, spanning a range, 173$\\pm$18 to 340$\\pm$64~Hz, similar to that covered by all previous twin kHz QPO detections in other sources taken together.  In order to further compare the kHz QPOs in Cir~X-1 with those in other sources we examine the locus of these QPOs in two frequency-frequency correlation diagrams that have been considered in previous works.  In Fig.~\\ref{pbk} we plot $\\nu_u^0$ and $\\nu_{LF}^0$ vs. $\\nu_\\ell^0$ in the \\citet{pbvdk} diagram containing data for both neutron stars and black holes.  The lower kHz QPO points of Cir~X-1 had previously been used by these authors to suggest a link between the lower kHz QPO in Z and atoll sources above a $\\nu_\\ell^0$ of 200~Hz, and a broad feature in black holes and weak X-ray burst sources below a $\\nu_\\ell^0$ of 10~Hz (both plotted in Fig.~\\ref{pbk} using {\\it open circles}).  This latter broad feature was dubbed $L_{\\ell ow}$ by \\citet{steve}, who based on the relations of amplitude and coherence with frequency, and on systematic shifts in the frequency-frequency correlations observed in some millisecond pulsars questioned the identification with $L_\\ell$ \\citep{straaten, steve}.  Our identification of the twin kHz QPOs in Cir~X-1 confirms that at least in Cir~X-1 this feature is indeed the lower kHz QPO as proposed by \\citet{pbvdk}.  Our newly discovered upper kHz QPO points smoothly extend the $\\nu_u^0$ vs. $\\nu_\\ell^0$ relation to frequencies a factor 4 lower than originally reported (\\citealt{pbvdk}; upper track in Fig.~\\ref{pbk}).  The best-fit power law to all points including Cir~X-1 is $\\nu_u^0=(23\\pm3) \\times {\\nu_\\ell^0}^{(0.57\\pm0.02)}$~Hz ($\\chi^2/dof=20$ for 82 degrees of freedom).  As can be seen in Fig.~\\ref{pbk}, this relation provides an excellent match to Cir~X-1 as well as to the data on the other sources.  In Fig.~\\ref{vstrfig} we plot the characteristic frequencies of all other power spectral components in the 11 twin kHz QPO power spectra vs. $\\nu_u$ in a diagram originating with \\citet{multilor} containing data for atoll sources, weak burst sources and Z sources. In this diagram it can be seen that the $\\nu_\\ell$ vs. $\\nu_u$ relation, as smoothly extended with our Cir~X-1 points, does {\\it not} connect to the $\\nu_{\\ell ow}$ vs. $\\nu_u$ relation, supporting van Straaten et al.'s suspicion that $L_\\ell$ and $L_{\\ell ow}$ may be different components.  Further work on the other sources contributing to these correlations is required to definitively resolve this issue; this is beyond the scope of this paper.  The large range over which the twin kHz QPO frequencies vary, the strong deviation from a 2:3 ratio and the frequent simultaneous appearance of the two peaks are very different from what is seen in black hole high-frequency QPOs.  Instead, the way in which the kHz QPO frequencies correlate to each other and to the frequencies of other variability components is consistent with what is seen in neutron stars.  These observations therefore support the identification of Cir~X-1 as a neutron star (cf. Section~1).  In Fig.~\\ref{vstrfig} it can be seen that the LF QPO in Cir~X-1 fits well on the $L_h$ track in this diagram, but the low-frequency noise or break component does not match well to the $L_b$ track.  Cir~X-1 shares both these characteristics with the Z sources \\citep{straaten}, giving further support to the proposal of \\citet{shirey} that Cir~X-1 is a (peculiar) Z source.  The low coherence, relative weakness, and low frequency range of our twin kHz QPOs as well as the detection of a LF QPO harmonic (such as also commonly seen in Z source HBOs) are also in accordance with a Z source classification.  However, the band-limited noise can be much stronger \\citep{atoll}, and the kHz QPOs can reach much lower frequencies in Cir~X-1 than has been observed in any (other) Z source.  The extended range of kHz QPO frequencies in Cir~X-1 allows unique tests of theoretical models.  Perhaps most strikingly, the increase in the frequency separation $\\Delta\\nu$ with upper kHz QPO centroid frequency $\\nu_u^0$, which has only been observed previously, somewhat marginally, in GX~17+2 \\citep{homan} and on a single occasion in 4U~1728$-$34 \\citep{migliari}, is in accordance with a prediction of the relativistic precession model \\citep{stella} for the behavior of $\\Delta\\nu$ at low QPO frequency.  In Fig.~\\ref{rel-prec} we show a plot of the centroid frequency separation $\\Delta\\nu\\equiv\\nu_u^0-\\nu_\\ell^0$ vs. $\\nu_u^0$. From an F-test for additional term we find that in our data, $\\Delta\\nu$ is inconsistent with being constant at the 5.2$\\sigma$ level; a linear fit yields a line with slope 0.6$\\pm$0.1.  The relation predicted by the precession model for negligible frame dragging, $\\Delta\\nu=\\nu_u^0\\sqrt{1-6\\left(2\\pi GM\\nu_u^0\\right)^{2/3}/c^2}$ describes our data points well ({\\it drawn curve}).  The best fit to just our Cir~X-1 points ($\\chi^2/dof=0.2$) gives a mass for the neutron star of $M=(2.2\\pm0.3) M_{\\sun}$.  This is high, but agrees with inferences drawn from this model for other sources \\citep{stella}, and indeed the best-fit curve passes through the cloud of scattered points corresponding to all other sources (which of course can be expected to have different masses).  Note, however, that as already discussed by \\citet{stella}, for some other individual sources no exact match to this model is possible without introducing additional physical assumptions.  We further note, that Alfv\\'en wave oscillation models \\citep{ZhangCM, rezania, ZhangCM2} which also predict $\\Delta\\nu$ to increase at low $\\nu_u^0$, can fit our data as well.   As is evident from the data presented in the rightmost column of Table~\\ref{khz-table}, while the twin kHz QPO frequencies are clearly inconsistent with a 2:3 ratio, a 1:3 ratio is closer, although still formally inconsistent with the data at 99\\% confidence.  A 1:3 ratio corresponds to one of the resonances that have been considered by \\citet{1235}.  However, since that ratio would arise from a combination frequency of the main 2:3 resonance, the weakness of the main resonance (which is not significantly detected) would in such an interpretation remain unexplained.  Beat-frequency models, based on a constant frequency separation $\\Delta\\nu$ related to the spin frequency, would have difficulties explaining the factor 2 change in frequency separation of the kHz QPOs of Cir~X-1.  The modified beat-frequency model proposed by \\citet{lambmil} to explain the modest {\\it decrease} of $\\Delta\\nu$ with QPO frequency previously observed in a number of sources would need further modification to explain the considerable {\\it increase} that we observe instead. The range of upper kHz QPO frequencies we observe (229--505~Hz) is nearly entirely below that expected in the sonic-point beat-frequency model (500--1200~Hz; \\citealt{coleI}).  Although our discovery of twin kHz QPOs in Cir~X-1, as well as the observed frequency-frequency correlations, put this system on a footing with the neutron-star LMXBs, the fact that frequency range over which the upper kHz QPOs are observed is mostly well below that of other sources further confirms the exceptional nature of the source.  These low frequencies might be taken to indicate the presence of a massive neutron star, if for example the upper kHz QPO would derive from the orbital frequency at some radius set by general relativity, scaling like $1/M$.  However, when taken at face value the fact that the observed frequencies are strongly variable makes this interpretation unlikely.  The possibility remains that the highest kHz QPO frequencies in a source derive through general relativity from $M$, and the observed range of frequencies in turn derives from this maximum.  Most proposed models interpret the strong variations of the frequencies as due to variations in the inner radius of the Keplerian disk flow.  This radius could be set by a number of different mechanisms, such as electromagnetic stresses (in which case the inner radius would be the radius of the magnetosphere), or alternatively radiative stresses (radial radiation force and radiation drag on the accreting matter; \\citealt{coleI}).  Within the framework of such ideas the low kHz QPO frequencies in Cir~X-1 would then imply a relatively stronger magnetic field for the neutron star, or alternatively a larger ratio of radiation flux density to matter density at the inner disk edge. In the description proposed by \\citet{parlines} for the kHz QPO 'parallel tracks' phenomenon, relatively larger radiative stresses might occur if a larger fraction of the accretion takes place not through the (thin) disk, but in a more radial inflow, leading to relatively more radiation, and hence more radiative stress, being produced at the inner disk edge per gram of accreting disk material.  The existence of such a stronger radial inflow might possibly be related to the unusual systemic characteristics of Cir~X-1, such as the periodic episodes of rapid mass transfer and/or the large companion mass and likely, radius.  It is interesting to note, that a tendency towards relatively low kHz QPO frequencies is already known to be one of the characteristics distinguishing Z sources from other neutron star LMXBs.  As noted above, Cir~X-1 shares several more characteristics with the Z sources. So, Cir~X-1 may possess the same physical property that makes Z sources stand out, but more extreme.  If that property would be a stronger magnetic field, as has been discussed before \\citep[e.g., ][]{psalt95}, then additional ingredients (such as smearing of rapid variability by scattering in material surrounding the neutron star) are required to explain why no coherent pulsations are observed. If instead a stronger \\hbox{(radial flow / disk flow)} ratio causes the Z-source phenomenon (possibly, like in Cir~X-1, related to systemic characteristics such as a larger companion radius), then the radiative transfer of the emerging radiation through this inflow may be what causes the spectral effects (the Z track), and the larger inner disk radius what causes some of the timing effects (lower kHz QPO frequencies) that set these sources apart."
    },
    "0608/astro-ph0608330_arXiv.txt": {
        "abstract": "s{This presentation is a Moriond version of our recent paper (Kravtsov, Vikhlinin \\& Nagai~\\cite{2006astro.ph..3205K}) where we discussed X-ray proxies for the total cluster mass, including the spectral temperature ($\\Tx$), gas mass measured within $r_{500}$ ($\\Mg$), and the new proxy, $\\Yx$, which is a simple product of $\\Tx$ and $\\Mg$. We use mock \\emph{Chandra} images constructed for a sample of clusters simulated with high resolution in the concordance $\\Lambda$CDM cosmology. The simulated clusters exhibit tight correlations between the considered observables and total mass. The normalizations of the $M_{500}-\\Tx$, $\\Mg-\\Tx$, and $M_{500}-\\Yx$ relations agree to better than $\\approx 10-15\\%$ with the current observational measurements of these relations. Our results show that $\\Yx$ is the best mass proxy with a remarkably low scatter of only $\\approx 5-7\\%$ in $\\M500$ for a fixed $\\Yx$, at both low and high redshifts and regardless of whether clusters are relaxed or not.  In addition, we show that redshift evolution of the $\\Yx-M_{500}$ relation is close to the self-similar prediction, which makes $\\Yx$ a very attractive mass indicator for measurements of the cluster mass function from X-ray selected samples.  } ",
        "introduction": "\\label{sec:intro}  The evolution of the cluster abundance is one of the most sensitive probes of cosmology. The potential and importance of this method have motivated efforts to construct several large surveys of high-redshift clusters during the next several years. However, in order to realize the full statistical power of the upcoming cluster surveys, it is paramount that the relation between cluster mass and observables and any potential biases are well known.  Several cluster observables based on the galaxy velocities, optical light, X-ray observables such as luminosity, temperature, mass of the intracluster medium (ICM), and Sunyaev-Zel'dovich flux have been proposed as proxies of the total cluster mass (see a recent review by Voit~\\cite{voit05}). In this study we focus on the mass indicators derived from cluster X-ray observables.  X-ray luminosity is the most straightforward mass indicator to measure observationally. However, $\\Lx$ is also the least accurate (internally) of all proposed X-ray proxies for $\\Mtot$ with a large scatter~\\cite{david_etal93,stanek_etal06} and deviations of the slope of the $\\Lx-M$ relation from the self-similar prediction~\\cite{allen_etal03}.  The most common choice of mass proxy in the cluster cosmological studies has been the X-ray temperature of the intracluster gas \\cite{henry_arnaud91,oukbir_blanchard92,markevitch98,ikebe_etal02}. The scatter in the $M-\\Tx$ relation is smaller compared to that in the $\\Lx-M$ relation (the upper limit from observations is $\\approx15\\%$ in $M$ for fixed $T$ for relaxed clusters~\\cite{vikhlinin_etal06}). In general, existence of a tight relation such as $M-\\Tx$ indicates that clusters are regular population of objects with their global properties tightly related to total mass, and scatter caused by secondary effects such as substructure in the ICM, non-gravitational processes, and mergers \\cite{ohara_etal06}. More recently, gas mass was used as a proxy for $\\Mtot$ \\cite{vikhlinin_etal03,voevodkin_vikhlinin04}. The practical advantage of $\\Mg$ over $\\Tx$ is that it can be measured from the X-ray imaging alone. Also, $M_g$ can be expected to be less sensitive to mergers which should translate into smaller scatter in the $M_{\\rm g}-M$ relation. The caveat is that trend of gas mass with cluster mass and evolution with redshift are not yet fully understood.  The use of clusters as efficient probes for precision cosmology puts stringent requirements on observable cluster mass proxies: 1) tight, low-scatter correlation between the proxy and mass, with the scatter insensitive to mergers etc.{}, and 2) simple power-law relation and evolution which can be described by a small number of parameters and be as close as possible to the prediction of the self-similar model. The last point is crucial to ensure that the self-calibration strategies for analyses of large cluster surveys \\cite{levine_etal02,hu03,majumdar_mohr03,lima_hu05,wang_etal04} are successful. This is because self-calibration is powerful when cluster scaling relations and their evolution have a simple form which can be parameterized with a small number of parameters.  In general, a mass proxy does not have to be a single cluster property, such as $\\Lx$, $\\Tx$ or $\\Mg$. Any physically-motivated combination of these variables that is expected to be tightly related to cluster mass can be used to construct a valid mass indicator. A hint for a better X-ray mass proxy is provided by recent studies based on cosmological simulations of cluster formation \\cite{motl_etal05,nagai06}, which show that integrated SZ flux, $Y_{\\rm SZ}$ is a good, robust mass indicator with low scatter in the $Y_{\\rm SZ}-M$ relation, regardless of the dynamical state of the cluster.  In addition, the $Y_{\\rm SZ}-M$ relation exhibits a simple, nearly self-similar evolution with redshift \\cite{dasilva_etal04,nagai06}. The physical reason for the robustness of the SZ flux is straightforward: $Y_{\\rm SZ}$ is directly related to the total thermal energy of the ICM and thus to the depth of the cluster potential well.  Here we show that a similar robust, low-scatter mass indicator can be constructed using X-ray observables. The indicator, which is simply the product of the X-ray derived gas mass and average temperature, {\\boldmath $Y_{\\rm x}=M_{\\rm g}\\,\\Tx$}, correlates strongly with cluster mass with only $\\approx 5-8\\%$ intrinsic scatter. The $\\Yx-\\Mtot$ relation is robust to mergers, in the sense that even for disturbed unrelaxed systems it gives unbiased estimates of mass with the statistical uncertainty similar to that for relaxed systems. In addition, we show that evolution of the slope and normalization of the $Y_{\\rm X}-M$ relation is nearly self-similar.  These properties make $Y_{\\rm X}$ particularly useful for measurements of cluster mass function using X-ray surveys. ",
        "conclusions": "\\label{s:discussion}   We presented comparison of several X-ray proxies for the cluster mass --- the spectral temperature $T_x$ and gas mass $M_g$ \\emph{derived from the X-ray data} within $r_{500}$, and the new proxy, $\\Yx$, defined as a simple product of $\\Tx$ and $\\Mg$.  Analogously to the integrated Sunyaev-Zel'dovich flux, $\\Yx$ is related to the total thermal energy of the ICM. To test these mass proxies, we use mock \\emph{Chandra} ``observations'' of a sample of clusters simulated in the concordance $\\Lambda$CDM cosmology.  The main result of this study is that $\\Yx$ is a robust mass indicator with remarkably low scatter of only $\\approx 5-7\\%$ in $\\M500$ for fixed $\\Yx$, regardless of whether the clusters are relaxed or not. In addition, the redshift evolution of the $\\Yx-M_{500}$ relation is close to the self-similar prediction given by equation~\\ref{eq:yszm}, which makes this indicator a very attractive observable for studies of cluster mass function with the X-ray selected samples.  The $\\Tx-\\M500$ relation has the largest scatter ($\\approx 20\\%$), most of which is due to unrelaxed clusters. The unrelaxed clusters have temperatures biased low for a given mass because a certain fraction of the kinetic energy of merging systems is still in the form of bulk motions of the ICM.  The $\\Mg-\\M500$ relation shows an intermediate level of scatter, $\\approx 10-12\\%$. This relation does not appear to be sensitive to mergers. It does, however, exhibit significant deviations from self-similarity in its slope, which is due to the dependence of gas fraction within $r_{500}$ on the cluster mass \\cite{kravtsov_etal05} (a similar dependence exists for the observed clusters \\cite{vikhlinin_etal06}).  Generally, all the observable--mass relations we tested demonstrate a remarkable degree of regularity of galaxy clusters as a population. $\\Tx$, $\\Mg$, and $\\Yx$ all exhibit correlations with $\\M500$ which are close to the expectation of the self-similar model, both in their slope and evolution with time, within the uncertainties provided by our sample. The only exception is the slope of the $\\Mg-\\M500$ relation.  Given that our analysis relies on cosmological simulations, it is reasonable to ask whether the simulated clusters are realistic. Although simulations certainly do not reproduce all of the observed properties of clusters, especially in their core regions, the ICM properties outside the core in simulations and observations agree quite well. We illustrate this in Fig.~\\ref{fig:tmg}, which shows that the $\\Mg-\\Tx$ relations for simulated and observed clusters \\cite{vikhlinin_etal06}. Clearly, both simulated and observed clusters exhibit tight correlations between $\\Mg$ and $\\Tx$ which agree remarkably in their slope ($M_g \\propto T^{1.75}$) and normalization. The normalizations derived from simulated and real clusters agree to $\\approx 10\\%$, while slopes are indistiguishable and both deviate significantly from the expected self-similar value of $1.5$. This is a consequence of significant trends in the gas fraction with cluster mass, $\\Mg/M_{500}\\propto M_{500}^{0.2\\div 0.25}$ for both simulated \\cite{kravtsov_etal05} and observed clusters \\cite{vikhlinin_etal06}. The deviations from the self-similar model also manifest themselves in the absence of any noticeable evolution with redshift\\footnote{Note that $\\Mg$ in Fig.\\ref{fig:tmg} is not multiplied by the $E(z)$ factor unlike the total mass in Fig.\\ref{fig:txm} and \\ref{fig:yxm}.}. Interestingly, the real clusters show a similarly weak evolution in the $\\Mg-\\Tx$ relation \\cite{vikhlinin_etal02}. Figure~\\ref{fig:res} shows that the likely explanation is that the clusters at $z=0.6$ tend to be colder for the fixed $\\Mtot$ but have higher estimated $\\Mg$ then their counterparts at $z=0$ because they are less relaxed.  \\begin{figure*}[t] \\begin{minipage}[t]{0.48\\linewidth} \\includegraphics[width=0.97\\linewidth,bb=18 177 547 670]{f05.eps} \\caption{Relation between X-ray spectral temperature and gas mass for the relaxed subsample of simulated clusters (circles) and for a sample of relaxed \\emph{Chandra} clusters (\\emph{stars}).  Both gas mass and temperature are the quanities derived from analysis of real and mock X-ray data. The error bars in the \\emph{Chandra} measurements are comparable to the symbol size and are not shown for clarity. The \\emph{dashed line} shows the best fit power law relation with the slope $1.75$.} \\label{fig:tmg} \\end{minipage} \\hfill \\begin{minipage}[t]{0.48\\linewidth} \\includegraphics[width=0.97\\linewidth,bb=18 177 547 670]{f06.eps} \\caption{$\\Yx-\\M500$ relation for the $z=0$ sample of the simulated clusters (\\emph{circles}) and for a sample of relaxed \\emph{Chandra} clusters.  The gas masses for the simulated clusters are appropriately rescaled (see caption to Fig.\\ref{fig:tmg}). The {\\it dot-dashed} line shows the best fit power law relation for the simulated clusters with the slope fixed to the self-similar value of $3/5$. The {\\it dashed} line shows the same best fit power law, but with the normalization scaled down by $15\\%$. } \\label{fig:yxmc} \\end{minipage} \\end{figure*}  A similar level of agreement between the simulations and latest \\emph{Chandra} measurements exists also for the total mass vs.{} temperature relation, $M_{500}-\\Tx$. In fact, the normalization for our simulated sample agrees with the observational results\\cite{vikhlinin_etal06} to $\\approx 10\\%$. This is a considerable improvement over the situation of just several years ago when there was $\\approx 30-50\\%$ discrepancy between observational measurements and cosmological simulations~\\cite{finoguenov_etal01,pierpaoli_etal03}. The $M-\\Tx$ normalization was revised both in simulations and observations due to (1) inclusion of more realistic physics in cosmological simulations (e.g., radiative cooling and star formation, \\cite{dave_etal02,muanwong_etal02}, (2) improved analyses of observed clusters using more realistic gas density profiles \\cite{borgani_etal04,vikhlinin_etal06}, (3) more reliable measurements of the cluster temperature profiles \\cite{markevitch_etal98,arnaud_etal05,vikhlinin_etal06}, and (4) the use of uniform definition of $\\Tx$ in observations and in simulations analyses \\cite{mazzotta_etal04,rasia_etal05,vikhlinin06}. The remaining systematic $10\\%$ difference observed at present is likely caused by non-thermal pressure support from bulk gas motions \\cite{faltenbacher_etal05,rasia_etal06,lau_etal06}, which is unaccounted for by the X-ray hydrostatic mass estimates. In Figure~\\ref{fig:yxmc} we compare the $\\Yx-\\M500$ relation for the simulated clusters and for the \\emph{Chandra} sample\\cite{vikhlinin_etal06}. The observed clusters show a tight correlation with the slope close to the self-similar value. There is $\\approx 15\\%$ difference in normalization, likely explained also by neglecting the turbulent pressure support in the \\emph{Chandra} hydrostatic mass estimates. The excellent agreement of simulations and observations in terms of the relation between the two X-ray observables used to compute $\\Yx$ ($\\Mg-\\Tx$) and a relatively good agreement in the $\\Tx-\\M500$ and $\\Yx-\\M500$ relations, gives us confidence that the results presented here are sufficiently realistic.  Our results show that $\\Yx$ is clearly most robust and most self-similar X-ray cluster mass indicator. The biases existing in mass estimates based on $\\Mg$ and $\\Tx$ anti-correlate both for a given redshift and in terms of evolutionary trends (see Figure~\\ref{fig:res}). This explains why their product, $\\Yx$, is a better mass indicator than $\\Tx$ and $\\Mg$ individually. The quality of $\\Yx$ compares well to that for the actual three-dimension integral of the ICM thermal energy (proportional to $Y_{\\rm SZ}$) in terms of its low scatter and self-similarity.  $\\Yx$ may prove to be an even better mass proxy than $Y_{\\rm SZ}$, given that we use ideal 3D measurement of the latter while reproducing the actual data analysis for the former.  Note also that $Y_{\\rm SZ}$ is more sensitive to the outskirts of clusters, because it involves gas mass-weighted temperature (as opposed to the spectral temperature more sensitive to the inner regions), and thus should be more prone to projection effects.  Note that $\\Yx$ is also an attractive mass proxy from the data analysis point of view. First, it reduces observational statistical noise by combining the two independently measured quantities, $\\Mg$ and $\\Tx$, into a single quantity. For example, a $10\\%$ measurement uncertainty in $\\Tx$ translates into a $\\sim 15\\%$ mass uncertainty through the $M-\\Tx$ relation and only $6\\%$ uncertainty through the $\\Yx-M$ relation. $\\Yx$ is also less sensitive to any errors in the absolute calibration of the X-ray telescope because the biases in the derived $T_X$ and $M_g$ tend to anticorrelate.   The robustness and low scatter make $\\Yx$ an excellent mass indicator for observational measurements of cluster mass function at both $z=0$ and higher redshifts. The necessary data --- an X-ray brightness profile and a wide-beam spectrum excluding the core --- are easily obtained with sufficiently deep observations with \\emph{Chandra}, \\emph{XMM-Newton}, and \\emph{Suzaku} (for low-redshift clusters). The small scatter and simple, nearly self-similar evolution of the $\\Yx-M$ relation hold promise for the self-calibration strategies for future large X-ray cluster surveys.  \\bigskip  This project was supported by the NSF under grants No.{} AST-0206216 and AST-0239759, and by NASA through grants NAG5-13274 \\& NAG5-9217 and contract NAS8-39073. D.N. is supported by a Sherman Fairchild postdoctoral fellowship at Caltech. Last but not least, we would like to thank the organizers of the XLIst \\emph{Rencontres du Moriond}."
    },
    "0608/astro-ph0608106_arXiv.txt": {
        "abstract": "We present results of an analysis of high-resolution spectra ($R \\sim 50\\,000$), obtained with the Subaru Telescope High Dispersion Spectrograph, of two Carbon-Enhanced Metal-Poor (CEMP) stars selected from the Hamburg/ESO prism survey, HE~1305$+$0007 and HE~1152$-$0355, and of the classical CH star HD~5223.  All of these stars have relatively low effective temperatures (4000--4750\\,K) and high carbon abundances, which results in the presence of very strong molecular carbon bands in their spectra. The stellar atmospheric parameters for these stars indicate that they all have surface gravities consistent with a present location on the red-giant branch, and metallicities of ${\\rm [Fe/H]} = -2.0$ (HE~1305$+$0007, HD~5223) and ${\\rm [Fe/H]} = -1.3$ (HE~1152$-$0355). In addition to their large enhancements of carbon (${\\rm [C/Fe]} = +1.8$, $+1.6$ and $+0.6$. respectively), all three stars exhibit strong enhancements of the s-process elements relative to iron.  HE~1305$+$0007 exhibits a large enhancement of the 3rd-peak s-process element lead, with ${\\rm [Pb/Fe]} = +2.37$, as well as a high abundance of the r-process element europium, ${\\rm [Eu/Fe]} = +1.97$. The 2nd-peak s-process elements Ba, La, Ce, Nd, and Sm are found to be more enhanced than the 1st-peak s-process elements Zr, Sr and Y. Thus, HE~1305$+$0007 joins the growing class of the so-called ``Lead Stars'', and also the class of objects that exhibit the presence of both r-process and s-process elements, the CEMP-r/s stars. The large enhancement of n-capture elements exhibited by HE~1152$-$0355 and HD~5223 are more consistent with the abundance patterns generally noticed in CH stars, essentially arising from pure s-process nucleosynthesis. The elemental abundance distributions observed in these stars are discussed in the light of existing theories of CH star formation, as well as the suggested formation scenarios of the CEMP-r/s group. ",
        "introduction": "Over the past few decades, detailed studies of the harvest of metal-poor stars from the HK survey of Beers and colleagues (Beers, Preston, \\& Shectman 1985; Beers, Preston, \\& Shectman 1992; Beers 1999) and the Hamburg/ESO Survey (HES) of Christlieb and collaborators (Christlieb 2003) have revealed them to possess a wide variety of elemental abundance patterns. Among the at-first surprising results was the large fraction of very metal-poor (${\\rm [Fe/H]} < -2.0$) stars with large overabundances of carbon relative to iron. Lucatello et al. (2006) estimates the fraction of carbon-enhanced metal-poor (CEMP) stars, defined as having ${\\rm [C/Fe]} > +1.0$, to be on the order of 20\\% to 25\\% of very metal-poor stars, once the likely dilution of carbon within their atmospheres is taken into account. Subsequent high-resolution spectroscopic studies of a number of CEMP stars (Norris et al. 1997a 1997b, 2002; Bonifacio et al. 1998; Hill et al. 2000; Aoki et al. 2002; Aoki et al. 2006) has revealed that the CEMP stars can be further subdivided into several classes. According to the suggested taxonomy of Beers \\& Christlieb (2005), these include: (a) the CEMP-r stars, which exhibit the presence of strong enhancements of r-process elements, (b) the CEMP-s stars, which exhibit the presence of strongly enhanced s-process elements, (c) the CEMP-r/s stars, which apparently have contributions from both of these neutron-capture sources, and (d) the CEMP-no stars, which exhibit no enhancements of neutron-capture elements.  Numerous questions remain as to the nucleosynthetic histories and astrophysical sites associated with the production of these classes of carbon-enhanced metal-poor stars. It already seems clear that a single mechanism for the production of the enhanced carbon in these stars would be unlikely to lead to such a diversity of heavy-element abundance patterns.  In order to obtain insight into the possible origins of the CEMP stars, and to place additional constraints on the nucleosynthesis of s-process and r-process elements at low metallicity, it is thus necessary to conduct high-resolution observations of as many types of the CEMP stars as possible. As carbon plays a central role in the nucleosynthesis reactions that drive the post main-sequence evolution of stars, one might also obtain clues into the evolutionary processes by studying the range of carbon abundances and carbon isotope ratios exhibited by CEMP stars.  Christlieb et al. (2001) selected a large sample of carbon-enhanced stars from the HES on the basis of the presence of molecular (e.g, CH, CN, and C$_2$) bands on the original prism plates, without placing any additional selection on their metallicities. Several teams have been obtaining low- to medium-resolution spectroscopic follow-up of these stars in order to investigate their nature in more detail. Marsteller et al. (in preparation) will report on medium-resolution spectroscopy of over 350 stars from the Christlieb et al. (2001) sample. Goswami (2005) reported spectral classifications for 91 of the Christlieb et al. (2001) stars based on low-resolution spectroscopy ($R = {\\lambda}$/${\\delta\\lambda} \\sim$ 1330). In this study some 30\\% of the sample were shown to be likely CEMP stars. We have undertaken a high-resolution study of these CEMP candidates (with additional targets from Goswami 2006, in preparation) in order to obtain confirmation of their identification with this class, as well as to carry out a comprehensive study of their heavy-element abundance distributions.  This paper reports the first results of our high-resolution spectroscopic campaign, and discusses two candidate CEMP stars from the Goswami (2005, 2006 in preparation) study, HE~1305$+$0007 and HE~1152$-$0355.  First abundance results for a number of heavy elements in the classical CH star HD~5223 are also presented.  In section 2 we present details of the observations that have been conducted, and describe the low-resolution spectra of our program stars. Section 3 presents $BVRIJHK$ photometry of these stars, and discusses our estimates of their effective temperatures based on this information. The procedures we have adopted for estimation of the stellar atmospheric parameters are discussed in section 4.  The elemental abundance results are presented in section 5. Results and discussions of our present study, and consideration of existing theories for the nucleosynthesis history of CEMP-s and CEMP-r/s stars are presented in section 6.  Section 7 provides some conclusions. ",
        "conclusions": "A high-resolution spectroscopic analysis of the CEMP stars HE~1305$+$0007 and HE~1152$-$0355 indicates large enhancements of C and the s-process elements relative to Fe. The 2nd-peak s-process elements are found to be more enhanced than those of the 1st-peak s-process elements. HE~1305$+$0007 is also found to be characterised by a large enhancement of the 3rd-peak s-process element lead (${\\rm [Pb/Fe]} = +2.37$). The large enhancement of europium (${\\rm [Eu/Fe]} = +1.97$) is another important feature of this star. The CEMP stars exhibiting enhancements of Pb constitutes a growing group of 29 objects; we raise this number to 30 by adding HE~1305$+$0007. In addition, this star is likely to be one of a small number that exhibit double enhancement of r-process and s-process elements, the CEMP-r/s group. The large enhancements of $n$-capture elements exhibited by the classical CH star HD~5223 is consistent with the abundance pattern generally noticed in other CH stars. The enhancement of s-process elements in this star, along with the large enhancement of carbon, indicate that a mass-transfer event in a binary system from a companion AGB star that underwent s-process nucleosynthesis during its lifetime. However, in HE~1305$+$0007, the enhanced Eu abundance, together with the large abundance of the 3rd-peak s-process element Pb, presents a challenge to understanding the formation mechanism of this star. Among several scenarios proposed, the possibility that a star with a high ratio of r-elements/Fe could, by accreting some s-rich matter, become an $r/s$ star with a large [Eu/Fe], has been discussed by many authors. A scenario invoking a triple system, one component polluting the presently-observed star with r-elements, and another polluting with s-elements has also been considered in the literature, but was discarded on the grounds that such a triple system may not be dynamically stable (Cohen et al. 2003). Qian \\& Wasserburg (2003) suggested that the double enhancement in the CEMP-r/s stars could be due to the s-process occurring in an AGB member of a binary system, followed by the r-process taking place in a subsequent accretion induced collapse (AIC) of the white dwarf remnant of the former AGB.  However, the expected rate of the occurence of AIC events in the Galaxy is low and highly uncertain. This rarity cannot justify the substantial fraction of CEMP-r/s stars ($\\sim$ 30\\%) observed among all the CEMP$-s$ stars (Bailyn \\& Grindlay 1990, Qian \\& Wasserburg 2003).  Zijlstra (2004) proposed a different scenario, in which the primary (evolved to an AGB star) transfers s-rich matter to the observed star, but does not suffer large mass loss (owing to its low metallicity). At the end of the AGB phase, the degenerate core explodes as an AGB supernova. However, r-processing in such a scenario is not expected to be significant.  A formation scenario for these stars involving high-mass AGB stars (8-10 M$_\\odot$) was discussed at length by several authors (e.g., Barbuy et al. 2005, Wanajo et al. 2006). The possibility of s-processing occurring in a 10 solar mass star with an O-Ne-Mg core has been studied by Ritossa et al. (1996); these authors demonstrated that the reaction $^{22}$Ne(${\\alpha}$, $n$)$^{25}$Mg is efficient in such high-mass stars, owing to the high temperature 3$\\times$10$^{8}$ K reached at the base of the He-convective shell. An astrophysical scenario associated with the AGB star in a binary system, in which the r-process might also occur, appears to be a positive step toward understanding the formation mechanisms of the CEMP-r/s stars. More quantitative studies of s-process nucleosynthesis in this mass range is needed in order to make definitive conclusions. A detailed and comprehensive discussion on several of the proposed formation scenarios of CEMP-r/s stars is provided by Jonsell et al. (2006).  While in the present work we discuss the carbon and neutron-capture element abundance distribution for only three CEMP stars, we intend to conduct additional analyses for a larger sample of stars in the near future. Aoki et al. (2006) provides one such study, including over 20 CEMP stars. Cohen et al. (2006) describes a sample of 16 CEMP stars. An analysis of the abundance patterns obtained from consideration of larger samples will permit one to put the role of CEMP stars in the early history of Galactic chemical evolution into better perspective. It is also important to seek measurements of the crucial species Li, O, Na, and the mixing diagnostic $^{12}$C/$^{13}$C ratio for these expanded samples of CEMP stars. In addition to many theoretical propositions (Gallino et al. 1998, Goriely \\& Mowlavi 2000, Goriely \\& Siess 2001), there now exists observational evidence that suggests, in contrast to the $r-$only hypothesis of Truran (1981), that the s-process could indeed operate even at very low metallicities, as low as ${\\rm [Fe/H]} = -3.1$ (Johnson \\& Bolte 2002). Sivarani et al. (2006) report on three unevolved (main-sequence turnoff) CEMP stars, at least one of which (CS~29528-041) exhibits moderately high ${\\rm [Ba/Fe]}$ ($\\sim$ +1.0) at a metallicity ${\\rm [Fe/H]} = -3.3$. Masseron et al. (2006) describe the remarkable discovery of an apparent thermally-pulsing AGB star, CS~30322-023, which exhibits a clear s-process signature, at ${\\rm [Fe/H]} = -3.5$. Obviously, there are many surprises yet to be revealed concerning the operation of the s-process in the early Galaxy. Although uncertain, it also seems likely that the impact of an early s-process on the chemical evolution of carbon, nitrogen, oxygen, and the neutron-capture elements in the Galaxy will critically depend on the Initial Mass Function of early-generation stars (Abia et al. 2001; Lucatello et al. 2005; Karlsson 2006).\\\\  {\\it Acknowledgements}\\\\ \\noindent This work made use of the SIMBAD astronomical database, operated at CDS, Strasbourg, France, and the NASA ADS, USA. \\\\ T.C.B. acknowledges partial funding for this work from grant AST 04-06784, as well as from grant PHY 02-16783: Physics Frontiers Center/Joint Institute for Nuclear Astrophysics (JINA), both awarded by the U.S. National Science Foundation. N.C. acknowledges financial support by Deutche Forschungsgemeinschaft through grants Ch~214/3 and Re~353/44. He is a research fellow of the Royal Swedish Academy of Sciences supported by a grant from the Knut and Alice Wallenberg Foundation.  J.E.N. acknowledges support from Australian Research Council grant DP0342613.\\\\"
    },
    "0608/hep-ph0608096_arXiv.txt": {
        "abstract": " ",
        "introduction": "One of the most attractive mechanisms for producing the baryon asymmetry of the Universe, is the out-of-equilibrium decay of coherent scalar field oscillations along nearly $F$- and $D$-flat directions of the scalar potential in supersymmetric theories \\cite{ad,lin}. The minimal supersymmetric standard model contains many of these flat directions \\cite{tony} and one or more can be expected to be excited during inflation \\cite{fluct,eeno}. There are many consequences of the evolution of  large vevs along flat directions \\cite{em}. If, in the context of a grand unified theory (GUT), there exists a scalar operator $\\mathcal{O}$, which violates baryon number and for which $\\langle \\mathcal{O} \\rangle \\ne 0$, the CP violating decay of the flat direction will produce a large baryon asymmetry. Alternatively, any operator violating lepton number could be used to generate a net lepton asymmetry \\cite{cdo} which is subsequently converted to a baryon asymmetry by sphaleron processes.  Neglecting nonperturbative effects, one typically finds that the scalar field oscillations of a flat direction persist after the decay of the inflaton \\cite{eeno,cgmo,am}. Indeed, quite generally, the Universe becomes dominated by flat direction oscillations until these decay.  It has also been recently argued \\cite{am} that flat directions can delay the thermalization of the inflaton decay products (resulting in a low reheat temperature) by providing a large mass to gauge bosons through their vev, and thus suppressing the rates of the scattering processes needed for thermalization.  Here, we examine the effects of preheating on the fate of flat directions. We find that nonperturbative effects can be important and lead to the rapid decay of the flat directions, long before they can dominate the expansion rate of the Universe. Similarly, the rapid decay may preclude any delay in the thermalization process after the decay of the inflaton. These effects are due to mixing between different MSSM excitations around many different sets of flat directions\\footnote{Two or more mutually non-exclusive flat directions are required. As we show below, only in exceptional cases does a single flat direction exclude all others from being excited.}, and have been overlooked in the literature so far.  We will work in the context where the inflaton, $\\psi$, couples to standard model fields only through gravity\\footnote{Reheating in the context of nonperturbative inflaton decay has been discussed in~\\cite{rehnonpert}.}.  We will also assume for simplicity that there is only a single mass scale associated with the inflaton potential set by the COBE normalization \\cite{cobe}, with $m_\\psi \\sim H_I \\sim 10^{-7} M_P$, where $H_I$ is the Hubble parameter during inflation and $M_P$ is the Planck scale. We also assume that after inflation, $\\phi$ oscillates with a Planck scale amplitude. The energy density in oscillations can be written as \\beq \\rho_\\psi = m_\\psi^2 \\psi^2 = m_\\psi^2 M_P^2 (R_\\psi/R)^3 \\label{rhopsi} \\eeq where $R_\\psi$ is the Robertson-Walker scale factor when oscillations begin.  For $R > R_\\psi$, the Universe expands as a matter dominated Universe with Hubble parameter, $H \\simeq m_\\psi (R_\\psi/R)^{3/2}$. Inflatons decay at $R_{d\\psi} \\simeq (M_P/m_\\psi)^{4/3} R_\\psi$ when their decay rate $\\Gamma_\\psi = m_\\psi^3/M_P^2 \\simeq H$. After inflaton decay, the Universe is radiation dominated and $\\rho_{r\\psi} \\simeq m_\\psi^{2/3} M_P^{10/3} (R_\\psi/R)^4$, where $\\rho_{r\\psi}$ is the energy density in the radiation produced by inflaton decay.  Thermalization occurs when the relevant scattering rates become comparable to the Hubble parameter.  Naive estimates indicate that this may be delayed due to suppressed cross sections of the form $\\sigma \\sim (\\alpha^2/m_\\psi^2)(R/R_{d\\psi})^2$ where $\\alpha$ is a characteristic coupling strength (presumably of order the gauge fine structure constant) and the suppression is due to the initially large particle energies produced in the decay ($E \\sim m_\\psi$) which redshift to lower energies as the Universe expands.  The resulting reheat temperature in this case would be expected to be low, $T_R \\sim 10^4$ GeV \\cite{eeno}. However it was shown \\cite{ds} that forward scattering in 2 $\\to$ 3 interactions, though of higher order in $\\alpha$, lead to rapid thermalization.  It was estimated that the inelastic rates were roughly $\\Gamma_{inel} \\sim \\alpha (M_P/m_\\psi)n_\\psi \\sigma$ in which case thermalization occurs almost immediately after decay and restores the simple estimate $T_R \\sim (\\Gamma_\\psi M_P) \\sim 10^8$ GeV. This will be true so long as $\\alpha^3 > m_\\psi/M_P$.  In supersymmetric theories, quantum fluctuations can drive several scalar fields to large expectation values along flat directions. In fact, in supergravity models with a Heisenberg symmetry \\cite{BG}, such as no-scale supergravity \\cite{ns}, we expect to generate vev's of order $\\phi_0 \\sim M_P$ for any flat direction  not involving stops \\cite{gmo}.  This vev may be smaller if the flat direction is regulated by non-renormalizable operators \\cite{cgmo}.  When the Hubble parameter after inflation is of order the supersymmetry breaking scale $m_\\phi$, sfermion oscillations along the flat direction will begin. This occurs at $R = R_\\phi \\simeq (m_\\psi/m_\\phi)^{2/3} R_\\psi$. The energy density in the oscillations of the flat direction field characterized by $\\phi$ is given by \\beq \\rho_\\phi = m_\\phi^2 \\phi^2 = m_\\phi^2 \\phi_0^2 (R_\\phi/R)^3 \\simeq m_\\psi^2 \\phi_0^2 (R_\\psi/R)^3 \\eeq Typically, oscillations of $\\phi$ will begin before inflaton decay, since $(R_{d\\psi}/R_\\phi) \\sim M_P^{4/3} m_\\phi^{2/3}/m_\\psi^2 \\sim 10^3$.   We note that although most standard model (SM) fields are massive when the flat direction is excited, inflaton decay is unimpeded if the inflaton can decay directly into the light fields making up the flat direction. Nevertheless, the scattering rates will be suppressed due to the large masses of exchanged particles \\cite{am}. To compute the thermalization rate of the decay products, we once again consider the inelastic 2 $\\to$ 3 scatterings by replacing $\\sigma$ with $(\\alpha^2/\\phi^2)$ and $\\phi^2 = (\\phi_0^2 m_\\psi^2/m_\\phi^2) (R_\\psi/R)^3$. In this case,  thermalization is not suppressed unless $\\alpha <m_\\psi^{5/3}/m_\\phi^{2/3} M_P$.  For $\\alpha^2 \\sim 10^{-3}$,  one finds again a reheat temperature of order $T_R \\sim \\alpha^{3/2} m_\\phi M_P/m_\\psi \\sim 10^8$ GeV. Note that we have assumed $\\phi_0 = M_P$.  For lower $\\phi_0$, reheating at this stage will always be given by the instantaneous rate yielding $T_R \\sim 10^8$ GeV. However as we will see, the energy density of the Universe will become dominated by the oscillations along the flat direction and final thermalization will depend on the decay of the flat direction vev\\footnote{As we show below, nonperturbative effects associated with mixing dramatically alter this conclusion.}.   If inflaton decay occurs before the decay of flat direction, the radiation produced by inflaton decay will redshift faster than the flat direction oscillations.  It is of interest, therefore, to compare the energy density in inflaton decay products with that in the flat direction oscillations at the time of decay (at $\\phi = \\phi_d$). The ratio of energy density in radiation to oscillations is \\beq \\frac{\\rho_{r\\psi}}{\\rho_\\phi} = \\left( \\frac{M_P^{10/3}}{m_\\psi^{4/3} \\phi_0^2} \\right) \\left( \\frac{R_\\psi}{R} \\right) \\eeq The oscillations end when the decay rate, $\\Gamma_\\phi \\sim m_\\phi^3/\\phi^2 \\sim H$, from which we deduce that $R_{d\\phi} = (m_\\psi^{7/15} \\phi_0^{2/5} M_P^{2/15} /m_\\phi) R_\\psi$. At this value of $R$, we see that inflaton decays dominate the energy density only if $\\phi_0 < m_\\phi^{5/12} M_P^{4/3} / m_\\psi^{3/4} \\sim 10^{-2} M_P$. For larger values of the initial vev along the flat direction, reheating and thermalization is entirely determined by the dynamics of the flat direction \\cite{cgmo}. In this case, we expect a baryon asymmetry of order $\\mathcal{O}(1)$ \\cite{ad,lin} unless it is diluted by other processes such as moduli decay or the vev is regulated by non-renormalizable or GUT scale operators \\cite{cgmo}.  For relatively large $\\phi_0$, the oscillations of the flat direction come to dominate the energy density.  In this case the decay of $\\phi$ is determined by $\\Gamma_\\phi = (m_\\phi^5 / \\phi_0^2 m_\\psi^2) (R / R_\\psi)^3 = (m_\\psi \\phi_0 / M_P) (R_\\psi/R)^{3/2} = H$. The scale factor at decay is $R_{d\\phi} = (\\phi_0^{2/3} m_\\psi^{2/3} / m_\\phi^{10/9} M_P^{2/9} ) R_\\psi$. Because this decay occurs relatively late,  the reheat temperature is suppressed, $T_R \\sim m_\\phi^{5/6} M_P^{1/6}$  $\\sim 10^6$~GeV.   The above argument for the decay of $\\phi$ ignores the possibility that the flat directions decay on a much quicker timescale in a nonperturbative fashion (preheating). In the following, we show that these effects can be indeed dominant, and lead to the quick depletion of the flat directions within their first ($\\sim 5-10$) rotations. If so, the Universe remains dominated by inflaton oscillations until inflaton decay. Reheating in this case is produced by inflaton decays, and the baryon asymmetry is expected to be large unless $\\phi_0$ is regulated by non-renormalizable operators or dilution due to moduli decay occurs.  The plan of the paper is as follows. In Section \\ref{sec2}, we discuss an argument, based on a toy model of preheating, which has been used to claim that preheating along flat directions is negligible. In the remainder of that Section, and in Section \\ref{sec3}, we explain why this conclusion might not remain valid when realistic models are considered. In Section \\ref{sec4}, we present an      explicit computation of the actual nonperturbative decay. Section \\ref{sec5} contains our conclusions. ",
        "conclusions": "\\label{sec5}  The scalar potential of the MSSM is characterized by several flat directions. This flatness may be lifted by supersymmetry breaking terms, nonrenormalizable operators, and/or supergravity corrections. The latter are of particular relevance in the early Universe. For instance, if the MSSM fields have a minimal Kahler potential, then supergravity terms provide a mass (of order H) for the flat directions which prevents them from being excited during inflation \\cite{drt}. However, if the Kahler potential exhibits a Heisenberg symmetry \\cite{BG}, such mass terms do not arise at the tree level and the flat directions can indeed develop a large vev during inflation \\cite{gmo}. No-scale supergravity \\cite{ns} is one such example. The development of a large vev may have interesting consequences for the subsequent evolution of the Universe.   Among the consequences of a large vev along a supersymmetric flat direction, is the possible generation of a baryon asymmetry \\cite{ad}. Furthermore,  the energy density stored along the flat direction may come to dominate over the inflaton decay products if the flat direction is long-lived. It is tempting to conclude that this is the case if one simply considers the perturbative (single quanta) decay of the flat direction. Indeed, the flat direction, due to its large vev, provides a large mass to the fields to which it is coupled, so that the perturbative decay proceeds by loop(s), with very heavy fields in the propagator. This typically results in a small decay rate.  However, this consideration does not rule out the possibility that the flat directions decay nonperturbatively. It has been shown \\cite{kls}, that in many models of inflation, the inflaton decays nonperturbatively through a collective effect called preheating. In this paper, we have investigated the possibility of preheating during the evolution of fields associated with supersymmetric flat directions. In particular, it is important to determine whether preheating can lead to a significantly quicker decay than the perturbative estimates may suggest. Prior to our investigation, this question was addressed only in a very limited number of papers, which concluded that preheating is negligible in this case. Their claim, however, was based on toy models, where the flat direction $\\phi$ is coupled to a single scalar $\\chi$ through a quartic $\\vert \\phi \\vert^2 \\vert \\chi \\vert^2$ interaction. The time varying amplitude of $\\phi$ results in a time varying mass for the real and imaginary components of $\\chi \\,$. However, in the specific case in which $\\phi$ is a flat direction (rather than the inflaton) the variation of this mass is always too slow to lead to significant nonperturbative particle production.  We have shown, however,  that the couplings of the flat direction encoded in the MSSM potential are more complicated than the simple $\\vert \\phi \\vert^2 \\vert \\chi \\vert^2$ interaction. In particular, the part of the potential from the $D-$terms results in couplings of the form $\\phi^2 \\chi_i \\chi_j$, leading a non-diagonal mass matrix for the fields $\\chi_i$. Even if the eigenmasses are slowly varying, the mass matrix ${\\cal M}^2$ itself varies with time. The dominant time dependence is due to the phase $\\sigma$ of the flat direction, which is rotating quasi periodically (with a small decrease of the amplitude) around the minimum of its potential. The quasi periodic variation of ${\\cal M}^2$ leads to a very strong parametric resonance effect. We found that when at least two flat directions are excited, the flat directions typically loses most of their energy density within its first $5-10$ rotations due to these couplings. From this point on, backreaction effects of the $\\chi_i$ fields can be expected to play a very relevant role in the evolution of the system. Their complete investigation requires the use of numerical simulations on the lattice \\cite{latticeeasy}, which are beyond the aim of the present work\\footnote{ Numerical simulations are also required to determine if non-linear interactions affect the early stages of preheating.}. However, due to the strong interactions between the MSSM quanta, these interactions are likely to lead to a fast depletion of the (remaining) vev of the flat direction. As such, it is very unlikely that the scalar field oscillations of the flat direction can ever come to dominate over the inflaton field and its decays product, particularly in the case in which the inflaton itself decays only gravitationally.  Although the final stage of the decay remains to be explored, we have thus showed that flat directions can potentially have a very fast decay channel which has been so far overlooked in the literature. There are several important consequences of this rapid decay.  If flat directions never dominate the energy density of the Universe, entropy production and reheating are determined by the dynamics of the inflaton. The baryon asymmetry will be measured with respect to the entropy density produced by inflaton decays and will typically be large requiring either further dilution of the baryon asymmetry by later entropy production (e.g. due to moduli decays) or regulating the initial vev along the flat direction with some non-renormalizable operator \\cite{cgmo}. Finally, in the absence of a large vev and hence large particle masses, inflaton decay and the subsequent thermalization also proceeds rather quickly \\cite{ds}, leading to a relatively large reheat temperature after inflation which could have dire consequences on the gravitino problem.  \\vspace{1cm} \\noindent{ {\\bf Note added} } : After the present work was completed, a new paper which discusses the decay of flat directions \\cite{am2} came to our attention. The first part of \\cite{am2} is devoted to the case of a single flat direction. The analysis, and the conclusions, coincide with those already given in the present work. We comment on two of the remarks made in the second part (devoted to the case of two or more flat directions), although neither of them has been supported by a concrete computation. The first remark claims that flat directions may have hierarchical vevs (a small ratio $r$ between the two vevs), so that the system may be effectively reduced to the single flat direction case. In general, the amount of produced particles will be suppressed by the ratio $r$. However, for any finite $r$, the number densities of the non-goldstone massless fields still grow exponentially with the number $N$ of rotations of the flat directions. For the toy model with two flat directions, we have verified numerically that most of the energy density in the flat directions decays after $N \\sim 10$ when $r = 10^{-3}$ (as opposed to $N \\sim 5$ when $r=1$). The second remark is that nonrenormalizable terms in the potential, which depend also on the phase of the flat directions, can in some cases result in chaotic motion for the phases, potentially suppressing the parametric amplification which takes place when the phases move with nearly constant velocity. First, nonrenormalizable interactions are very model dependent, and do not necessarily lead to chaotic motion for the phases. Second, even if they do, nonrenormalizable terms are relevant only for the first few oscillations of the flat directions. As soon as the flat directions evolve, their vev is decreased by the expansion of the universe, entering into a region where their potential can be safely approximated by the quadratic (phase-independent) part. When this happens, the phases will evolve with a nearly constant velocity, as the cases considered here.    \\vskip 0.5in \\vbox{ \\noindent{ {\\bf Acknowledgments} } \\\\ \\noindent M.P. thanks Lorenzo Sorbo for collaboration in obtaining the formalism used in this paper and in its study. We also thank Joel Giedt, Lev Kofman,  Erich Poppitz, and Misha Shifman for very useful discussions. We would particularly like to thank Rouzbeh Allahverdi for very pertinent comments on the original version of this manuscript. This work was partially supported by DOE grant DE--FG02--94ER--40823. M.P. also acknowledges a grant from the Office of the Dean of the Graduate School of the University of Minnesota. }"
    },
    "0608/astro-ph0608383_arXiv.txt": {
        "abstract": "We present a detailed analysis, including photoionization modeling, of the intrinsic absorption in the Seyfert 1 galaxy NGC 4151 using ultraviolet (UV) spectra from the {\\it Hubble Space Telescope}/Space Telescope Imaging Spectrograph and the {\\it Far Ultraviolet Spectrographic Explorer} obtained 2002 May as part of a set of contemporaneous observations that included {\\it Chandra}/High Energy Transmission Grating Spectrometer spectra. In our analysis of the {\\it Chandra} spectra, we determined that the X-ray absorption was dominated by two components: a high-ionization absorber, revealed by the presence of H-like and He-like lines of Mg, Si, and S, and a lower-ionization absorber, in which inner-shell absorption lines from lower-ionization species of these elements formed. We identified the latter as the source of the saturated UV lines of O~VI, C~IV, and N~V associated with the absorption feature at a radial velocity of $\\sim$ $-$ 500 km s$^{-1}$, which we referred to as component D$+$E. In the present work, we have derived tighter constrains on the the line-of-sight covering factors, densities, and radial distances of the absorbers. We confirm the presence of the 3 sub-components of D$+$E described in our previous paper, with line-of-sight covering factors ($C_{los}$) ranging from $\\sim$ 0.5 -- 0.9, and find evidence for a fourth component, D$+$Ed characterized by low ionization and a $C_{los}$ $\\sim$ 0.2. The complexity of the UV absorption in NGC 4151 may be a consequence of the fact that we are viewing the black hole/accretion disk system at a relatively high inclination and, therefore, may be detecting the densest part of the flow. Our deconvolution of the underlying C~IV emission indicates that D$+$E must lie outside the Intermediate Line Region, hence at a radial distance of $\\sim$ 0.1 pc. We find that the Equivalent Widths (EWs) of the low-ionization lines associated with D$+$E varied over the period from 1999 July to 2002 May. Although over part of this time, the variations were correlated with changes in the UV continuum, the drop in the EWs of these lines between 2001 April and 2002 May are suggestive of bulk motion of gas out of our line-of-sight. Over this period, $C_{los}$ for the low-ionization absorption lines dropped from $\\sim$ 0.7 to $\\sim$ 0.2. If these lines from these two epochs arose in the same sub-component, the transverse velocity of the gas is $\\approx$ 2100 km s$^{-1}$. This is similar to the constraint on transverse velocity derived from the drop in the X-ray absorbing column between 2000 March and 2002 May. Transverse velocities of this order are consistent with an origin in a rotating disk, at the roughly radial distance we derived for D$+$E. As we suggested in our previous study, it is likely that the absorption arises in a disk-driven wind. ",
        "introduction": "Understanding the nature of mass outflow in active galactic nuclei (AGN) has been an area of extensive recent research activity. The UV and X-ray spectral properties of such outflows and their overall physical properties have been detailed by Crenshaw, Kraemer, \\& George (2003). To summarize, the ubiquity of blue-shifted absorption among Seyfert galaxies indicates that the global covering factors of the absorbers must be high. The total mass in the outflows can exceed the mass of gas in the broad line region (BLR). Furthermore, the mass outflow rates are on the same order as the accretion rates required to fuel the supermassive black holes that power AGN. Hence, determining the nature of these outflows is critical to understanding the structure and evolution of AGN. However the origin of the outflows, e.g. an accretion disk (Konigl \\& Kartje 1994) or the putative torus (Krolik \\& Kriss 1995, 2001), and the physical means by which they are accelerated, e.g. thermal winds (Krolik \\& Begelman 1986), radiation pressure (Murray et al. 1995), or magneto-hydrodynamic (MHD) flows (Bottorff, Korista, \\& Shlosman 2000), remain uncertain.  NGC 4151 ($cz =  995$ km s$^{-1}$) was the first Seyfert galaxy to show intrinsic absorption. Oke \\& Sargent (1968) found non-stellar He~I $\\lambda$3889 and Anderson \\& Kraft (1969) detected H$\\beta$ and H$\\gamma$ self-absorption. These lines were blue-shifted with respect to the host galaxy, indicating mass outflow. Observations in the UV obtained with the {\\it IUE} (Boksenberg et al. 1978) and far-UV using the Hopkins UltraViolet Telesocope (HUT; Kriss et al., 1992, 1995) revealed absorption lines from a wide range of ionization states, including fine-structure and metastable lines. The intrinsic lines were found to vary in response to the UV continuum (Bromage et al. 1985), suggesting ionic column density changes resulting from changes in the ionization state of the absorbers. Monitoring of NGC 4151 in the far-UV with HUT (Kriss et al. 1995) and the {\\it Orbiting and Retrievable Far and Extreme Ultraviolet Spectrometers-Shuttle Pallet Satellite II} mission (Espey et al. 1998) also showed evidence for short timescale variations in absorption due to changes in ionization. Early X-ray spectra revealed the presence of a large column of absorbing gas (Barr et al. 1977; Holt et al. 1980), which subsequent observations revealed to be ionized (Yaqoob, Warwick, \\& Pounds 1989; George et al. 1998). As with the UV absorption, the X-ray absorption is highly variable (George et al. 1998; and references therein), but the changes are not always correlated with the X-ray continuum flux and may indicate transverse motion of the gas across the line of sight (Weaver et al. 1994).  In the {\\it Hubble Space Telescope/HST} survey of intrinsic UV absorption in Seyfert galaxies (Crenshaw et al. 1999), NGC 4151 stood out as a unique object. It is one of the few objects whose spectra that show Mg~II absorption, and it is the only one that showed strong metastable C~III* $\\lambda$ 1175 absorption, until recent detection of a weak C~III* absorption line in NGC 3783 (Gabel et al. 2003). One possible reason for the complexity of the intrinsic absorption in NGC 4151 is the fact that we are viewing the black hole/accretion disk system at a relatively high inclination of $\\sim$ 45\\deg~with respect to the accretion disk (Evans et al. 1993; Das et al. 2005), and are likely detecting material close to the densest part of the outflow. The implication is that this material cannot be detected in absorption if the AGN is viewed more face-on.  The first high-resolution ($\\sim$ 15 km s$^{-1}$) UV spectra obtained by Weymann et al. (1997) using the Goddard High Resolution Spectrograph (GHRS) aboard {\\it HST} revealed six major kinematic components of C~IV and Mg~II absorption, which were stable over the period 1992--1996.  Following this, we obtained an {\\it HST}/Space Telescope Imaging Spectrograph (STIS) GTO spectrum in 1999 July (Kraemer et al. 2001; hereafter K01). These spectra revealed that the kinematic components detected by Weymann et al. were still present, although the strongest kinematic components, D and E using the nomenclature of Weymann et al.,  were broad and we were unable to separate them. Also, we found complex line-of sight covering factors\\footnote{Here, we generalize partial covering factor $C_{los}$ to include either unocculted regions in our direct view to the nuclear emission sources, or light that is filling in the absorption troughs due to an extended scatterer.} ($C_{los}$) in the broad D+E component. Based on the relative depths of individual members of doublets and multiplets, different ions in this component were found to have different covering factors: the strong high-ionization \\ion{C}{4} and \\ion{N}{5} doublets had $C_{los} \\approx$ 0.9, the \\ion{Si}{4} and \\ion{Mg}{2} doublets yielded $C_{los} \\approx $0.85 and 0.62, respectively, and the many \\ion{Fe}{2} metastable lines indicated even lower covering factors. From our photoionization modeling analysis, we argued that the complex, low ionization absorption arose in dense gas, within 0.03 pc of the central source. The presence of these low ionization lines was a consequence of the relatively low continuum state during this observation, as originally suggested by Bromage et al. (1985). However, the lack of contemporaneous X-ray observations prevented us from fully constraining the physical state and geometry of the circumnuclear absorbers.  In order to better characterize the absorbers, we obtained contemporaneous {\\it FUSE}, {\\it HST}/STIS, and {\\it Chandra}/High Energy Transmission Grating Spectrometer(HETGS) spectra in 2002 May (see next section). In Kraemer et al. 2005 (hereafter Paper I), we presented our analysis of the {\\it Chandra} observations. To summarize, we were able to model the intrinsic X-ray absorption using five distinct components (see Table 1). The H-like Mg, Si, and S lines were consistent with the presence of a highly ionized component, X-High, which was not detected in the UV spectra. The bulk of the soft X-ray absorption is associated with the UV kinematic component D$+$E. Based on our preliminary analysis of the UV absorption, we found evidence for three subcomponents associated with D$+$E. The saturated O~VI, N~V, and C~IV lines have $\\sim$ 10\\% residual flux in the deepest part of their troughs. We identified this component, D$+$Ea, as the source of the inner-shell lines of S, Si and Mg detected in the HETG spectrum. The \\ion{C}{3}*~$\\lambda$1175 absorption, arising from the $2s2p~^{3}P$ metastable levels, over this range of velocity shows a higher residual flux ($\\sim$ 20\\%), hence, we added a second component, D$+$Eb. Finally, the \\ion{P}{5} doublet indicated the presence of a third component, D$+$Ec, with a much lower covering factor ($C_{los}  \\approx 0.5$). A fifth component of X-ray absorption was identified with UV kinematic component E$'$.  In Paper I, we also analyzed the 2000 March 3 {\\it Chandra}/HETGS observation of NGC 4151. We found that the 2-10 keV continuum flux was a factor of $\\sim$ 2 lower than in 2002 and that the soft X-ray continuum was much more heavily absorbed. Although higher soft X-ray opacity would be a consequence of the drop in ionization of the absorbers, due to the weaker ionizing flux, in order to fit the spectrum we required an additional column of gas, which we associated with component D$+$Ea, in which case it must cover the continuum source and the Broad Line Region (BLR). This column must have moved out of our line-of-sight in the time between the two observations, which requires a transverse velocity\\footnote{in Paper I, we had used the radius of the broad line region to estimate v$_T$, while, here, we correctly use the diameter.} $v_{T} > 2500$ km s$^{-1}$.  Based on our photoionization models, we found that X-High could not be accelerated via radiation pressure or as part of a thermal wind. Although the other components could be radiatively driven, the overlap in velocities of X-High and the components of D$+$E suggests a common mechanism. We suggested that the most likely means for driving the outflow is by an MHD wind arising from the accretion disk. Interestingly, a disk origin would provide a natural mechanism for the observed transverse bulk motion: i.e., rotation.   In the present study, we have used the parameters determined in Paper I as the initial model inputs, subject to modest adjustments in order to refine the fit to the UV absorption, with the main goal being to more tightly constrain the radial distances of the absorbers. We also track the changes in the UV absorption over the period from 1997-2002 to examine the effects of the response to the ionizing continuum and/or bulk motion. ",
        "conclusions": "Based on our modeling of the UV and X-ray absorbers, we have developed a set of constraints on the physical conditions and evolution of the absorbers. The main results are as follows.  1) The bulk of the soft X-ray absorption and the strong Li-like CNO absorption lines associated with kinematic component D$+$E can be modeled as forming in a single component, D$+$Ea. There are three additional lower-ionization sub-components of D$+$E, which do not contribute significantly to the X-ray absorption.  2) Our deconvolution of the C~IV profile indicates that at least some of the absorbing gas associated with D$+$E must lie outside the ILR. If the ILR emission arises at a radial distance of 0.1 pc, from our derived values for $U$ we can constrain the densities of the absorbers. Using these densities, and the predicted ionic columns from the metastable states of C~III when these densities are input to the models, we were able to fit the observed C~III* profile, hence our assumptions are self-consistent.  3) The EWs of the low ionization lines associated with D$+$E have varied over the period from 1999 July to 2002 May. Before 2001 April, these changes were correlated with variations in the UV continuum, as expected for changes in ionic columns due to the response of the gas to the ionizing continuum. However, between 2001 April and 2002 May, the drop in the EWs of these lines are more suggestive of bulk motion, as does the change in the X-ray opacity between between 2000 March and 2002 May (Paper I). The weakness of the metastable Fe~II lines in 2002, compared to 1999 when NGC 4151 was in a similar flux state, must also be due to bulk motion. However, it is uncertain how the column densities of the individual sub-components of D$+$E have varied.  4) We have obtained constraints on the transverse velocity, $v_{T}$ $\\approx$ 2100 km s$^{-1}$, for the low ionization absorption associated with D$+$E, which we dubbed D$+$Ed. While this is roughly comparable to the lower limit derived for D$+$Ea in Paper I, the connection between these sub-components and the origin of the strong low ionization absorption present in the 1999 July STIS spectra is not clear.   In summary, while we have modified some of the details of the models used in Paper I, and included additional components of low-ionization gas, the UV analysis has confirmed our two principal conclusions. First, the bulk of the absorption in NGC 4151 is much closer to the central source than in other Seyferts, e.g. NGC 3783 (Gabel et al. 2005b) and, second, that the variations in the absorption are due to both ionization and bulk motion. The latter could be a consequence of an origin within a rotating disk, which is consistent with our suggestion that the outflow is MHD-driven.  An interesting characteristic of the absorbers which is yet to be addressed is the origin of low covering factor sub-components, such as D$+E$c. Although one can envision denser, lower ionization knots embedded in a more tenuous flow, the dynamics of such a system are not straightforward. One major problem is that there are multiple components, characterized by different $U$'s and, thus different Force Multipliers (see Paper I and references therein), with the same radial velocities. While this may be consistent with an MHD flow, it requires that the low ionization knots are coupled to the flow in such a way that they are not radiatively driven through the more tenuous gas."
    },
    "0608/astro-ph0608660_arXiv.txt": {
        "abstract": "{Numerical simulations of planet-disk interactions are usually performed with hydro-codes that -- because they consider only an annulus of the disk, over a 2D grid -- can not take into account the global evolution of the disk. However, the latter governs planetary migration of type~II, so that the accuracy of the planetary evolution can be questioned.} {To develop an algorithm that models the local planet-disk interactions together with the global viscous evolution of the disk.} {We surround the usual 2D grid with a 1D grid ranging over the real extension of the disk. The 1D and 2D grids are coupled at their common boundaries via ghost rings, paying particular attention to the fluxes at the interface, especially the flux of angular momentum carried by waves. The computation is done in the frame centered on the center of mass to ensure angular momentum conservation.} {The global evolution of the disk and the local planet-disk interactions are both well described and the feedback of one on the other can be studied with this algorithm, for a negligible additional computing cost with respect to usual algorithms.} {} ",
        "introduction": "Planetary formation occurs in disks of gas and dust around protostars. In particular, giant planets -- whose mass is mostly made of hydrogen and helium -- must have formed before the dissipation of the gas disk. Consequently, they must have exerted tidal forces on the gas and their orbits must have evolved in response to the gas.  In particular, the presence of a planet on a circular orbit leads to the formation of a spiral density wake in the disk. \\citet{GT80} and \\citet{LinPapaloizou1979} showed that, through the wake, the planet gives angular momentum to the part of the disk exterior to its orbit, while it takes angular momentum from the inner part. For small mass planets, \\citet{Ward1997} showed that the net result is a loss of angular momentum for the planet, which makes its orbit decay on a short timescale. This is usually referred to as \\emph{type~I migration}.  As the planet gives angular momentum to the outer part of the disk, it repels it outward\\,; symmetrically it repels the inner part inward. If the planet is massive enough (the threshold mass depending on the disk's viscosity and scale height\\,; see \\citet{Crida2006} and references therein) a clear gap opens around the planet's orbit, effectively splitting the disk into an internal and an external part. In this situation, the planet gets locked in the middle of the gap, because both the outer and the inner part of the disk are repellent for it. Thus, the migration of the planet has to go in parallel with the migration of the gap \\citep{LinPapaloizou1986b,Ward2003}, which in turn has to follow the viscous evolution of the disk (characterized by a radial spreading of the disk as the gas is accreted onto the central star \\citep[see][]{LBP74}. This is usually referred to as \\emph{type~II migration}.  Although analytic theories have brought a great deal of understanding of the fundamentals of planet-disk interactions, it has become increasingly evident in the last years that numerical simulations are an essential tool of investigation \\citep[see][ for the most recent review]{Papaloizou-etal-2006}. Numerical simulations, however, are difficult and time-consuming, and they inevitably involve some simplifications, which might turn out to produce artifacts. In this paper we are concerned with the simulation of type~II migration.  As we have seen above, a proper simulation of type~II migration requires not only the correct calculation of planet-disk interactions -- which are essentially \\emph{local} -- but also of the \\emph{global} evolution of the disk.  Unfortunately, planet migration is typically simulated with hydro-dynamical codes using a 2 dimensional polar grid which, for numerical reasons, is truncated at an inner and an outer radius. This enables to describe well the local interaction of the planet with the disk but not the global evolution of the disk.  Indeed, the boundary conditions at the extreme rings of the polar grid can not take into account what happens in the whole disk, outside of the grid. For instance, the use of \\emph{open} boundary conditions allows the gas to leave the region covered by the grid\\,; this makes the studied disk annulus behave as if it were surrounded by vacuum, so that it empties very rapidly, which is obviously not realistic. In the opposite extreme case, boundary conditions which impose that the mean density on the extreme rings remains constant with time, disable the accretion and spreading of the gas. Some other prescriptions between these two extremes may be used, (for instance allowing inflow, or imposing an outflow given by an analytical model, or setting the flow on the last ring equal to the viscous flow measured in the neighboring rings, etc.). However, it is very difficult to adapt these prescriptions to the changing behavior of the disk with time, in particular when the disk undergoes perturbation from the planets. In any case these prescriptions are rather arbitrary and, thus they may introduce possible artifacts in the planetary evolution.  In this paper we present a novel idea for the correct calculation of the global evolution of the disk. It consists in surrounding the 2D polar grid with a 1D radial grid. The 1D grid extends from the real inner edge of the gas disk (e.g. the X-wind truncation radius at a few tenths of AU) to the real outer edge (e.g. the photo-dissociation radius at hundreds of AU). This 1D grid has open boundaries at the inner and outer edges, and exchanges information with the 2D grid for the definition of realistic, time-dependent boundary conditions of the latter. Our algorithm for the interfacing between the 1D and 2D grids is driven by the requirement that the angular momentum of the global system (the disk in the 2D section, plus the disk in the 1D section plus the planet-star system) is conserved.  We will describe it in detail in Sect.~\\ref{sec:coupling}.  First, however, we need to revisit the algorithm used for modeling the gas evolution and planet-disk interactions on the 2D grid, to ensure that it also conserves angular momentum, which is often not the case in standard implementations. We discuss this issue in Sect.~\\ref{sec:cons2D}. In Sect.~\\ref{sec:results}, the results of our new algorithm will be discussed, in terms of CPU time and robustness with respect to the positioning of the interfaces between the two grids. Finally in Sect.~\\ref{sec:applics}, we will describe some interesting astrophysical applications of this new hydrodynamical code. ",
        "conclusions": "It is well known that the migration of giant planets (i.e. planets massive enough to open significant gaps in the disk's density distribution) is governed -- or at least strongly affected -- by the global evolution of the disk under its own viscosity.  Usual simulation algorithms solve the hydrodynamical equations over a 2 dimensional polar grid that is truncated at an inner and outer radius to keep the computing time within reasonable limits. Thus, they cannot reproduce correctly the evolution of the disk, and no more the planet's migration.  In this paper, we have shown how the use of a 1D grid surrounding a classical 2D grid allows us to simulate the global evolution of the disk, while resolving the local planet-disk interactions with the accuracy of the usual algorithms. Coupling the two grids via a system of ghost rings, with special attention paid to the conservation of the angular momentum, leads to a smooth evolution of the disk from its inner radius (the truncation radius or the surface of the star) to its outer edge, several hundreds of AU away. This increases by far the accuracy of the simulation with essentially no additional CPU cost.  Consequently, this algorithm is a tool of choice to properly simulate type~II migration, and related problems. More generally, it also enables the study of the effect of the presence of planets on the disk's global evolution, which in turn affects the migration of the planets in a feedback effect.  In our simulations, we observe that the disks slowly disappear through the open boundaries of the 1D grid. However, disks are believed to disappear rapidly after only a few million years, under the action of photo-evaporation. This phenomenon could easily be introduced in our code via a simplified prescription consisting in removing a fraction of the gas in each cell, depending on time and location. This will open the possibility of simulating the evolution of planets in a globally evolving gas disk, over the disk's lifetime."
    },
    "0608/astro-ph0608456_arXiv.txt": {
        "abstract": "We report the detection of rest--frame 6.2 and 7.7~\\micron\\ emission features arising from Polycyclic Aromatic Hydrocarbons (PAH) in the Spitzer/IRS spectrum of an infrared-luminous Lyman break galaxy at z=3.01.  This is currently the highest redshift galaxy where these PAH emission features have been detected.  The total infrared luminosity inferred from  the MIPS 24~\\micron\\ and radio flux density is  2$\\times$10$^{13}$ L$_{\\odot}$, which qualifies this object as a so--called hyperluminous infrared galaxy (HyLIRG). However, unlike local HyLIRGs which are generally associated with QSO/AGNs and have  weak or absent PAH emission features, this HyLIRG has very strong 6.2 and 7.7~\\micron\\ PAH emission. We argue that intense star formation dominates the  infrared emission of this source, although we cannot rule out the presence of a deeply obscured AGN.  This LBG appears to be a distorted system in the HST ACS F606W and F814W images, possibly indicating that a significant merger or interaction is driving the large IR luminosity. ",
        "introduction": "Intensive star formation is a very common phenomenon in galaxies at high redshifts. Most high redshift galaxies are detected because of their high star formation rates, revealed by either the rest--frame UV luminosity  in Lyman break galaxies (LBGs, Steidel et al. 2003) or at far--infrared wavelengths in, for example, submillimeter--selected galaxies (SMGs, Chapman et al. 2003).  It was long debated if LBGs and SMGs constituted different coeval populations.   Recently \\citet{cha05} found that some SMGs have similar $u-g$ and $g-R$ colors as LBGs. The Spitzer Space Telescope permits exploration of the rest-frame IR properties of LBGs.   For example, Reddy et al.\\ (2006) studied the 24~\\micron\\ properties of $1.5<z<2.5$ LBGs and considered their relation to other galaxies at that epoch, including SMGs.   \\citet{hua05} reported that 5\\% of the LBGs at $z\\sim 3$ in the sample of \\citet{ste03} are detected at 24~\\micron\\, and they defined this sub-population as Infrared Luminous Lyman Break Galaxies (ILLBG).   ILLBGs have infrared properties very similar to sub--mm selected sources \\citep{ega04,hua05,pope06}.   Huang et al.\\ used Spitzer [3.6]--[8.0] and [8.0]--[24] colors to argue that the 24~\\micron\\ emission from the majority of ILLBGs is powered by intensive star formation.  Observation of high redshift galaxies  with the Infrared Spectrograph (IRS, Houck et al. 2004) on board Spitzer permits the study of their 24~\\micron\\ emission mechanism.  Early IRS studies \\citep{hou05, yan05} have focused on 24~\\micron\\ luminous (f$_{24}>$1 mJy) and optically faint sources. Most of the objects in both samples of \\citet{hou05} and \\citet{yan05}  are identified as dusty AGNs at $1.5<z<2.5$ whose spectra show power-law continua and strong silicate absorption but weak or absent emission features arising from the polycyclic aromatic hydrocarbon (PAH) molecules. \\citet{lut05} observed two lensed SCUBA sources at $z\\sim2.8$ with the IRS. Both sources are so--called Hyper--luminous IR galaxies (HyLIRGs, L$_{8-1000\\mu m} > 10^{13}$~L$_\\odot$) with strong PAH features.  For galaxies at $z>2$, the Spitzer MIPS 24~\\micron\\ band detects rest-frame mid--IR continuum emission.  At $z=2.1$ and $z=2.8$, the 6.2 and 7.7\\micron\\ PAH emission features respectively enter into the 24~\\micron\\ band. Thus, a high redshift, 24~\\micron\\-selected-sample may preferentially select objects with strong emission-line features in the mid--IR bands.   We have been studying just such a sample of LBGs, which have a redshift distribution centered at $\\langle$z$\\rangle$=3.0 \\citep{ste03}.  This letter presents the IRS spectrum and multiwavelength properties of one ILLBG at $z=3.01$ detected in the Extended Groth Strip (EGS) area.   The multi--wavelength data set from All-wavelength Extended Groth Strip International Survey (AEGIS, Davis et al. 2006)  including ACS images permits the study of Spectral Energy Distribution (SED) and morphology of this source. We adopt $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "IRS spectra of EGS20 J1418+5236, a $z=3.01$, u--dropout LBG in AEGIS show the 6.2 and 7.7\\micron\\ PAH emission features. This is currently the highest redshift galaxy where the PAH emission features have been detected. This source has a total infrared luminosity of L$_{8-1000\\mu m}$=2$\\times$10$^{13}$ L$_{\\odot}$ inferred from both the 24~\\micron\\ and radio flux densities.    Nearby galaxies with these IR luminosities that are predominantly QSO/AGN, but given the presence of strong PAH emission features and the lack of any evidence for an AGN, we conclude that star--formation dominates the emission from this z=3 ILLBG.  The 24~\\micron-to-radio flux ratio, strong IRS continuum emission, and upper limits on the flux density at 70 and 160~\\micron\\ implies that hot dust dominates the far--IR emission in this object, very similar to Mrk 231 and model SEDs of HyLIRGs. Interestingly, the theoretical models that best describe the shape and relative strength of the mid--IR SED require very high extinction values ($A_V \\gtrsim 70$~mag).  Because this object has strong UV emission, we suggest that this LBG contains a deeply embedded starburst in addition to the relatively unobscured stars that dominate the rest--frame UV optical emission.  This scenario may explain the general UV--IR properties of the ILLBG population.   This example suggests that HyLIRGs at high redshifts can appear as either ILLBGs, or SMGs, or both. Larger samples of $z\\sim 3$ galaxies with IRS spectroscopy and measurements of their far-IR SED are clearly needed to confirm these hypotheses."
    },
    "0608/nucl-th0608024_arXiv.txt": {
        "abstract": "{The neutron radius of a heavy nucleus is a fundamental nuclear-structure observable that remains elusive. Progress in this arena has been limited by the exclusive use of hadronic probes that are hindered by large and controversial uncertainties in the reaction mechanism. The Parity Radius Experiment at the Jefferson Laboratory offers an attractive electro-weak alternative to the hadronic program and promises to measure the neutron radius of $^{208}$Pb accurately and model independently via parity-violating electron scattering. In this contribution we examine the far-reaching implications that such a determination will have in areas as diverse as nuclear structure, atomic parity violation, and astrophysics. \\PACS{ {21.10.Gv}{Mass and neutron distributions}   \\and {26.60.+c}{Nuclear matter aspects of neutron stars} } % } % ",
        "introduction": "\\label{introduction}  The nucleus of ${}^{208}$Pb is 18 order of magnitudes smaller and 55 orders of magnitude lighter than a neutron star. Yet remarkably, both the neutron radius of ${}^{208}$Pb as well as the radius of a neutron star depend critically on our (incomplete) knowledge of the equation of state of neutron-rich matter. The emergence of such a correlation among objects of such a disparate size is not difficult to understand. Heavy nuclei develop a neutron-rich skin as a result of a large neutron excess and a large Coulomb barrier that reduces the proton density at the surface of the nucleus. Thus the thickness of the neutron skin depends on the pressure that pushes neutrons out against surface tension~\\cite{Brown:2000}. It is this same pressure that supports a neutron star against gravitational collapse~\\cite{Lattimer:2000nx,Steiner:2004fi}.  Thus models with thicker neutron skins often produce neutron stars with larger radii~\\cite{Horowitz:2001ya}.  Attempts at mapping the neutron distribution of ${}^{208}$Pb, or of any other heavy nucleus, using hadronic probes have been met with limited success. Although highly mature and successful, the hadronic program will never attain the precision status that the electro-weak program enjoys. This is due to the large and controversial uncertainties in the reaction mechanism~\\cite{Ray:1985yg,Ray:1992fj}. A particularly illustrative example of this situation is provided by the proton and neutron radii of ${}^{208}$Pb. While elastic electron scattering experiments have determined the charge radius of ${}^{208}$Pb to better than 0.001~fm~\\cite{Fricke:1995}, realistic estimates place the uncertainty in the neutron radius at about 0.2~fm~\\cite{Horowitz:1999fk}.  The enormously successful parity violating program at the Jefferson Laboratory~\\cite{Aniol:2005zf,Aniol:2005zg} provides an attractive electro-weak alternative to the hadronic program. Indeed, the Parity Radius Experiment (PREX) at the Jefferson Laboratory aims to measure the neutron radius of $^{208}$Pb accurately (to within $0.05$~fm) and model independently via parity-violating electron scattering~\\cite{Horowitz:1999fk}. Parity violation at low momentum transfers is particularly sensitive to the neutron density because the $Z^0$ boson couples primarily to neutrons. Moreover, the parity-violating asymmetry, while small, can be interpreted with as much confidence as conventional electromagnetic scattering experiments. PREX will provide a unique observational constraint on the thickness of the neutron skin of a heavy nucleus. We note that since first proposed in 1999, many of the technical difficulties intrinsic to such a challenging experiment have been met~\\cite{Michaels:2005}. For further details on the status of the experiment, see the contribution from Robert Michaels to these proceedings. ",
        "conclusions": "\\label{conclusions}  The neutron radius of a heavy nucleus is a fundamental nuclear-structure observable that remains elusive, mainly due to our inability to use electro-weak probes to sample the neutron distribution. While a highly mature hadronic program has been used to map the neutron distribution, the clean extraction of the neutron radius has been marred by controversial uncertainties in the reaction mechanism. The established and successful parity-violating program at the Jefferson Laboratory provides an attractive electro-weak alternative to the hadronic program. The Parity Radius Experiment at the Jefferson Laboratory will take advantage of the strong coupling of the $Z^0$ boson to neutrons to measure the neutron radius of $^{208}$Pb accurately and model independently. While the intrinsic achievement of this experiment is undeniable, in this contribution we have examined the far-reaching consequences that such a measurement could have over fields as diverse as atomic parity violation and astrophysics. First, a tight correlation was found between the neutron skin of ${}^{208}{\\rm Pb}$ (the aim of the PREX experiment) and the neutron radius of a variety of elements (Barium, Dysprosium, and Ytterbium) of relevance to atomic parity violation program. Although the calculations presented here neglect both deformation and pairing correlations, we have argued --- based on more sophisticated calculations --- that being a bulk nuclear property, the neutron radius remains largely insensitive to these effects.  Second, we demonstrated that PREX will also have a strong impact on various astrophysical observables. Indeed, a model-independent (or ``data-to-data'') relation between the neutron skin of $^{208}$Pb and the transition density from the uniform liquid mantle to the non-uniform solid crust was established. This correlation emerged as a result of the similar composition of the neutron skin of a heavy nucleus and the crust of a neutron star: neutron-rich matter at similar densities. Further, we showed how the measurement of the neutron skin in $^{208}$Pb is strongly correlated to the radius of a $M\\!=\\!1.4~M_{\\odot}$ neutron star. This result emerges as a direct consequence of the density dependence of the symmetry energy, as the same pressure that pushes neutrons out against surface tension in the nucleus of $^{208}$Pb is responsible for pushing neutrons out against gravity in a neutron star. Yet in contrast to the ``data-to-data'' relation described above, this correlation is model dependent --- as the radius of the neutron star depends on both the low- and high-density component of the equation of state. Thus, we eagerly await observational results to constrain the high-density component of the equation of state. Fortunately, earth- and space-based telescopes have started to place important constrains on the high-density component of the equation of state. New telescopes operating at a variety of wavelengths are turning neutron stars from theoretical curiosities into powerful diagnostic tools. Significant advances in observational astronomy will soon yield the first combined measurement of mass-radius relations for a variety of neutron stars. These results --- combined with the Parity Radius Experiment at the Jefferson Laboratory --- will provide the most complete information to date on the long sought equation of state of neutron-rich matter.   \\subsection*"
    },
    "0608/astro-ph0608061_arXiv.txt": {
        "abstract": "We present a pilot study of the X-ray properties of intermediate-mass ($\\sim 10^5-10^6$ \\msun) black holes in active galaxies using the \\chandra\\ X-ray telescope.  Eight of the 10 active galaxies are detected with a significance of at least $3\\,\\sigma$, with X-ray luminosities in the range $L_{0.5-2 \\mathrm{keV}} \\approx 10^{41}-10^{43}$ ergs s$^{-1}$.  The optical-to-X-ray flux ratios are consistent with expectations, given the known correlations between $\\alpha_{\\rm ox}$ and ultraviolet luminosity, while a couple of objects appear to be anomalously X-ray weak.  The range of 0.5--2 keV photon indices we measure, $1 < \\Gamma_{\\mathrm{s}} < 2.7$, is entirely consistent with values found in samples of more luminous sources with more massive black holes.  Black hole mass evidently is not a primary driver of soft X-ray spectral index.  On the other hand, we do find evidence for a correlation between X-ray power-law slope and both X-ray luminosity and Eddington ratio, which may suggest that X-ray emission mechanisms weaken at high Eddington ratio.  Such a weakening may explain the anomalous X-ray weakness of one of our most optically luminous objects. ",
        "introduction": "Astrophysical black holes (BHs) are typically found in two mass ranges: stellar-mass BHs have masses $\\sim 10$ \\msun, and are the remnants of massive star death, while supermassive BHs have masses in the range $\\sim 10^6-10^9$ \\msun, and are a ubiquitous component of galaxy bulges.  There are thus $\\sim 5$ orders of magnitude in BH mass that remains unexplored; BHs in this mass range have been dubbed ``intermediate-mass'' BHs, and their existence remains the subject of debate. In particular, there are anomalously luminous (``ultraluminous''; $L_{\\mathrm{X}} \\geq 10^{39}$ ergs s$^{-1}$) off-nuclear extragalactic X-ray sources that may be BHs with masses of $100-1000$ \\msun\\ (see, e.g., van der Marel 2004 for a review).  We have a complementary program to find intermediate-mass BHs with masses of $10^4-10^6$ \\msun\\ in active galactic nuclei (AGNs).  These objects are of interest not only because they begin to fill the mass gap between supermassive and stellar-mass BHs, but also as potential analogues of the primordial seeds of supermassive BHs. Furthermore, the mergers of BHs with masses $\\sim 10^5$~\\msun\\ are expected to provide a strong signal for the {\\it Laser Interferometric Space Antenna} (e.g,~Hughes 2002).  The tight correlations between supermassive BH mass (\\mbh) and both the luminosity (Kormendy \\& Richstone 1995) and stellar velocity dispersion (the \\msigma\\ relation; Gebhardt \\etal\\ 2000; Ferrarese \\& Merritt 2000; Tremaine \\etal\\ 2002; Barth et al. 2005; Greene \\& Ho 2006) of spheroids suggest that BHs play an important role in the evolution of bulges. On the other hand, very little is known about the prevalence of nuclear BHs in late-type, bulgeless galaxies.  Unfortunately, intermediate-mass BHs are difficult to find, because we are currently unable to resolve the gravitational sphere of influence of a $\\sim 10^5$~\\msun\\ BH outside the Local Group, and thus we cannot detect them directly through resolved kinematics.  We are forced to rely on indirect evidence of the presence of a BH from radiative signatures. Two prototypical low-mass AGNs are known, in the late-type spiral NGC 4395 (Filippenko \\& Ho 2003) and in the dwarf elliptical galaxy POX 52 (Barth \\etal\\ 2004).  Greene \\& Ho (2004) systematically defined a sample of similar objects using the First Data Release of the Sloan Digital Sky Survey (SDSS; Abazajian \\etal\\ 2003).  Their sample of 19 galaxies forms the parent sample for the present paper, and throughout we will refer to the sample using the identification numbers from Greene \\& Ho (2004).  The Greene \\& Ho objects represent a very homogeneously selected sample of BHs with low masses, and thus presents an ideal sample to investigate how the broad-band spectral properties of AGNs depend on BH mass.  The spectral properties of intermediate-mass BHs are important not only for the insight they may provide into accretion processes, but also to place observational constraints on the radiative properties of ``mini-quasars'' that may have contributed significantly to the reionization of the Universe (e.g,~Madau \\etal\\ 2004).  We are thus in the process of measuring the multiwavelength properties of the sample. Using the Very Large Array, we found that the objects are very faint at in the radio (Greene et al. 2006).  Here we present the results of a pilot study to constrain the X-ray properties of the sample.  Throughout we assume the following cosmological parameters to calculate distances: $H_0 = 100~h = 71$~\\kms~Mpc$^{-1}$, $\\Omega_{\\rm m} = 0.27$, and $\\Omega_{\\Lambda} = 0.75$ (Spergel \\etal\\ 2003).  \\begin{figure*} \\vbox{ \\vskip +0.2truein \\hbox{ \\psfig{file=table1_v7.epsi,width=0.37\\textwidth,keepaspectratio=true,angle=90} }} \\end{figure*}  \\vskip 5mm ",
        "conclusions": "We present a pilot study of the X-ray properties of intermediate-mass BHs in active galaxies.  We have detected eight out of the 10 objects surveyed with 5 ks \\chandra\\ observations.  From an analysis of their hardness ratios, we find that the objects have a range in 0.5--2 keV photon index of $1<$\\gams$<3$ [$N(E) \\propto E^{-\\Gamma_{\\mathrm{s}}}$], in keeping with general AGN samples previously studied.  In the five objects for which there are sufficient counts to extract spectra, there is no evidence for neutral absorption, although we do see evidence for spectral complexity at energies $\\leq 1$ keV, which may indicate ionized absorption.  The X-ray-to-optical flux ratios for the majority of the sample obey the same correlation with optical/UV luminosity seen in the literature, but there are two X-ray-weak sources.  Given the apparent correlation between Eddington ratio and spectral slope, we speculate that X-ray weakness in at least one of these sources may be related to a high Eddington ratio.  We find that the soft X-ray slope correlates strongly with soft X-ray luminosity, but not with BH mass.  This study demonstrates the feasibility of detecting $\\sim 10^5-10^6$~\\msun\\ BHs with relatively short \\chandra\\ observations.  With them, we have begun to probe the accretion properties of AGNs in an unexplored mass regime.  It is our goal to build broad spectral energy distributions for the entire sample, in order to properly measure their bolometric luminosities.  We will then be in a good position to search for mass-dependent spectral properties, as well as calibrate various common average bolometric corrections (such as \\oiii\\ luminosity) for objects in this mass regime.  Finally, while we are unable to address the variability properties of the sample with short observations, there are strong indications that the X-ray variability amplitude (e.g.,~O'Neill \\etal\\ 2005) and detailed power spectrum (McHardy \\etal\\ 2005) can serve as indicators of the BH mass. Extended X-ray observations of this sample would provide a powerful new test of these important BH mass measurement techniques."
    },
    "0608/astro-ph0608127_arXiv.txt": {
        "abstract": "{We present arguments based on the measured abundances in individual damped Lyman $\\alpha$ systems (DLAs) and sub-damped Lyman $\\alpha$ systems (sub-DLAs), and also the average abundances inferred in large samples of QSO absorption line systems, to suggest that the amount of dust in intervening QSO absorbers is small and is not responsible for missing many QSOs in magnitude limited QSO surveys. While we can not totally rule out a bimodal dust distribution with a population of very dusty, metal rich, absorbers which push the background QSOs below the observational threshold of current optical spectroscopic studies, based upon the current samples it appears that the metallicity in QSO absorbers decreases with increase in H I column densities beyond $10^{19}$ cm$^{-2}$. Thus the sub-DLA population is more metal rich than the DLAs, a trend which may possibly extend to the non-damped Lyman limit systems (NDLLS).  Based on the recently discovered mass-metallicity relation for galaxies, we suggest that most sub-DLAs and possibly NDLLS, are associated with massive spiral/elliptical galaxies while most DLAs are associated with low mass galaxies.  The sub-DLA galaxies will then contribute a larger fraction of total mass (stellar and ISM) and therefore metals, to the cosmic budget, specially at low redshifts, as compared to the DLAs.} ",
        "introduction": "Damped Lyman $\\alpha$ systems (DLAs), having column density of {H~I} (N$_{\\rm H\\;I}$) $\\ge\\;2\\times 10^{20}$ cm$^{-2}$ have been suggested to be the high-redshift analogs of disks of nearby luminous galaxies. In spite of having a high amount of neutral hydrogen these absorbers have very little H$_2$ (Ledoux et al. 2003), which raises questions about their being associated with high star formation activity. A large fraction of DLAs do indeed appear to have low star formation rates based on deep emission-line imaging searches (e.g. Kulkarni et al. 2006a and references therein).  Metallicity measurements, using intermediate to high resolution observations, are available for $>$ 100 DLAs. These indicate that the metallicity evolution of DLAs is weak; most of the DLAs, even at $z_{abs}\\sim0$ are found to be metal poor (Kulkarni et al. 2005). It is thus possible that the DLAs may not trace the bulk of star formation in the Universe and therefore may not be among the leading metal carriers.  Compilation of observed metallicities of individual DLA systems shows a trend of decreasing abundance with increasing H I column density (Boisse et al. 1998; Akerman et al. 2005; Meiring et al. 2006a). The trend possibly exists to lower H I column densities covering the sub-DLAs, having N$_{\\rm H\\;I}$ between 10$^{19}$ and $2\\times 10^{20}$ cm$^{-2}$ (P\\'eroux et al. 2003). A similar conclusion was drawn by York et al. (2006, hereafter, Y06) based on the average abundances inferred from the composite spectra of sub-samples drawn from a sample of 809, intervening Mg II QSO absorption line systems (rest equivalent width of $\\lambda$2796 $>$ 0.3 {\\AA} and 1 $<z_{abs}<$ 1.86) compiled from the Sloan Digital Sky Survey (SDSS) Data Release 1 (DR1). The sub-sample with highest estimated (near solar) Zn metallicity, was found to have a low average inferred N$_{\\rm H\\;I}$ of $\\sim10^{20}$ cm$^{-2}$. Several arguments were given in Y06 to suggest that the DLAs may not be the analogs of the local bright galaxies, rather, the sub-DLAs and possibly the non-DLA Lyman limit systems (NDLLS) with N$_{\\rm H\\;I}$ between 10$^{17}$ and 10$^{19}$ cm$^{-2}$ may be associated with large galaxies. Here we consider this scenario further, using the currently available data on individual absorbers and the results of Y06. We consider various selection effects at work and discuss the implications of the recently established mass-metallicity relation for the nature of DLAs and sub-DLAs.  In section 2 we study the dependence of the Zn abundance on the column density of H I, for the observations of individual systems taken from the literature and for the average values of these quantities obtained by Y06 for large samples. We present other relevant observational information and discuss the obscuration threshold and other selection effects to argue for the reality of the observed dependence.  The implication of this dependence for the nature of the absorbing galaxies are presented in section 3 and conclusions are presented in section 4. ",
        "conclusions": "We have studied the dependence of metallicity on N$_{\\rm H\\;I}$ in QSO absorbers with N$_{\\rm H\\;I}>10^{19}$ cm$^{-2}$ and have discussed various selection effects that may be giving rise to the observed trend. We have argued that the selection effects are not important and that the observed trend is real. Our conclusions based on these arguments, and subject to confirmation by future observations, are as follows: \\begin{enumerate} \\item{} The amount of dust in QSO absorbers is small and is not responsible for missing many QSOs in magnitude limited surveys. We can not, however, rule out the possibility of a bimodal distribution of dust columns such that there may exist a population of dusty absorbers which push the background QSOs below the observational threshold of current optical spectroscopic studies and is completely invisible. \\item{} The metallicity in QSO absorbers with N$_{\\rm H\\;I}>10^{19}$ cm$^{-2}$, decreases with increase in H I column density of these absorbers. The trend possibly continues to lower H I column densities.  The sub-DLAs thus have higher metal abundances as compared to the DLAs at redshifts between 0 and 2. \\item{} The observed mass-metallicity relation suggests that most DLAs are associated with low mass ($<10^9$ M$_\\odot$) galaxies while most sub-DLAs are associated with massive spiral/elliptical galaxies.  It is possible that the non-DLA LLS may also be metal rich and may be associated with massive galaxies. \\item{} The sub-DLA galaxies will contribute a larger fraction of total mass (stellar and ISM) and therefore metals, to the cosmic budget, specially at low redshifts, as compared to the DLAs. The Sub-DLAs and possibly, non-DLA LLS galaxies, may contain a much larger fraction of the metals at $z< 1$ than has been appreciated. \\item{} The few imaging studies, of galaxies responsible for quasar absorption line systems done so far, are ambiguous on the morphology of DLA and sub-DLA galaxies. More systematic, deep imaging in r-band, Ks-band and in narrow emission lines is essential to confirming the inferences of this paper. \\end{enumerate}"
    },
    "0608/astro-ph0608311_arXiv.txt": {
        "abstract": "``Rotating RAdio Transients'' (RRATs) are a newly discovered astronomical phenomenon, characterised by occasional brief radio bursts, with average intervals between bursts ranging from minutes to hours.  The burst spacings allow identification of periodicities, which fall in the range 0.4 to 7 seconds.  The RRATs thus seem to be rotating neutron stars, albeit with properties very different from the rest of the population. We here present the serendipitous detection with the {\\em Chandra X-ray Observatory}\\ of a bright point-like X-ray source coincident with one of the RRATs. We discuss the temporal and spectral properties of this X-ray emission, consider counterparts in other wavebands, and interpret these results in the context of possible explanations for the RRAT population. ",
        "introduction": "\\label{intro}  Astronomers are still coming to terms with the fact that the zoo of isolated neutron stars harbours an increasingly diverse population. Objects of considerable interest include radio pulsars, anomalous X-ray pulsars, soft gamma repeaters, central compact objects in supernova remnants (SNRs), and dim isolated neutron stars. A startling new discovery has now forced us to further expand this group: \\cite{mll+06} have recently reported the detection of eleven ``Rotation RAdio Transients'', or ``RRATs'', characterised by repeated, irregular radio bursts, with burst durations of 2--30~ms, and intervals between bursts of $\\sim4$~min to $\\sim3$~hr. The RRATs are concentrated at low Galactic latitudes, with distances implied by their dispersion measures of $\\sim2-7$~kpc.  For ten of the eleven RRATs discovered by \\cite{mll+06}, an analysis of the spacings between repeat bursts reveals an underlying spin period, $P$, and also in three cases, a spin period derivative, $\\dot{P}$. The observed periods fall in the range 0.4~s~$<P<$~7~s, which generally overlap with those seen for the radio pulsar population.  For the three RRATs with values measured for both $P$ and $\\dot{P}$, a characteristic age, $\\tau_c$, and a dipole surface magnetic field, $B$, can both be inferred, as listed in Table~\\ref{tab_rrats}. These sources can also be placed on the standard ``$P-\\dot{P}$ diagram'', and can thus be compared to other populations of rotating neutron star: \\begin{itemize} \\item RRAT~J1317--5759 has properties typical of radio pulsars, except for a relatively long spin period. \\item RRAT~J1819--1458 is very young, with a high magnetic field. On the $P-\\dot{P}$ diagram, it is located in the upper-right region, in the same area occupied by the magnetars and by the high-field radio pulsars. \\item RRAT~J1913+3333 has spin properties indistinguishable from the bulk of radio pulsars. \\end{itemize}  We here present the X-ray detection of RRAT~J1819--1458, which provides a further point of comparison with the various neutron star population. The details of this detection and its interpretation are discussed by \\cite{rbg+06}.  \\begin{table}[hbt] \\caption{Properties of the three RRATs with measured period derivatives.} \\centering \\label{tab_rrats}       % \\begin{tabular}{lccc} \\hline\\noalign{\\smallskip} Source & $P$ (sec) & $\\tau_c$ (Myr) & $B$ ($10^{12}$~G)  \\\\[3pt] \\tableheadseprule\\noalign{\\smallskip} RRAT~J1317--5759 & 2.6 & 3.3 & 5.83 \\\\ RRAT~J1819--1458 & 4.3 & 0.12 & 50 \\\\ RRAT~J1913+3333 & 0.92 & 1.9 & 2.7 \\\\ \\noalign{\\smallskip}\\hline \\end{tabular} \\end{table} ",
        "conclusions": "A variety of forthcoming observations should be able to cast more light on the possibilities discussed above, and on the spatial and spin distributions of these sources. Several groups have already undertaken new radio searches for RRAT-like emission from other classes of neutron stars. Near-infrared observations with the VLT have also been obtained for some RRATs, to provide stronger constraints on the X-ray to infrared flux ratio of these sources, and to look for hints of a fossil / fallback disk \\citep[see][]{wck06}. Meanwhile, in a project called ``Astro Pulse'', the data from the Arecibo SETI@Home project are being reanalysed for short transient signals. With the sensitivity of Arecibo, many RRATs may be found in this analysis.  Finally, following on from our X-ray detection discussed here, {\\em XMM-Newton}\\ observations of two RRATs have been approved for Cycle~5: of RRAT~J1819--1458 (to obtain a better spectrum and to properly search for pulsations) and of RRAT~J1317--5759 (to see if it too emits detectable X-rays).  A final point to make is that it is already clear from the 11 RRATs known that the birth rate of these objects is $\\sim3-4$ times that previously estimated for all radio pulsars \\citep{mll+06,ptp06}. We are thus only seeing the upper tip of a large transient distribution. When one combines this with the populations of nullers \\citep{bac70}, ``bursters'' \\citep[e.g., PSR~J1752+2359;][]{lwf+04}, ``winkers'' \\citep[PSR~B1931+24;][]{klo+06} and ``burpers'' \\citep[GCRT J1745-3009;][]{hlk+05}, it becomes clear that a full study of the transient radio sky is needed \\citep{clm04}. The next generation of radio telescopes, with the capability of monitoring very wide fields of view (e.g., LOFAR, xNTD and the SKA) should make many further discoveries of such phenomena."
    },
    "0608/astro-ph0608477_arXiv.txt": {
        "abstract": "Using the Wide Field Grism Spectrograph 2 (WFGS2), we have carried out slit-less spectroscopy, $g'r'i'$ photometry, and slit spectroscopy on the L1014 dense core.  We detected three H$\\alpha$ emission line stars.  We interpret one as weak-line T Tauri star (WTTS) and the others as classical T Tauri stars (CTTS). Since their $g'-i'$ colors and/or classified spectral types are consistent with those of T Tauri stars and two of them show less extinction than the cloud, these three stars are likely to be T Tauri stars associated with L1014. Adopting an age range for T Tauri stars, 1--10 Myr, the  color-magnitude diagram suggests a  distance of  $\\sim$ 400--900 pc, rather than the previously assumed distance, 200 pc. This could strongly affect on the mass estimate  of L1014-IRS, which is thought to be either a very young protostar or proto-brown dwarf. ",
        "introduction": "In the early evolution of substellar objects, the formation of brown dwarfs in molecular clouds is of particular interest; several formation mechanisms have been proposed (e.g., \\cite{ML06, luhman06, whitworth06}). One possibility is the same mechanism as low-mass stars, but with much smaller cores (e.g., \\cite{PN04}). The L1014 dense core is a potential site for such a process.  L1014 was previously thought to be starless because it lacks IRAS point sources  \\citep{LM99}.   \\citet{yng04} discovered a very faint infrared source (L1014-IRS) toward the center of L1014 with the Spitzer Space Telescope. Recent detections of a molecular outflow and a reflection nebula revealed that this central object is really embedded within the cloud core and is not a background source associated with the Perseus arm at 2.6 kpc \\citep{bourke05, huard06}. Assuming the distance of L1014 is $\\sim$ 200 pc, as adopted in recent literature, this embedded source becomes an ideal candidate for a very low mass protostar, i.e., a proto-brown dwarf.  However, the mass estimate of a protostellar object depends on its adopted distance. It is, in general,  not easy to determine the distance to  the surrounding dark cloud, particularly  that of a small, nearby dense core as  its small size implies a small number of background/foreground stars. In fact, the assumption that L1014 is 200 pc away seems to originate from the distance estimate  of B362, which is located $\\sim10\\arcmin$ north of L1014, but this estimate is based upon the similarity of $V_{LSR}$ \\citep{dieter1973} to that of nearby clouds with a kinematic distance $\\sim$ 100--200 pc \\citep{dickman1976}.   The L1014 dense core is thought to be a site of very  recent  ($<$ 1 Myr) star formation. We interpret the slightly extinct region extending southwest of L1014 in the DSS-II red image as thinner molecular gas extending from the dense core.  In fact, CO observations \\citep{RP92, crapsi05} showed the existence of molecular material in this region. We suspected that the thin molecular gas and the dense core are remnants of previous low-mass star formation over the past few Myr. If so, pre-main-sequence stars, i.e., T Tauri stars,  could be detected toward L1014 and the validity of the distance estimate of L1014 can be examined with the color-magnitude diagram by using theoretical isochrones of pre-main-sequence stars.   In order to substantiate the presence of previous star formation, we surveyed H$\\alpha$ emission line stars by slit-less spectroscopy around this remarkable core. Slit spectroscopy and $g'r'i'$ photometry were then performed on the selected stars. In this letter, we present our spectroscopic results and discuss whether the detected H$\\alpha$ emission line stars are T Tauri stars associated with the L1014 core. We also discuss the distance to L1014. ",
        "conclusions": "We identified three T Tauri stars toward the L1014 core by slit-less spectroscopy, $g'r'i'$ photometry and slit spectroscopy. One of them is a WTTS and the others are CTTS. The estimation of the extinction and the ages of T Tauri stars suggests that they are associated with L1014. This indicates that star formation took place previously in L1014's  progenitor cloud.  An examination of the color-magnitude diagram suggests that the distance to L1014 is $\\sim$ 400--900 pc. Since the previously assumed distance of 200 pc was based on the kinematic distance, our estimated distance is plausible if these stars are associated with L1014.  The luminosity of the very low mass protostar formed in the L1014 core is reported to be $\\sim$ 0.09 $L\\odot$, for a distance of 200 pc \\citep{yng04}. Assuming a distance of 400--900 pc, the mass estimate of the protostellar object (L1014-IRS) significantly increases.   \\bigskip  We are grateful to the staff of the UH 2.2-m telescope for supporting our WFGS2 observations. We also appreciate the great support of the Subaru Telescope. Use of the UH 2.2-m telescope for the observations is supported by the National Astronomical Observatory of Japan. We thank  Joel F. Koerwer for proof reading. A. M. thanks the Daiko Foundation for their financial support. WFGS2 was supported by the Daiko Foundation and by the Sumitomo Foundation. This instrument was also supported by Grants-in-Aid from the Ministry of Education, Culture, Sports, Science and Technology (14540228 and 15540238). This work was financially supported in part by Grants-in-Aid from the Ministry of Education, Culture, Sports, Science and Technology (16740121 and 17039011)."
    },
    "0608/astro-ph0608688_arXiv.txt": {
        "abstract": "We calculate the evolution of heavy element abundances from C to Zn in the solar neighborhood adopting our new nucleosynthesis yields. Our yields are calculated for wide ranges of metallicity ($Z=0-Z_\\odot$) and the explosion energy (normal supernovae and hypernovae), based on the light curve and spectra fitting of individual supernovae. The elemental abundance ratios are in good agreement with observations. Among the $\\alpha$-elements, O, Mg, Si, S, and Ca show a plateau at [Fe/H] $\\ltsim -1$, while Ti is underabundant overall. The observed abundance of Zn ([Zn/Fe] $\\sim0$) can be explained only by the high energy explosion models, which requires a large contribution of hypernovae. The observed decrease in the odd-Z elements (Na, Al, and Cu) toward low [Fe/H] is reproduced by the metallicity effect on nucleosynthesis. The iron-peak elements (Cr, Mn, Co, and Ni) are consistent with the observed mean values at $-2.5 \\ltsim$ [Fe/H] $\\ltsim -1$, and the observed trend at the lower metallicity can be explained by the energy effect. We also show the abundance ratios and the metallicity distribution functions of the Galactic bulge, halo, and thick disk. Our results suggest that the formation timescale of the thick disk is $\\sim 1-3$ Gyr. ",
        "introduction": "Elemental abundance ratios are the treasure house of information on star formation and galaxy formation because different types of supernovae produce different heavy elements with different timescales \\citep[e.g.,][]{tin80,pagel1997,matteucci2001}. High resolution spectroscopy gives the elemental abundance patterns of individual stars in our Galaxy and nearby dwarf spheroidal galaxies, and of high-redshift quasar absorption line systems. These observations have shown different abundance patterns, which suggest different chemical enrichment histories of these objects and probably different initial mass functions (IMFs) depending on environment. The re-ionization of the universe at very early epoch as suggested by the WMAP result may require a different population for the first stars \\citep[e.g.,][]{cir03}. However, to discuss these issues, nucleosynthesis yields have involved too many inconsistencies with recent observations.  There exist two distinct types of supernova explosions \\citep[e.g.,][]{arnett1996,alex1997}: One is Type II supernovae (SNe II), which are the core collapse-induced explosions of massive stars ($\\gtsim \\, 8M_\\odot$) with short lifetimes of $10^{6-7}$ yrs, and produce more $\\alpha$-elements (O, Mg, Si, S, Ca, and Ti) relative to Fe with respect to the solar ratios (i.e., [$\\alpha$/Fe] $>0$). The other is Type Ia supernovae (SNe Ia), which are the thermonuclear explosions of accreting white dwarfs (WDs) in close binaries and produce mostly Fe and little $\\alpha$-elements \\citep[e.g.,][]{nom94}. Although the lifetime of SNe Ia has been estimated to be the order of Gyr from the chemical evolution of our Galaxy (e.g., Yoshii, Tsujimoto \\& Nomoto 1996), the simulations of the SN Ia progenitor systems predict that the lifetime of the majority of SNe Ia is shorter than this (Hachisu, Kato, \\& Nomoto 1996; Kobayashi et al. 1998, hereafter K98). The metallicity effect of SNe Ia that is also predicted by the simulations can solve this problem (K98).  Following the development of observations of individual supernovae and elemental abundances, theory of stellar nucleosynthesis has been improved from Woosley \\& Weaver (1995) and Nomoto et al. (1997a, hereafter N97). From the light curve and spectral fitting of individual supernovae, the mass of progenitors stars $M$, explosion energy $E$, and produced $^{56}$Ni mass (which decayed to $^{56}$Fe) have been obtained (e.g., Nomoto et al. 2004 for a review). There exist two distinct types of core-collapse supernovae: One is normal SNe II (including Ib and Ic), which have the explosion energy of $E_{51} \\equiv E/10^{51}$ erg $\\sim 1$ and produce little iron relative to $\\alpha$-elements. Some SNe II with $M>25M_\\odot$ may have smaller energy ($E_{51}<1$), which are called ``faint SNe'' \\citep{tur98}. The other is hypernovae (HNe), which have more than ten times larger explosion energy ($E_{51}\\gtsim10$) and produce a certain amount of iron. HNe may be related to gamma-ray bursts \\citep{gal98,iwa98}.  Umeda \\& Nomoto (2002, 2005, hereafter UN05) updated the progenitor star models with a metallicity range of $Z=0-0.02$ (see also Limongi \\& Chieffi 2003), and calculated the explosive nucleosynthesis yields for larger energies ($E_{51}>1$). To meet the observed abundance patterns for extremely metal-poor (EMP) stars \\citep{cay04,aoki2004,beers2005}, they introduced the mixing and fallback mechanism. They then succeed in reproducing the observed abundance patterns of some EMP stars (Umeda \\& Nomoto 2003, hereafter UN03; UN05), based on the idea that the interstellar medium is not mixed and the EMP stars are enriched only by a single supernova \\citep{aud95}. The trends of iron-peak element abundance patterns for the EMP stars \\citep{mcw95,rya96} have been explained by changing a parameter mass-cut \\citep{nak99}, while they were explained with the energy dependence: [(Zn,Co)/Fe] increases and [(Mn,Cr)/Fe] decreases for metal poor stars because of higher energy and a larger amount of swept hydrogen mass (UN05).  In this paper, however, we focus on typical yields of core-collapse supernovae that is responsible for the plateau values of the [$\\alpha$/Fe]-[Fe/H] relations, which is important when we discuss the galaxy formation and evolution. Physical parameters in our model are constrained from observations, i.e., the light curve and spectra of individual supernovae and the abundance patterns of metal poor stars in the solar neighborhood. There may be a scatter in the observed mass-energy relation and [$\\alpha$/Fe]-[Fe/H] relations that arise from inhomogeneity of the explosion and of the mixing of the interstellar medium. However, at $-2.5 \\ltsim$ [Fe/H] $\\ltsim -1$, chemical enrichment of the interstellar medium is proceeded with many supernova explosions, and the average contribution of supernovae with different mass is imprinted in the [X/Fe]-[Fe/H] diagrams. First, we give a table set of nucleosynthesis yields both for SNe II and HNe as functions of progenitor mass and metallicity (\\S 2). We then show the evolution of elemental abundance ratios from carbon to zinc in the solar neighborhood using a one-zone chemical evolution model (\\S 3). In \\S 4, we construct models for the bulge, halo, and thick disk of the Milky Way Galaxy to meet their metallicity distribution functions, and show the age-metallicity relations. Showing the different evolution of elemental abundance ratios, we extend our discussion to the formation of the thick disk. \\S 5 gives our conclusions. ",
        "conclusions": "We calculate the evolution of heavy element abundances from C to Zn in the solar neighborhood adopting our new nucleosynthesis yields. Our new yields are based on the new developments in the observational/theoretical studies of supernovae and extremely metal-poor stars in the Galactic halo. We use the light curve and spectra fitting of individual supernova to estimate the mass of the progenitor, explosion energy, and produced $^{56}$Ni mass.  The elemental abundance ratios are in good agreement with observations. Figure \\ref{fig:imf} provides a summary of our new yield table set. The solid and long-dashed lines show the abundance patterns in our chemical evolution model for the solar neighborhood at [Fe/H] $=0$ and $-1.1$, respectively, which should correspond to the solar abundance [X/Fe] $=0$ and the IMF weighted SN II yield without SN Ia contribution, respectively. The metallicity and energy dependencies of the yields are demonstrated by the short-dashed and dotted lines, which show the IMF-weighted yield with $Z=0$ and the HN yield with $M=20M_\\odot, E_{51}=10, Z=0$, respectively.  \\begin{enumerate} \\item The solar abundance (i.e., [X/Fe] $=0$) is basically well reproduced with our chemical evolution model at [Fe/H] $=0$ (solid line). Some of the discrepancies between our model and observations may be explained as follows:  \\noindent (1) The underabundances of C and N may suggest that AGB stars and Wolf-Rayet star winds are the dominant sources of these elements.  \\noindent (2) F, K, Sc, Ti, and V are underabundant, which cannot be increased by changing our parameters such as metallicity and energy. A jet-like explosion \\citep{mae03} can efficiently increase Ti and Zn abundances and a low-density model (UN05) or neutrino process \\citep{fro06,wan06,yos06} may increase Sc abundance.  \\noindent (3) Ni overproduction by SNe Ia can be reduced \\citep{iwa99}.  \\item The observed plateau values at $-1.5 \\ltsim$ [Fe/H] $\\ltsim -1$ \\citep[dots, ][]{sne91,mel02,gra03} are in good agreement with our model at [Fe/H] $=-1.1$ (long-dashed line), where SNe Ia do not contribute. Only Ti, Sc and V are underabundant by 0.4, 0.6 and 0.2 dex, respectively, relative to Fe.  \\item The observed scatter at $-3.5 \\ltsim$ [Fe/H] $\\ltsim -2.5$ (errorbars) and the observed trends toward [Fe/H] $\\sim -4$ \\citep[arrows,  ][]{mcw95,rya96,cay04,hon04} may be due to {\\sl inhomogeneous} enrichment, and could be explained with the variations of the properties of individual supernovae such as energy and metallicity (UN05).  \\item The metallicity effect is seen in our IMF-weighted yield of metal-free stars (short-dashed line), where the abundance ratios of the odd-Z elements [(Na, Al, Mn, Cu, ...)/Fe] are smaller than those of the even-Z elements by $\\sim 0.6$ dex. The observed trends (arrows) are well reproduced with this metallicity effect in our chemical evolution model.  \\item The energy effect is seen in our HN yield for $M=20M_\\odot, E_{51}=10$, and $Z=0$ (dotted line), where the abundance ratios of iron-peak elements [(Cr, Mn, Co, and Zn)/Fe] are different from normal SNe II. The observed trends (arrows) could be explained with the variation of the explosion energy of individual supernovae.  \\item  [$\\alpha$/Fe] (O, Mg, Si, S, Ca, and Ti) depends on the mass of the progenitor star and is larger for massive SNe II. For HNe, however, because the iron yield is as large as the yields of $\\alpha$-elements, [$\\alpha$/Fe] is almost constant. This may account for the observed small scatter of [$\\alpha$/Fe] in the solar neighborhood, being independent of the mixing process of interstellar medium. The small [$\\alpha$/Fe] (O, Mg, Si, S, Ca, and Ti) in some anomalous stars is due to i) SNe Ia, ii) relatively large Fe production from low-mass SNe II with $M=13-15M_\\odot$, iii) large Fe production from HNe, or iv) the SN I.5 explosion of some AGB stars (see \\S \\ref{sec:xfe}). The large [$\\alpha$/Fe] is due to i) small Fe production from SNe II with $E_{51} \\le 1$, or ii) large fallback mass for HNe. From Zn and iron-peak elements, we can distinguish the enrichment source of such stars; HNe produce large [(Zn,Co)/Fe], and SNe Ia and I.5 produce relatively large [Mn/Fe].  \\end{enumerate}"
    },
    "0608/astro-ph0608531_arXiv.txt": {
        "abstract": "We present the results of a MOPED analysis of $\\sim 3 \\times 10^5$ galaxy spectra from the Sloan Digital Sky Survey Data Release Three (SDSS DR3), with a number of improvements in data, modelling and analysis compared with our previous analysis of DR1.  The improvements include: modelling the galaxies with theoretical models at a higher spectral resolution of 3\\AA; better calibrated data; an extended list of excluded emission lines, and a wider range of dust models. We present new estimates of the cosmic star formation rate, the evolution of stellar mass density and the stellar mass function from the fossil record. In contrast to our earlier work the results show no conclusive peak in the star formation rate out to a redshift around 2 but continue to show conclusive evidence for `downsizing' in the SDSS fossil record. The star formation history is now in good agreement with more traditional instantaneous measures. The galaxy stellar mass function is determined over five decades of mass, and an updated estimate of the current stellar mass density is presented. We also investigate the systematic effects of changes in the stellar population modelling, the spectral resolution, dust modelling, sky lines, spectral resolution and the change of data set. We find that the main changes in the results are due to the improvements in the calibration of the SDSS data, changes in the initial mass function and the theoretical models used. ",
        "introduction": "The quality of spectra of the observed light of unresolved stellar populations has reached enough accuracy that it is possible to make detailed studies of the physical properties of the stellar populations in these galaxies. An excellent example of this new generation of data-sets is given by the Sloan Digital Sky Survey \\citep{Gunn+98,York+00,Strauss+02} at low redshift, not only by the size of the spectroscopic sample (about $10^6$ spectra) but by the quality and wavelength coverage of the spectra. At higher redshift the DEEP2 survey \\citep{Davis+03} is providing a similar database, albeit with a  smaller wavelength coverage. Future spectroscopic surveys (e.g. WFMOS) will yield larger samples at even deeper redshifts. Given the quality of the spectra, it is interesting to ask the question of whether reliable information about the stellar population of these galaxies can be inferred from the spectra.  Indeed, several attempts have been made previously to study in detail the physical properties of the SDSS galaxies either by using selected features in the spectra \\citep{Kauffmann+04,Brinchmann04, Tremonti+04} or using the full spectrum \\citep{ PHJ03,HPJD04,PHJ04, Cid05,Mathis06,Ocvirk06}. These studies have led to interesting conclusions about the physical properties of these galaxies. In particular analysis of the SDSS sample \\citep{HPJD04} and other local galaxies \\citep{Thomas05} show very clear evidence for `downsizing' - the process by which star formation at low redshift takes place predominantly in low-mass galaxies, whereas more massive galaxies have the bulk of their star-formation activity at high redshift. also observed using other methods. In addition one can also determine the global star formation history in the Universe from this fossil record.  Our previous study found broad agreement with observations of contemporary star formation at large lookback times, but suggested that there was a peak in star formation activity at $z<1$. Agreement between the fossil record and contemporary star formation indicators is expected if the Cosmological Principle holds; this expected agreement offers an opportunity in principle to test the assumptions in the modelling of both the fossil record and the instantaneous star formation rates. Of most importance is that both are dependent on the stellar initial mass function, but in different ways, with the instantaneous rates being determined very much by the upper end of the IMF, while the fossil record is determined by a wider range of stars.  This offers an opportunity in principle to test the evolution of the IMF as a function of cosmic time.  In our previous work, we have presented results from around $10^5$ galaxies from the SDSS DR1.  This paper enhances the previous results through the introduction of a number of improvements to data and analysis methods, and through an investigation of systematic effects and sensitivity to assumptions.  On the data side, we have analysed $\\sim 3 \\times 10^5$ galaxies from the SDSS DR3 sample. The extra size of this sample is not particularly important, but the calibration of the data has been improved since DR1.  On the methods side, we no longer rebin the data to 20\\AA\\ resolution, but compare the data with a wider range of theoretical models, including the Bruzual \\& Charlot models \\citep{BC03} at 3\\AA\\ resolution.  We consider two stellar initial mass functions (a Salpeter and a Chabrier IMF), and a wider range of dust models. We also make some improvements to the treatment of emission lines, by removing additional weaker lines, since interstellar emission lines are not included in the stellar modelling.  Finally, we have also explored the effect of removal of sky lines, using a PCA-based method \\citep{Wild05}. These studies give us a reasonable estimate of the sensitivity of the results to the assumptions.  As is expected from a sample of this size, changes in the assumptions give rise to much larger variations in the results than the statistical errors. The uncertainties in star formation rates from the stellar models used are typically about a factor of two, which makes it difficult to distinguish between different IMFs, although it may be possible to constrain extreme IMF variations.  The layout of the paper is as follows.  In Section 2 we outline the enhancements to the method, and the improvements to the data, and describe the assumptions which are made for analysis of the DR3 sample.  In Section 3 we present the new results on star formation history, downsizing and galaxy stellar mass function from the DR3 dataset. We also introduce a method to assess the particular areas where current models are lacking. In Section 4 we show the sensitivity of our results to changes in the assumptions, through analysis of a subset of the data, and in Section 5 we draw conclusions.  Throughout this work we assume a concordance cosmology with $\\Omega_v=0.73$, $\\Omega_m=0.27$, $H_0=71$ kms$^{-1}$Mpc$^{-1}$\\citep{SpergelWMAP03}. ",
        "conclusions": "We have used MOPED to probe the fossil record of star formation encoded in the spectra of more than $300,000$ SDSS-DR3 galaxies to determine stellar populations, metallicity evolution and dust content. We have also investigated the impact of systematics on recovering physical information from the fossil record.  Our main conclusions are: \\begin{itemize} \\item The main impacts on systematic variation in the estimation of SFR from the fossil record are due to the the stellar population model, the calibration of the observed spectra and the choice of the IMF.  \\item We find strong evidence for downsizing, independent of model choice.  \\item The overall star formation history of galaxies recovered from the fossil record agrees well with instantaneous formation measurements.  \\item The mass build-up recovered from our analysis is in good agreement with that predicted from both high redshift studies and the semi-analytic simulations of galaxy formation based on Millennium Run.  \\item We have identified the cause of many of the residuals from the spectral fits. Stellar population models that provide extra freedom in terms of alpha enhancement should provide better fits. \\end{itemize}    The fossil record continues to provide a useful tool to unveil the physical conditions of galaxies in the present and the past. With improved models that incorporate alpha-enhancement it should be possible to constrain further models of galaxy formation and evolution and the initial mass function of galaxies."
    },
    "0608/gr-qc0608095_arXiv.txt": {
        "abstract": "We have demonstrated displacement- and frequency-noise free laser interferometry (DFI) by partially implementing a recently proposed optical configuration using bi-directional Mach-Zehnder interferometers (MZIs).  This partial implementation, the minimum necessary to be called DFI, has confirmed the essential feature of DFI: the combination of two MZI signals can be carried out in a way that cancels displacement noise of the mirrors while maintaining gravitational wave signals. The attained maximum displacement noise suppression was 45\\,dB. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608707_arXiv.txt": {
        "abstract": "LS~I+61 303 has been detected by the Cherenkov telescope MAGIC at very high energies, presenting a variable flux along the orbital motion with a maximum clearly separated from the periastron passage. In the light of the new observational constraints, we revisit the discussion of the production of high-energy gamma rays from particle interactions in the inner jet of this system. The hadronic contribution could represent a major fraction of the TeV emission detected from this source. The spectral energy distribution resulting from $pp$ interactions is recalculated. Opacity effects introduced by the photon fields of the primary star and the stellar decretion disk are shown to be essential in shaping the high-energy gamma-ray light curve at energies close to 200 GeV. We also present results of Monte Carlo simulations of the electromagnetic cascades developed very close to the periastron passage. We conclude that a hadronic microquasar model for the gamma-ray emission in LS~I +61 303 can reproduce the main features of its observed high-energy $\\gamma$-ray flux. ",
        "introduction": "\\label{intro} LS~I +61 303 is a Be X-ray binary with strong and variable radio emission. It has been suggested as the potential counterpart of the gamma-ray sources 2CG 135+01 and 3EG J0241+6103 (Gregory \\& Talor 1978, Kniffen et al. 1997). Massi et al. (2001) detected relativistic jets in the source, which was then included in the microquasar class. The existence of jets was confirmed by the observations presented by Massi et al. (2004). Very recently, LS~I +61 303 was detected by MAGIC, a very large atmospheric imaging Cherenkov telescope located at La Palma (Albert et al. 2006). This makes LS~I +61 303 the second high-energy microquasar detected by ground-based telescopes. Moreover, MAGIC found clear evidence for variability in the detected signal: the source was stronger at orbital phases 0.5-0.6, whereas the periastron passage occurs at phase 0.23.  At X-rays the source has been detected by XMM-Newton, Beppo-SaX and INTEGRAL (Sidoli et al. 2006, Chernyakova et al. 2006) showing variability on short timescales and a harder spectrum when the source was brighter (Sidoli et al. 2006).  The orbital parameters of LS~I +61 303 have been determined by Casares et al. (2005). The eccentricity of the system is 0.72$\\pm$0.15 and the orbital inclination is $i\\sim30\\pm20$ deg. The nature of the compact object is not well-established. A low-mass black hole cannot be ruled out (Massi 2004). If a mass of $\\sim$ 2.5 M$_{\\odot}$ is assumed for the compact object, then the mass of the Be star should be $\\sim 12$ M$_{\\odot}$, for $i=30$ deg.  From a theoretical point of view, both accreting and non-accreting models have been proposed to explain the high-energy emission of LS~I+61 303. In the accreting scenario (e.g. Bosch-Ramon et al. 2006a) a relativistic jet is launched from the surroundings of the compact object. Strong magnetic fields are supported by an underluminous accretion disk, which is advection-dominated. High-energy emission can be produced in the jet by leptonic (e.g. Bosch-Ramon et al. 2006b, Gupta \\& B\\\"ottcher 2006, Bednarek 2006a) or hadronic (Romero et al. 2003, 2005) interactions. In non-accreting models (Dubus 2006b, Chernyakova et al. 2006) particles are accelerated up to relativistic energies in the interacting region where a pulsar wind collides with the stellar wind. Inverse Compton (IC) cooling of electrons (Maraschi \\& Treves 1981, Dubus 2006b) or $pp$ interactions (Chernyakova et al. 2006) could result in the production of high-energy photons.  A colliding wind model has been proposed as the mechanism producing the high-energy emission in the $\\gamma$-ray binary PSR B1259-63, a system that contains a radio pulsar in an eccentric orbit around a Be star (Aharonian et al. 2005). In the case of LS~I+61 303, the existence of jets extending up to $\\sim 400$ AU from the core seems to support an accretion/jet model (see the discussion in Mirabel 2006). Albert et al. (2006) favor a microquasar model and suggest that the detection of the high-energy signal after the periastron passage would support a leptonic origin for the emission. However, $\\gamma$-ray propagation effects alone could be responsible for such a behavior (Bednarek 2006a).  In this paper we revisit the hadronic model for high-energy emission in LS~I+61 303 proposed by Romero et al. (2005) before the MAGIC detection. Using an improved accretion model and more sophisticated calculations for the opacity effects inside the binary system we show that a hadronic origin for the emission detected by MAGIC can not be ruled out. ",
        "conclusions": "\\label{sec:3} The luminosity obtained at phase 0.5 implies that the constant coupling the jet kinetic power and the accretion luminosity should be $q_{\\rm jet}\\sim 0.1$ in order to explain the $7\\times 10^{33}$ erg s$^{-1}$ inferred from MAGIC detections. Powerful jets can be present in microquasars, as showed by Gallo et al. (2005) for the case of Cygnus X-1.  At energies below the break of the hadronic spectrum ($\\sim$ GeV) the emission of leptonic origin should change the shape of the SED yielding the EGRET  spectrum. The hadronic interactions that produce neutral pions also generate relativistic electrons and positrons by the decay of charged $\\pi$-mesons. We have not computed here the emission of such leptons, for we are concerned only with the very high energy range. Some estimations of the synchrotron luminosity of secondaries were presented in Romero et al. (2005).  Both leptonic and hadronic microquasar models for the gamma-ray emission in LS~I +61 303 can reproduce the main features of the observed gamma-ray emission. The detection of the source at phase 0.5 is a consequence of the opacity effects and is independent of the dominant radiative process in the jet. Even in hadronic models like the one presented here, near the periastron passage gamma-rays will initiate electromagnetic cascades with IC emission, hence the actual situation could be a mixture of radiative processes. More realistic simulations of the photospheric cascades should include the effects of circumstellar disk absorption.  In a pure leptonic model the radiative losses equating the acceleration rate fix the maximum energies of the $\\gamma$-ray photons at $\\sim 1$ TeV (Bosch-Ramon et al. 2006b). Detection of $\\gamma$-ray emission of LS~I +61 303 at higher energies than those observed by MAGIC (e.g. at $E\\geq 10$ TeV) could give support to a hadronic model, and to the presence of intense magnetic fields at the base of the jet.  A clear distinction about the nature of most of the primary gamma-rays can be established through the detection of neutrinos, which are produced only in hadronic models. In the context of our model, we expect around 3-4 muon neutrinos per squared km per year at Earth, so the source might be detectable by IceCube (Christiansen et al. 2006, see also the discussion in Torres \\& Halzen 2006).  As we mentioned in the Introduction, an alternative model for the gamma-ray production in LS~I +61 303 involves an energetic pulsar which could generate a strong wind that would stop the accretion. Particles might be accelerated at the colliding wind region, producing gamma-rays through IC and $pp$ interactions (Dubus 2006b, Chernyakova et al. 2006). However, it is not clear how the morphological and variability properties of the system might be explained in this scenario. What is clear is that the high accretion rates and the observed low X-ray luminosity seem to exclude a scenario based upon an accreting pulsar, since there is no trace of the emission from the heated surface. Future observations with MAGIC will help to detect the source close to the periastron and the spectral evolution along the orbit, providing in this way more constraints to the models."
    },
    "0608/astro-ph0608713_arXiv.txt": {
        "abstract": "Blazars are the most violent steady/recurrent sources of high-energy gamma-ray emission in the known Universe. They are prominent emitters of electromagnetic radiation throughout the entire electromagnetic spectrum. The observable radiation most likely originates in a relativistic jet oriented at a small angle with respect to the line of sight. This review starts out with a general overview of the phenomenology of blazars, including results from a recent multiwavelength observing campaign on 3C279. Subsequently, issues of modeling broadband spectra will be discussed. Spectral information alone is not sufficient to distinguish between competing models and to constrain essential parameters, in particular related to the primary particle acceleration and radiation mechanisms in the jet. Short-term spectral variability information may help to break such model degeneracies, which will require snap-shot spectral information on intraday time scales, which may soon be achievable for many blazars even in the gamma-ray regime with the upcoming GLAST mission and current advances in Atmospheric Cherenkov Telescope technology. In addition to pure leptonic and hadronic models of gamma-ray emission from blazars, leptonic/hadronic hybrid models are reviewed, and the recently developed hadronic synchrotron mirror model for TeV $\\gamma$-ray flares which are not accompanied by simultaneous X-ray flares (``orphan TeV flares'') is revisited. ",
        "introduction": "\\label{introduction}  Blazars (BL~Lac objects and $\\gamma$-ray loud flat spectrum radio quasars [FSRQs]) are the most extreme class of active galaxies known. They have been observed at all wavelengths, from radio through very-high energy (VHE) $\\gamma$-rays. 46 blazars have been identified with high confidence as sources of $> 100$~MeV emission detected by the {\\it EGRET} telescope on board the {\\it Compton Gamma-Ray Observatory} \\cite{hartman99,mattox01}, and about one dozen blazars have now been detected at VHE $\\gamma$-rays ($> 350$~GeV) by ground-based atmospheric \\v Cerenkov telescopes (ACTs). Many of the {\\it EGRET}-detected $\\gamma$-ray blazars appear to emit the bulk of their bolometric luminosity at $\\gamma$-ray energies. Blazars exhibit variability at all wavelengths on various time scales. Radio interferometry often reveals one-sided kpc-scale jets with apparent superluminal motion.  \\begin{figure*} \\centering \\includegraphics[width=0.9\\textwidth]{bbspectra.eps} \\caption{\\it SEDs of 3C~279 \\cite{hartman01a}, BL~Lacertae \\cite{boettcher03}, and Mrk~501 \\cite{petry00}. For each object, two simultaneous broadband spectra at two different epochs are shown. The curves show model fits, using a leptonic jet model.} \\label{bbspectra} \\end{figure*}  \\subsection{Spectral classification of blazars} \\label{classification}  The broadband continuum spectra of blazars are dominated by non-thermal emission and consist of two distinct, broad components: A low-energy component from radio through UV or X-rays, and a high-energy component from X-rays to $\\gamma$-rays. A sequence of blazar sub-classes, from FSRQs to low-frequency peaked BL Lac objects (LBLs) to high-frequency peaked BL Lacs (HBLs) can be defined through the peak frequencies and relative $\\nu F_{\\nu}$ peak fluxes, which also seem to be correlated with the bolometric luminosity \\cite{fossati98}. The sequence FSRQ $\\to$ LBL $\\to$ HBL is characterized by increasing $\\nu F_{\\nu}$ peak frequencies, a decreasing dominance of the $\\gamma$-ray flux over the low-frequency emission, and a decreasing bolometric luminosity (see Fig. 1). LBLs are intermediate between the FSRQs and the HBLs. The peak of their low-frequency component is located at IR or optical wavelengths, their high-frequency component peaks at several GeV, and the $\\gamma$-ray output is of the order of or slightly higher than the level of the low-frequency emission. However, the existence and physical significance of this blazar sequence has recently been questioned and the apparent sequence attributed to selection effects due to the use of flux-limited samples (for a recent review see, e.g., \\cite{padovani06}).   \\begin{figure*} \\centering \\includegraphics[width=0.9\\textwidth]{bllac_variability.eps} \\caption{Spectral variability fitting of BL~Lacertae in 2000 \\cite{br04}: (a) Time-dependent leptonic fits to the Nov. 1 high state; (b) Various hadronic fits, differing mainly in their co-moving magnetic-field and synchrotron photon energy density. (c) Comparison of the simulation results corresponding to fits from panel (a) and (b) to the optical color-magnitude correlation; (d) comparison of the leptonic fit results to the X-ray hardness-intensity diagrams during a short flare observed by {\\it BeppoSAX} on Nov. 1. Based on count rates in the three {\\em BeppoSAX} NFI energy channels LECS [0.5 -- 2 keV], MECS [2 -- 4 keV], and MECS [4 -- 10 keV], and X-ray hardness ratios: HR1 = MECS[2-4]/LECS[0.5-2], HR2 = MECS[4-10]/MECS[2-4]. } \\label{bllac_var} \\end{figure*}  \\subsection{Spectral variability of blazars} \\label{variability}  Fig. \\ref{bbspectra} already illustrates that in particular the high-energy emission from blazars can easily vary by more than an order of magnitude between different {\\it EGRET} observing epochs \\cite{vmontigny95,mukherjee97,mukherjee99}. However, high-energy variability has been observed on much shorter time scales, in some cases less than an hour \\cite{gaidos96}.  BL~Lac objects occasionally exhibit X-ray variability patterns which can be characterized as spectral hysteresis in hardness-intensity diagrams (e.g., \\cite{takahashi96,kataoka00,fossati00,zhang02}). This has been interpreted as the synchrotron signature of gradual injection and/or acceleration of ultrarelativistic electrons in the emitting region, and subsequent radiative cooling (e.g., \\cite{kirk98,gm98,kataoka00,kusunose00,li00,bc02}). While spectral hysteresis has so far only been clearly identified in HBLs, it should also occur in the soft X-ray emission of LBLs if their synchrotron component extends into the soft X-ray regime. However, LBLs are generally fainter at X-ray energies than HBLs, making the extraction of time-resolved spectral information observationally very challenging. Fig. \\ref{bllac_var}d shows the results of a {\\em BeppoSAX} observation of BL~Lacertae in 2000 \\cite{ravasio03,boettcher03}. Rapid flux and spectral variability of blazars is also commonly observed in the optical regime, often characterized by a spectral hardening during flares (see, e.g., Fig. \\ref{bllac_var}c, or \\cite{lainela99,villata02}). ",
        "conclusions": ""
    },
    "0608/astro-ph0608525_arXiv.txt": {
        "abstract": "{We derive the distribution of maser stars in the inner Milky Way based on an analysis of diagrams of longitude versus \\los-velocity (= lV-diagrams) for two samples of maser stars: 771 OH/IR stars and 363 SiO-maser stars. The stars are all close to the plane of the Milky Way and have longitudes from $-45^\\circ$ to $+45^\\circ$.  The two \\lVd s are compared qualitatively and found to be very similar. They also compare well with the \\lVd\\ of interstellar CO, but there are significant differences in detail between the stellar \\lVd s and that of the ISM.  Based on the qualitative discussion we divide the \\lVd s into seven areas. In each area we count the number of stars as observed and compare these numbers with those predicted by an assumed set of orbits in a galactic potential. This potential is axially symmetric but a weak rotating bar has been added. We conclude that the maser stars move on almost circular orbits outside of about 3.5 kpc, but that the orbits become more and more elongated when one goes deep inside our MW. We find a strong effect of the Corotation (=CR) resonance at 3.3 kpc, we see a small but noticeable effect of the Outer Lindblad Resonance at 5 kpc and no effect of the Inner Lindblad Resonance (=ILR) at $r=0.8$ kpc.  We find a set of 6 groups of orbits that together predict counts in agreement with the counts of stars observed. We then calculate the trajectory of each orbit and so find the distribution of the maser stars in the plane of the MWG.  This distribution has two new (but not unexpected) features. The first is a bar-like distribution within 2 kpc from the GC outlined. These orbits explain the high-velocity stars near $l=0^\\circ$ in the forbidden and the permitted quadrants. The second feature are two \"croissant''-like voids in the distribution close to the CR radius (3.3 kpc).  These voids are the consequence of the presence of the co-rotation resonance.  We find excellent agreement with an earlier reconstruction by Sevenster (1999) based on partially the same data but on a completely different analysis. ",
        "introduction": "We cannot see into the inner parts of our Milky Way Galaxy (=MWG) close to the galactic plane because of the extinction of light by dust particles in intervening interstellar clouds. Radiation with a wavelength of a few microns and longer is not affected by the dust and that makes it possible to detect and to study the objects in the inner Galaxy. In the 1950's, under the leadership of J.H.  Oort, the distribution of the interstellar matter (=ISM) throughout the MWG was derived in an elegant and convincing manner using the results of a survey of the 21 cm line of HI. The success was not complete, however: about one third of the galactic plane map was blank: in directions within $30^\\circ$ from the galactic center (=GC) and galactic anticenter the kinematics of the gas could not be used to find the distance to the gas velocities and in addition the gas motions in the GC direction are complex and were not understood at that time. The problem is apparently solved by the introduction of a weak, rotating bar (see e.g.  Bissantz, Englmaier \\& Gerhard, 2003). \\nocite{bissa:03} The 21 cm HI line data have been complemented in the 1970's by surveys of the lowest rotational transition of CO at 115 GHz ($\\lambda = 2.6$ mm), a molecule that acts as substitute for molecular hydrogen. A very conspicuous feature is the (almost complete) absence of ISM inside the so-called ``molecular ring'' with a radius of about 3 kpc, except of the gas within a few hundred parsec of SgrA*.  All these results are well known and have been described in text books- e.g.  Binney \\& Merrifield 1998 (hereafter BM) \\nocite{binne:98}.\\\\  In 1968 Wilson and Barrett anounced the detection of a red giant as the source of a strong emission line of the hydroxyl radical (OH) at 1612 MHz (18 cm).\\nocite{wilso:68} The strength of the line by itself and in comparison to other lines from the same multiplet indicated that the line had been amplified by a maser process in the circumstellar environment of this star. Quickly thereafter Elld{\\'e}r, R{\\\"o}nn{\\\"a}ng \\& Winnberg (1969) \\nocite{ellde:69} announced the discovery of several similar maser sources in a systematic survey of a small part of the Milky Way in Cygnus. These detections suggested that a \"blind\" survey at 1612 MHz of a large fraction of the Milky Way might find enough masers to derive the distribution of these stars in the inner MWG, similar to what had been achieved for the ISM through the 21 cm line. In the following years two more maser-amplified molecular lines were detected: H$_2$O (Knowles et al., 1969 \\nocite{knowl:69}) and SiO (Kaifu, Buhl \\& Snyder, 1975) \\nocite{kaifu:75}.\\\\  The suggestion of a larger, blind 1612 MHz survey led to two searches for maser stars along the northern Milky Way (Bowers, 1978a, 1978b; Baud et al., 1981a). \\nocite{bower:78a, bower:78b, baud:81a} Some 200 stars were detected and from the observed relation between their longitude and latitude coordinate and their \\los\\ velocity\\footnote{We will use the word \\textit{line-of-sight velocity} instead of the more common term \\textit{radial velocity} because the latter will be reserved for stellar motions to and from the GC} it appeared that the stars form a thin, rotating disk with low density (a \"hole\") within about 3 kpc from the GC, much like the hole in the ISM.  An unexpected result was the discovery of a star with a high \\textit{negative} velocity ($-341$ \\kms ) at a \\textit{positive} longitude $(l,b)= (0.3^\\circ,-0.2^\\circ$) (Baud et al., 1975). \\nocite{baud:75} More high-velocity stars were discovered later.  They play an important role in this paper.\\\\  This paper deals with a renewed search for the galactic distribution of maser stars, using sources detected in two recent surveys, one in the OH 1612 MHz line, the other in the 86 GHz line of SiO. The data obtained in these surveys have already been analyzed into quite some detail- as we will discuss below.  Nevertheless we felt that further analysis might still yield some interesting results- hence this paper.\\\\  The paper is structured as follows. (1) We discuss the astrophysical properties of the maser star population. (2) A summary is given of what is already known about the relation between maser stars and the bar of our MWG. (3) The two diagrams of longitude versus line-of-sight velocity (=\\lVd s) of the maser stars are compared qualitatively one with another and then with that of interstellar CO.  (4) After adopting a galactic potential we calculate possible orbits for the maser stars in the galactic plane and their projection on the \\lVd . (5) Using the \\lVd\\ of the stars as a guide we select the orbits in the galactic plane that the maser stars are likely to follow. (6) Using these orbits we derive the spatial distribution of the maser stars in the inner MWG.\\\\ ",
        "conclusions": "We have at our disposal a sample of 1134 maser stars (771 OH, 363 SiO masers), all at low galactic latitudes ($\\mid b\\mid\\leq 3.5^\\circ$) and in the inner MWG ($ |l|\\leq 45^\\circ$).The OH/IR sample contains all OH/IR maser stars with a peak flux density above 0.5 Jy in the area specified and with a \\los\\ velocity within $300$ \\kms\\ from zero velocity in the Local Standard of Rest; see Sevenster et al. (1997a, 1997b, 2001) \\nocite{seven:97a, seven:97b, seven:01}. The SiO maser sample has been obtained by pointing the telescope at infrared objects from the ISOGAL survey (see Messineo et al., 2002)\\nocite{messi:02}. Not all candidate infrared objects have been observed and the longitude distribution may not be representative; all infrared objects were within $1^\\circ$ from the galactic plane. The OH- and the SiO-. survey had a sensitivity limit low enough to find stars beyond the GC.\\\\  We limited ourselves to these two surveys because we had all the data readily available and because we understood the selection criteria and thus the completeness of the surveys. We therefore did not include the large sample of SiO-maser stars detected by Japanese astronomers, e.g. Deguchi et al. (2004b) \\nocite{deguc:04b} who used the IRAS catalogue as the source of candidate stars, and this catalog is quite incomplete near the galactic plane. For the same reason we did not use the earlier OH/IR survey by te Lintel Hekkert et al. (1991). \\nocite{telin:91} One other limitation concerns the immediate surroundings of the GC, roughly the area within $1.5^\\circ$ from Sgr A*. Results from deeper OH surveys in this area are available (Lindqvist, Habing \\& Winnberg, 1992a, Lindqvist et al. 1992b \\nocite{lindq:92a,lindq:92b} and Sjouwerman et al., 1998) \\nocite{sjouw:98} but as we believe that the structure in this area has its own difficulties that are not of importance for the remainder of the MWG we did not discuss these surveys and leave those for a future paper.\\\\  Consider now the conclusions already reached concerning the distribution in position and in velocity of the maser stars as far as these are relevant for our present study.  \\begin{enumerate} \\item Sevenster et al. (2000) \\nocite{seven:00} calculated orbits of stars in a three-dimensional MWG model under the assumption that there are two integrals of motion, the energy, $E$, and the angular momentum around the $z$-axis, $J$. A quite significant conclusion follows from their figure 6: the distribution of the angular momenta changes gradually at energies $E$ corresponding to circular orbits with a radius $r= 2.0$ kpc. At higher energies (i.e. farther outside in the disk) the stars move all on almost-circular orbits and in the clock-wise direction as seen from the galactic north pole. At lower energies and thus deeper inside the MWG the distribution of the angular momenta becomes flat: for a given energy all angular momenta have the same probability and there are as many circular orbits in the anti-clock-wise direction as in the clock-wise direction. \\item If in a galaxy with a rotating bar the density distribution is described as a function of polar coordinates $(r,\\phi)$ the dependence on time must be such that $\\phi=\\phi_0-\\Omega t$. Tremaine and Weinberg (1984) \\nocite{trema:84} have shown that as a consequence of this symmetry the value of $\\Omega$ can be derived from the measurement of \\los\\ velocities in an edge-on galaxy. Using a data base of OH/IR stars almost the same as ours Debattista, Gerhardt \\& Sevenster (2001)\\nocite{debat:02} calculated the pattern speed of a bar in the inner MWG  and obtained the value $\\Omega= -59\\,\\pm 5$ km s$^{-1}$ kpc$^{-1}$. The tangential velocity of the maser stars thus contains a systematic component that is proportional to the distance of the star to the GC. This is a strong argument in favor of the statement that the OH/IR stars are part of the galactic bar. \\item Sevenster (1999b) \\nocite{seven:99b} analyzed a sample of 509 OH/IR stars in the longitude range from $l=+10^\\circ$ down to $l=-45^\\circ$. She concludes: \"Unequivocal morphological evidence is presented for the existence of a central Bar...\". Corotation (CR) is at 3.5 kpc, the inner Lindblad resonance (ILR) is between 1 and 1.5 kpc and the inner ultra-harmonic radius is at 2.5 kpc.  \\item Sevenster, Saha, Valls-Gabaud \\& Fux (1999) \\nocite{seven:99a} have compared several $N$-body simulations of the MWG with the OH/IR stars data obtained with the ATCA. In these simulations a bar is formed spontaneously. Good agreement between predictions and observations is found for a viewing angle of $44^\\circ$, CR at 4.5 kpc.  \\item In the same paper as in item 3 Sevenster points out that the distribution of the OH/IR stars over positive and negative longitudes are slightly different and that this difference has to be attributed to the presence of the bar. We have now a larger data base and our conclusion confirms the one previous; see the next subsection. \\end{enumerate}  We conclude that all these items, except 1, argue that the maser stars are part of the bar of the MWG.  \\subsection{Differences between the distribution at positive and at negative longitudes}  \\begin{figure}\\label{fig:lplusminlmin} \\includegraphics[width=6cm, angle=270]{4480fg1.ps} \\caption{The cumulative distribution of the sign of the longitude, as a function of absolute longitude.The solid curve indicates the distribution when correcting for the lower completeness of the VLA section of the survey (see Sevenster et al 2001). The dashed curve indicates the distribution when no correction is applied, i.e. when effectively every star in the sample counts as 1. A decreasing sum indicates an overdensity at negative longitudes, an increasing sum indicates an overdensity at positive longitudes.  Horizontal sections in the curve indicate a symmetrical distribution of stars with respect to l=0. The dash-dotted curves give the two-sided 95\\% interval assuming the stars are drawn from a binomial distribution. Only OH/IR stars (614 in total) with outflow velocities between 2 and 20 km/s are included, to ensure a homogeneous population of stars that traces the Bar} \\end{figure}  For a long time, the interpretation of asymmetries in the integrated-light distribution observed toward the inner Galaxy was limited by the degeneracy of the deprojection. In external galaxies, edge-on bars do not show up as asymmetries in the morphology, so asymmetries have to be interpreted as $m=1$ distortions (or odd-numbered in general). In the Galaxy, both $m=1$ and $m=2$ non-axisymmetries would result in an asymmetric surface morphology. Most observations before COBE (Weiland et al. 1994, Dwek et al. 1995)\\nocite{weila:94,dwek:95} showed asymmetries that could not be unambiguously attributed to an $m=2$ distortion.\\\\  The COBE observations themselves showed a surface-density distribution with a clear hint of the Bar. The image is, however, still limited by extinction. As explained by Blitz \\& Spergel (1991)\\nocite{blitz:91}, when observing a tracer population that can be sampled out to $d>12$ kpc, so that the bar is fully included, a signature would be seen that would point unambiguously toward the existence of a bar. This signature was first seen in a sample of OH/IR stars (Sevenster 1999). With the extended sample of OH/IR stars, we now repeat the test.\\\\  Figure~1 shows this signature. After arranging all OH/IR stars according to their absolute longitude we show $N=\\Sigma l_i/|l_i|$. If the distribution were axisymmetric $N$ would always remain close to 0. This is clearly not the case.  Out to about $|l|=10^\\circ$ there is an overdensity at negative longitudes. Between $|l|=10^\\circ$ and $30^\\circ$ there is a strong overdensity at positive longitudes. At higher longitudes, the curve approximates a constant value, indicating a certain lopsidedness in the distribution of stars in our MWG. For a more elaborate discussion of this test, see Sevenster (1999) \\nocite{seven:99b}. An important conclusion for this paper is that the sample of OH/IR stars does sample the entire Bar, both the near and the far end. It provides a good sample for kinematical analysis of bar orbits, especially with the accurate line-of-sight velocities.\\\\  The same asymmetric distribution of stars has been found by Benjamin et al. (2005) \\nocite{benja:05} in a much larger sample of red giants detected at mid-infrared wavelengths with the SPITZER satellite; the authors also argue for the existence of the bar.\\\\  \\subsection{Foreground extinction and distance from the Sun}  Stellar masers are formed in the dust-rich envelopes of AGB stars.  In OH masers these envelopes are optically thick to wavelengths as large as 10 \\micron . SiO masers form in thinner envelope and if near-IR photometry is available one can derive \\ak , the sum of interstellar and circumstellar extinction of each maser star. As the maser stars are variables one needs simultaneously measured magnitudes. Existing surveys provide this information: $ I- ,J-$ and $K-$magnitudes by DENIS and $J- , H-$ and $K-$magnitudes by 2MASS.  In all fields the maser star is (one of) the brightest sources and can easily be identified (Messineo et al., 2004, 2005)\\nocite{messi:04,messi:05}. When we separate groups of maser stars based on their value of \\ak\\ we see systematic differences between their \\lVd s; see Fig.~2. These differences we attribute to differences in distance: larger values of \\ak\\ imply a larger distance. The few stars at longitudes larger than $20^\\circ$ and with medium or large extinction may be located behind dense clouds . The systematic change in character between the three diagrams is in agreement with the model that we will develop later in this paper.  \\begin{figure}\\label{fig:AvSiOstars} \\centering \\includegraphics[width=6cm]{4480fg2.ps} \\caption{\\lVd s of different samples of SiO-maser stars; the samples have been selected using an estimate of the amount of the foreground extinction} \\end{figure}"
    },
    "0608/gr-qc0608048_arXiv.txt": {
        "abstract": "We present analytic solutions of Maxwell equations for infinitely long cylindrical conductors with nonvanishing electric charge and currents in the external background spacetime of a line gravitomagnetic monopole. It has been shown that vertical magnetic field arising around cylindrical conducting shell carrying azimuthal current will be modified by the gravitational field of NUT source. We obtain that the purely general relativistic magnetic field which has no any Newtonian analog will be produced around charged gravitomagnetic monopole. ",
        "introduction": "General relativity predicts that so called gravitomagnetic monopole or exotic NUT charge~\\cite{nut63} creates a gravitomagnetic field as in electrodynamics Dirac hypothetic monopole generates magnetic field. In recent years, there exists an interest in problems related to the general relativistic effects in the gravitational field of NUT sources. For example, some interesting results devoted to the physical aspects of the Kerr-NUT spaces have been discussed in papers~\\cite{bini03}$^{,}$ \\cite{rh03}$^{,}$ \\cite{nz02}. In particular in~\\cite{bini03} Klein-Gordon and Dirac equations in slowly rotating NUT spacetime have been investigated and in~\\cite{zonoz03} the solutions of the Maxwell equations for the electromagnetic waves in Kerr-NUT space have been found. In recent papers~\\cite{rnz03},\\cite{rh03} the theoretical effect of magnetic mass in NUT space on the microlensing light curve and the possibility of magnetic mass detection using the gravitational microlensing technique have been studied.  In our previous research~\\cite{99pla,00adp,03fp,ram01} the effect of gravitomagnetic field of spinning gravitating body on the electric current and electromagnetic field has been investigated on the theoretical level. In fact, the influence of the angular momentum of the rotating gravitational source may appear as a galvanogravitomagnetic effect in the current carrying conductors~\\cite{99pla} and as a general-relativistic effect of charge redistribution inside conductors in an applied magnetic field ~\\cite{00adp}. It has been shown in~\\cite{03fp} that the gravitomagnetic interaction with electric field can lead to the applicable general-relativistic effects. The general relativistic effect of ''dragging of inertial frames'' on electromagnetic fields produced by rotating magnetic dipole has been investigated in our papers~\\cite{ram01}.  Here we extend our research to the electromagnetic fields of conducting shell embedded in the NUT space around cylindric gravitomagnetic monopole and find analytical solutions of the Maxwell equations in this space-time.  The paper is organized as follows. In the section~\\ref{meq} we write the Maxwell equations in the exterior metric of cylindrical NUT source. In the next section we solve Maxwell equations in NUT space around i) conducting cylinder carrying electric current and ii) electric charge put on line gravitomagnetic monopole. We summarize our conclusions in the section~\\ref{concl}.  Throughout, we use a space-like signature $(-,+,+,+)$ and a system of units in which $G = 1 = c$ (However, for those expressions with a physical application we have written the speed of light explicitly.). Greek indices are taken to run from 0 to 3 and Latin indices from 1 to 3; covariant derivatives are denoted with a semi-colon and partial derivatives with a comma. ",
        "conclusions": "\\label{concl}  We have presented analytic general relativistic expressions for the electromagnetic fields external to the cylindrical conductor either carrying electric current or having electric charge in the space-time exterior to line gravitomagnetic monopole which is considered isolated and in vacuum.  In the case of the carrying current conductor we have shown that the magnetic field arising around it will be modified by the gravitational field of NUT source. In the case of the charged conductor an additional purely general relativistic magnetic field will be induced by the magnetic mass.  In our future research  we will present analytical solutions for the electromagnetic fields of slowly rotating magnetized spherical NUT stars arising from the interplay between gravitomagnetic charge and electric field."
    },
    "0608/astro-ph0608463_arXiv.txt": {
        "abstract": "We report an analysis of the archival Rossi X-ray Timing Explorer (RXTE) data from the December 2004 hyperflare from SGR 1806-20.  In addition to the $\\approx 90$ Hz QPO first discovered by Israel et al., we report the detection of higher frequency oscillations at $\\approx 150$, 625, and 1,840 Hz.  In addition, we also find indications of oscillations at $\\approx 720$, and 2,384 Hz, but with lower significances. The 150 Hz QPO has a width (FWHM) of about 17 Hz, an average amplitude (rms) of $6.8\\%$, and is detected in average power spectra centered on the rotational phase of the strongest peak in the pulse profile. This is approximately half a cycle from the phase at which the 90 Hz QPO is detected. The 625 Hz oscillation was detected in an average power spectrum from nine successive cycles beginning approximately 180 s after the initial hard spike. It has a width (FWHM) of $\\approx 2$ Hz and an average amplitude (rms) during this interval of $8.5\\%$. We find a strong detection of the 625 Hz oscillation in a pair of successive rotation cycles begining about 230 s after the start of the flare.  In these cycles we also detect the 1,840 Hz QPO. When the 625 Hz QPO is detected we also confirm the simultaneous presence of 30 and 92 Hz QPOs.  The centroid frequency of the 625 Hz QPO detected with RXTE is within 1 Hz of the $\\approx 626$ Hz oscillation recently found in RHESSI data from this hyperflare by Watts \\& Strohmayer.  We argue that these new findings provide further evidence for a connection of these oscillations with global oscillation modes of neutron stars, in particular, the high frequency signals may represent toroidal modes with at least one radial node in the crust. We discuss their implications in the context of this model, in particular for the depth of neutron star crusts. ",
        "introduction": "Recent studies of high time resolution data from two magnetar hyperflares recorded by the Proportional Counter Array (PCA) onboard the Rossi X-ray Timing Explorer (RXTE) have resulted in the discovery of a new phenomenon associated with these events. Both the December, 2004 event from SGR 1806-20, and the August, 1998 hyperflare from SGR 1900+14 produced fast, rotation-phase-dependent X-ray oscillations in the 20 - 150 Hz range.  First, Israel et al. (2005) reported the discovery of $\\approx 18$, 30, and 90 Hz quasi-periodic oscillations (QPO) in the December, 2004 event, and suggested that the 30 and 90 Hz QPOs could be linked with seismic vibrations of the neutron star crust. Strohmayer \\& Watts (2005) then reported the discovery of a sequence of QPOs in the SGR 1900+14 hyperflare.  They found a set of frequencies; 28, 53.5, 84, and 155 Hz, which could plausibly be associated with a sequence of low $l$ toroidal modes (denoted $_l t_0$) of the elastic neutron star crust (see, for example, Hansen \\& Cioffi 1980; McDermott, van Horn \\& Hansen 1988; Duncan 1998; Piro 2005).  In both hyperflares the oscillations are episodic, that is, their amplitudes vary considerably with time and rotational phase. Most recently, Watts \\& Strohmayer (2006) examined Ramaty High Energy Solar Spectroscopic Imager (RHESSI) data for the SGR 1806-20 event, confirmed the presence of the 18 Hz and 90 Hz QPOs, and found evidence for additional oscillations at 26 Hz and 626 Hz.  The similar phenomenology of the oscillations in the two sources, as well as the closeness of some of the measured frequencies argues rather convincingly that we are seeing the same physical process in each case.  The connection with torsional modes of the crust seems plausible for several reasons; 1) The observed frequencies are consistent with theoretical expectations for such modes, and can be more or less self-consistently associated with a sequence of modes with varying spherical harmonic index, $l$. 2) The magnetic instability which powers the hyperflares is likely associated with large scale fracturing of the neutron star crust \\citep{flo77, tho95, dun98, td01, sch05}, and will almost certainly generate seismic motions within the star. 3) The strong phase dependence argues for a mechanism associated with particular sites on the stellar surface, such as a fracture zone or magnetic field bundle. 4) Mechanical motions provide a natural explanation for the relatively high coherence of the oscillations.  5) Horizontal motions of the crust could modulate the beaming pattern associated with the strong magnetic field, providing a mechanism to modulate the X-ray flux. Moreover, beaming can act as an ``amplifier,'' producing potentially large X-ray modulations from modest horizontal displacements.  Although the present evidence for torsional modes is very suggestive, it is not yet definitive. Levin (2006), for example, argues that toroidal modes may damp too quickly to account for the detection of oscillations some minutes after the onset of the flare.  In addition some of the detected frequencies (18 Hz, 26 Hz) do not fit easily into current torsional mode models, without invoking magnetic splitting or other complications.  In our previous study of the RXTE data from the SGR 1900+14 hyperflare we computed average power spectra for different rotational phases and time intervals during the event.  This method should be more sensitive to oscillations which are localized in rotational phase, as borne out in the SGR 1900+14 study.  In this paper we report the results of a similar, phase averaged timing study on the now public archival RXTE data from the SGR 1806-20 hyperflare.  Using this method we find strong evidence for additional, higher frequency oscillations during the flare. ",
        "conclusions": "Our study of the RXTE data from the SGR 1806-20 hyperflare indicates that a complex pulsation phenomenology is associated with magnetar hyperflares.  The discovery of new kHz-range frequencies consistent with theoretical predictions for $n>0$ torsional modes provides strong evidence that we may be seeing vibration modes of the neutron star crust excited by these catastrophic events. If this is true, then it opens up the exciting prospect of probing the interiors of neutron stars in a manner analogous to helioseismology.  Additional excitement is warranted when we consider that all the current datasets used to explore these oscillations have been purely serendipitous.  That is, they have not been optimized in any way for studying these signals. This suggests that a wealth of additional information would likely be found from instruments better optimized to capture with high time resolution the flood of X-rays produced by these events.  Based on the findings presented here we are also strongly convinced that more theoretical work is definitely needed to make more accurate mode identifications, to better understand the excitation and damping mechanisms of modes and how they can couple to the X-ray emission, and to make more precise inferences on neutron star structure."
    },
    "0608/astro-ph0608149_arXiv.txt": {
        "abstract": "Observations of the gamma-ray sky reveal the most powerful sources and the most violent events in the Universe. While at lower wavebands the observed emission is generally dominated by thermal processes, the gamma-ray sky provides us with a view on the non-thermal Universe. Here particles are accelerated to extreme relativistic energies by mechanisms which are still poorly understood, and nuclear reactions are synthesizing the basic constituents of our world. Cosmic accelerators and cosmic explosions are the major science themes that are addressed in the gamma-ray regime.  With the INTEGRAL observatory, ESA has provided a unique tool to the astronomical community revealing hundreds of sources, new classes of objects, extraordinary views of antimatter annihilation in our Galaxy, and fingerprints of recent nucleosynthesis processes. While INTEGRAL provides the global overview over the soft gamma-ray sky, there is a growing need to perform deeper, more focused investigations of gamma-ray sources. In soft X-rays a comparable step was taken going from the Einstein and the EXOSAT satellites to the Chandra and XMM/Newton observatories. Technological advances in the past years in the domain of gamma-ray focusing using Laue diffraction and multilayer-coated mirror techniques have paved the way towards a gamma-ray mission, providing major improvements compared to past missions regarding sensitivity and angular resolution. Such a future Gamma-Ray Imager will allow to study particle acceleration processes and explosion physics in unprecedented detail, providing essential clues on the innermost nature of the most violent and most energetic processes in the Universe. ",
        "introduction": "As introductory remark, it is worth emphasizing some unique features of gamma-ray astronomy: the specific character of the emission processes, the diversity of the emission sites, and the penetrating nature of the emission.  First, the emission process that leads to gamma-rays is in general very specific, and as such, is rarely observable in other wavebands. At gamma-ray energies, cosmic acceleration processes are dominant, while in the other wavebands thermal processes are generally at the origin of the emission. For example, electrons accelerated to relativistic energies radiate gamma-ray photons of all energies through electromagnetic interactions with nuclei, photons, or intense magnetic fields. Accelerated protons generate secondary particles through nuclear interactions, which may decay by emission of high-energy gamma-ray photons. At gamma-ray energies, nuclear deexcitations lead to a manifold of line features, while in the other wavebands, it is the bound electrons that lead to atomic or molecular transition lines. For example, the radioactive decay of tracer isotopes allows the study of nucleosynthesis processes that occur in the deep inner layers of stars. The interaction of high-energy nuclei with the gas of the interstellar medium produces a wealth of excitation lines that probe the composition and energy spectrum of the interacting particles. Finally, annihilation between electrons and positrons result in a unique signature at 511~keV that allows the study of antimatter in the Universe.  Second, the sites of gamma-ray emission in the Universe are very diverse, and reach from the nearby Sun up to the distant Gamma-Ray Bursts and the cosmic background radiation. Cosmic acceleration takes place on all scales: locally in solar flares, within our Galaxy (e.g.~in compact binaries, pulsars and supernova remnants), and also in distant objects (such as active galactic nuclei or gamma-ray bursts). Cosmic explosions are another site of prominent gamma-ray emission. They produce a wealth of radioactive isotopes, are potential sources of antimatter, and accelerate particles to relativistic energies. Novae, supernovae and hypernovae are thus prime targets of gamma-ray astronomy.  Third, gamma-rays are highly penetrating, allowing the study of otherwise obscured regions. Examples are regions of the galactic disk hidden by dense interstellar clouds, or the deeper, inner, zones of some celestial bodies, where the most fundamental emission processes are at work. New classes of sources become visible in the gamma-ray domain, that are invisible otherwise.  In summary, gamma-ray astronomy provides a unique view of our Universe. It unveils specific emission processes, a large diversity of emission sites, and probes deeply into the otherwise obscured high-energy engines of our Universe. The gamma-ray Universe is the Universe of particle acceleration and nuclear physics, of cosmic explosions and non-thermal phenomena. Exploring the gamma-ray sky means exploring this unique face of our world, the face of the evolving violent Universe. ",
        "conclusions": "The gamma-ray band presents a unique astronomical window that allows the study of the most energetic and most violent phenomena in our Universe. With ESA's INTEGRAL observatory, an unprecedented global survey of the soft gamma-ray sky is currently performed, revealing hundreds of sources of different kinds, new classes of objects, extraordinary views of antimatter annihilation in our Galaxy, and fingerprints of recent nucleosynthesis processes. While INTEGRAL provides the longly awaited global overview over the soft gamma-ray sky, there is a growing need to perform deeper, more focused investigations of gamma-ray sources, comparable to the step that has been taken in X-rays by going from the EINSTEIN satellite to the more focused XMM-Newton observatory. Technological advances in the past years in the domain of gamma-ray focusing using Laue diffraction techniques have paved the way towards a future gamma-ray mission, that will outreach past missions by large factors in sensitivity and angular resolution. Such a future {\\em Gamma-Ray Imager} will allow to study particle acceleration processes and explosion physics in unprecedented depth, providing essential clues on the intimate nature of the most violent and most energetic processes in the Universe."
    },
    "0608/hep-th0608168_arXiv.txt": {
        "abstract": "We develop the general program of the unification of matter-dominated era with acceleration epoch for scalar-tensor theory or dark fluid. The general reconstruction of single scalar-tensor theory is fulfilled. The explicit form of scalar potential for which the theory admits matter-dominated era, transition to acceleration and (asymptotically deSitter) acceleration epoch consistent with WMAP data is found. The interrelation of the epochs of deceleration-acceleration transition and matter dominance-dark energy transition for dark fluids with general EOS is investigated. We give several examples of such models with explicit EOS (using redshift parametrization) where matter-dark energy domination transition may precede the deceleration-acceleration transition. As some by-product, the  reconstruction scheme is applied to scalar-tensor theory to define the scalar potentials which may produce the dark matter effect. The obtained modification of Newton potential may explain the rotation curves of galaxies. ",
        "introduction": "The dark energy era origin remains to be one of the challenges of modern cosmology. It has been proposed number of various models aiming at the description of dark energy universe (for recent review, see \\cite{review}). However, even being consistent with recent WMAP data, such models turn out to be problematic at the early and intermediate universe. In other words, we are lacking the unified theory which describes at once the sequence of the well-known cosmological epochs: inflation, radiation/matter dominated epoch and the current universe speed-up. In such a situation, the attempt to unify several known cosmological phases may be considered as some approximation for such a complete theory. From another side, it may rule out number of dark energy models which being consistent with recent WMAP data cannot describe correctly the earlier cosmological phases.  In the present paper we develop the general program of the search of dark energy models (scalar-tensor theory or ideal fluid)  where the sequence of matter dominated era and acceleration era is realized. The paper is organized as follows. Section \\ref{sec2} is devoted to the reconstruction program of scalar-tensor theory admitting matter dominated era before the current acceleration. In the class of theories with single scalar it is explicitly demonstrated for which scalar potentials the matter dominated stage precedes the acceleration era. Clearly, the method may be generalized for multi-scalar case (compare with \\cite{rec}). In section \\ref{sec3} we study the interrelation of the epochs of deceleration-acceleration transition and the matter-dark energy domination transition in the class of dark energy fluids with a general EOS. In the subsection \\ref{sec3.1} we develop the general setting for this class of models, present some general results and the interesting special cases. In subsection \\ref{sec3.2} we give examples of dark energy fluids, defined by the concrete EOS, in which in certain parametric regimes the epoch of the matter-dark energy domination transition may precede the epoch of the deceleration-acceleration transition. In subsection \\ref{sec3.3} we analyze a redshift parametrization of the dark energy EOS and elaborate the conditions on the parameters to observe the matter-dark energy domination transition before the deceleration-acceleration transition. In the section \\ref{sec4} we show that the same reconstruction method of section II may be applied to explain the qualitatively different phenomena of dark matter from scalar-tensor theory with specific potential. In particulary, it is shown which types of scalar potential lead to requested modification of Newton potential at large distances so that rotation curves of galaxies may be explained. Some summary and outlook are given in the Discussion. ",
        "conclusions": "\\label{sec5}  In summary, we demonstrated that matter dominated era may be combined with acceleration era within some class of scalar-tensor theory. The explicit program of scalar-tensor theory reconstruction is presented. Clearly, this makes non-realistic single scalar-tensor theory with other types of potentials as dark energies due to appearence of cosmological bounds. Nevertheless, the situation may be improved for several scalars (several dark energy types) where the realization of sequence of matter dominated and acceleration era should be investigated from the very beginning. This will be studied in other place.  The same problem is discussed for dark energy fluid in terms of redshift parametrization. It is shown that for some examples of EOS the matter domination-dark energy transition may occur prior to deceleration-acceleration transition. Hence, intermediate and late-time stages of the universe may be unified also in frames of dark energy fluid. As a by-product the reconstruction method which was used to get the form of scalar potential from cosmological bounds is applied for the same scalar-tensor theory to get the dark matter effect from it. The scalar-tensor theory with requested modified Newton law at large scales is constructed. Such a theory may explain the rotation curves of the galaxies and the formation of galaxy clusters. It is interesting that the obtained scalar potential seems to be qualitatively similar to the one obtained from the cosmological bounds.  It seems that we are still at the beginning of the road to complete theory which unifies  all the known epochs being compatible with astrophysical/cosmological data and Solar System tests. Nevertheless, the fact that such unification is possible at least for intermediate-time and late-time universe is quite promising."
    },
    "0608/astro-ph0608305_arXiv.txt": {
        "abstract": "In the 1970s, the Narrabri intensity interferometer was used to measure 32 stellar diameters; some as small as 0.4 milli-arc-seconds (mas). The interferometer consisted of a pair of 6.5\\,m telescopes ($\\rm 30\\,m^2$) with relatively crude optics, similar to those currently in use as Atmospheric Cherenkov Telescopes (ACT). We explore the possibility of implementing a modern intensity interferometer on an ACT array. Developments in fast digital signal processing technology now make such a system relatively easy to implement, and provide improved sensitivity. Allowing measurements at short wavelength ($\\rm <400\\,nm$), with long baselines ($\\rm > 100\\,m$), which are still challenging for Michelson interferometers, present ACT arrays could be used to probe angular structures as small as $\\rm \\sim 0.2\\,mas$, and smaller with large array projects already being discussed. This would provide measurements of stellar diameters, binary systems, circumstellar environments and, possibly, stellar surface features. ACT arrays could be used as intensity interferometers during bright moon periods, providing valuable scientific output for little expense and no impact on the $\\gamma$-ray observing schedule. ",
        "introduction": "The first successful measurement of the angular diameter of a star was made using the Michelson optical interferometer on Mount Wilson, over a baseline of 20 feet, \\cite{michelson, michelsonpease}. Difficulties encountered in the subsequent operation of a 50 foot baseline interferometer brought the development of stellar interferometry to a halt for thirty years. No further progress was made until the sixties, with the construction and operation of the Narrabri intensity interferometer.  Intensity interferometry was originally developed in the early fifties in the field of radio astronomy by R.~Hanbury~Brown, who soon extended the idea to the optical domain. After several laboratory tests and actual stellar intensity interferometer prototypes, the technique culminated in the 60s with the interferometer constructed under the direction of Hanbury Brown and Twiss in Narrabri. The Narrabri intensity interferometer was successfully used to study 32 stars \\cite{brown1974}, all brighter than B=+2.5, providing angular diameter measurements as small as $0.41\\pm0.03$ milli-arc-seconds (mas). Despite these successes, and plans for a larger instrument, the technique was then abandoned to the benefit of Michelson interferometry, which had become attractive again because of technological developments in photo-detection and optics. Since then, to our knowledge, intensity interferometry has not been used for astronomical measurements.  In this paper, we revisit intensity interferometry and the scientific potential it could offer if implemented on large Atmospheric Cherenkov Telescopes using modern signal processing technology. In section 2, we briefly review the principles of intensity interferometry. In section 3, we describe how an intensity interferometer could be implemented with an ACT array and in the two following sections, we comment on some specific points regarding the telescope's optics and the electronics. In the last section, we conclude and comment on the possible future developments and scientific potentials of intensity interferometry. ",
        "conclusions": "We have previously presented the possibility of implementing an intensity interferometer using atmospheric Cherenkov telescope arrays \\cite{lebohec05a,lebohec05b}. Here, we have shown that existing atmospheric Cherenkov telescopes could be used as the receivers of intensity interferometers to study stars of magnitude as large as 5 without taking advantage of the technological developments since the time of the Narrabri interferometer. Higher sensitivities are achievable with the faster and more stable electronics available today. These improvements, however, will not affect the sensitivity limit arising from the quality of the receiver optics. We have seen that existing ACTs, with a point spread function of $\\sim0.05^o$ are restricted to objects of magnitudes smaller than $\\rm \\sim 9.6$ with a $\\rm \\sim 5$ hour integration. The optical requirements for interferometric studies could be taken into account in the design of future arrays of larger Cherenkov telescopes to make them more effective.  We have considered the constraints and technical difficulties associated with such a project and proposed some possible technical solutions without finding major difficulties. We conclude that available telescope arrays can, in principle, be used for the interferometric study of stars at wavelengths shorter than $\\rm 400\\,nm$ with baselines close to $\\rm 200\\,m$. This would give access to angles as small as $\\rm \\sim0.1\\,mas$. Figure~\\ref{ang_diam_veritas} shows that, for example, the HESS or VERITAS baselines allow for measurements of the correlation curve for stars with angular sizes of $\\rm \\sim0.5\\,mas$, corresponding to a star the same radius as the sun at a distance of $\\rm \\sim20\\,pc$. Future Cherenkov telescope arrays being discussed extend over $\\rm \\sim 1\\,km$ and could be used to measure angles ten times smaller. One should then question the importance and role intensity interferometry could play in the coming years.  We note that the measurements provided by the Narrabri intensity interferometer of 32 stars, in the spectral range O5f to F8 \\cite{brown1974}, are still relevant and competitive today \\cite{kervalla} although the instrument ceased taking measurements as long ago as 1972. The most useful observations which could be made with a modern-day intensity interferometer are guided mainly by the unique strengths of the technique.  The relative ease of long-baseline measurements allows for very high angular resolution.  A further consideration is the ability to make measurements at shorter wavelengths than those usually probed by Michelson devices. Finally, long exposures using ground-based $\\gamma$-ray telescopes under moonlight are effectively free of competition for observing time, allowing the possibility of very long-term monitoring of variable sources.  The scientific potential of optical interferometry in general is outlined in e.g. \\cite{monnier}. Figure \\ref{stars} illustrates the capabilities of a ground-based $\\gamma$-ray telescope array at measuring stellar angular diameters. Knowledge of the angular diameter of stars is important for the understanding of their fundamental properties. Angular measurements are affected by the limb darkening in a way that depends on the wavelength \\cite{mozurkewich}. Intensity interferometers, with their capability of working over a range of short visible wavelengths, could contribute to the measurements of the stellar diameter wavelength dependence. The long base-lines offer the possibility of measuring main sequence late-type stars for which data is still missing\\cite{kervalla}. Measurements of Cepheid stars can be used to obtain improved calibration of the period-luminosity relation \\cite{kervalla2006}. Arrays of more than two telescopes could also be used to measure the oblateness of rapid rotators, providing further parameter constraints on stellar evolution models\\cite{vanbelle2006}. Future arrays with many base lines simultaneously available could even measure finer details on stellar surfaces as in \\cite{ young2000}. The observation of spectroscopic binaries is also of interest; resolving the components of these systems would determine all of the system parameters, including the respective members' masses \\cite{davis2005}. It is also simple, with an intensity interferometer, to select emission lines so as to study the circumstellar environment. The NPOI was used recently to study $\\rm H-\\alpha$ emitting regions in Be stars $\\gamma$ Cassiopeiae and $\\phi$ Persei with base lines of $\\rm 40\\,m$ and less. Using an ACT intensity interferometer with an $H-\\beta$ emission line filter could possibly reveal details ten times smaller\\cite{tycner2006}. Objects such as $\\zeta$-Tauri stars \\cite{tycner2004}, Wolf-Rayet stars, Luminous-Blue-Variables \\cite{chesneau2003} or even novae \\cite{quirrenbach} could also be monitored at short wavelength with unprecedented interferometric baselines.  In summary, atmospheric Cherenkov telescope arrays provide an inexpensive and practical way by which to implement a modern intensity interferometer. The technical difficulties are few, and the successful operation of an intensity interferometer system on currently operating ACT arrays would significantly increase the scientific output of these instruments. The experience gained from operating such a system would also allow for direct input, relevant to interferometry measurements, into the design of next generation systems.  \\begin{figure}[stars] \\plotone{f6.eps} \\caption{Visual magnitude - angular diameter relationship for the main sequence, the giant and the super-giant branches for distances of 3\\,pc, 30\\,pc and 300\\,pc. The magnitude limitation resulting from a $0.05^o$ optical point-spread function is indicated, as well as the VERITAS baselines as an example. Future atmospheric Cherenkov telescope arrays might provide baselines of 1\\,km.} \\label{stars} \\end{figure}"
    },
    "0608/astro-ph0608419_arXiv.txt": {
        "abstract": "We present results from new {\\it Chandra} and archival {\\it HST} and VLA imaging observations of the circumnuclear extended emission in the nearby Type 2 Seyfert galaxy NGC 2110. We find resolved soft-band X-ray emission $\\sim4''$ ($\\sim160$ pc) north of the nucleus, which is spatially coincident with \\ot emission, but lies just beyond the northern edge of the radio jet in the source. We find that shock-heating of multi-phase gas clouds can successfully account for this extended emission, although we cannot rule out alternative models, such as the scattering of nuclear radiation by ionized material, or pure photoionization from the nucleus. In addition, we detect kpc-scale ($\\sim30''$) extended soft-band X-ray emission south of the nucleus. Finally, we compare our results for NGC~2110 with the prototypical Type 2 Seyfert galaxy NGC~1068, and suggest that different physical processes could produce extended circumnuclear X-ray emission in Seyfert galaxies. ",
        "introduction": "X-ray observations of the circumnuclear environments of Seyfert galaxies with small-scale radio jets are important diagnostics for understanding the interplay between the ionizing radiation from the nucleus, radio ejecta, and galaxy ISM. In some of the brightest sources, e.g., NGC~1068 (e.g.,~\\citealt{you01}), Mrk~3 (\\citealt{sak00}), and the Circinus galaxy (\\citealt{sam01}), high-resolution \\Ch and \\XMM grating spectroscopy reveals their X-ray spectra to be dominated by He-like emission lines most likely to be produced in a plasma photoionized by nuclear radiation. However, alternative models, such as shock-heating induced by the interaction of the radio jet and ambient medium may be applicable to other Seyfert galaxies. High resolution X-ray imaging observations, combined with multiwavelength data, are important probes of circumnuclear environments, and are complementary to X-ray grating data.  NGC 2110, the subject of this paper, is a nearby ($z=0.007579$, $D_{\\rm L}=32.7$ Mpc) Seyfert type 2 galaxy (\\citealt{shu80}). Historically, it is classified as a Narrow Emission Line Galaxy (NELG), which is the subclass of Seyfert galaxies with strong 2--10 keV X-ray emission and narrow ($<600$km s$^{-1}$) optical emission lines. The source exhibits complex morphology on a range of angular scales, and shows a rich variety of structure across the electromagnetic spectrum. VLA observations (\\citealt{ulv89}) show that it hosts an `S'-shaped jet-counterjet system that extends $\\sim4''$ ($\\sim$600 pc) in the N-S direction, while HI spectral-line mapping (\\citealt{gal99}) indicates the presence of a small-scale disk that may partially obscure the southern radio jet. Optically, NGC 2110 is classified as a S0 galaxy, although broad-band {\\it HST} observations (\\citealt{mul94,gon02,fer04}) have revealed a series of complex, spiral-like dust lanes with a distribution consistent with a tightly wrapped single-armed spiral.  Emission-line \\ot$\\lambda$5007 and H$\\alpha$ + [N~{\\sc ii}] {\\it HST} observations of the inner regions of NGC 2110 (\\citealt{mul94}) have revealed a variety of structures. The dominant line emission is concentrated into a strongly curved ``jet-like'' feature (as characterized by \\citealt{mul94}) that extends $\\sim1.4''$ ($\\sim230$ pc) North of the nucleus and becomes increasingly broadened with distance. The \\ot emission is noticeably more asymmetric about the nucleus than the H$\\alpha$ emission, and \\cite{mul94} suggest that this is at least in part due to a north-south gradient in excitation. Detailed spectroscopy of the ``jet-like'' feature with the {\\it HST} Faint Object Spectrograph (FOS) (\\citealt{fer99}) suggests that photoionizing-shock models (e.g.,~\\citealt{ds96}) provide a good explanation of the origin of this emission, although photoionization by a central source cannot be ruled out. Recently, OASIS integral field spectroscopy of the same region (\\citealt{fer04}) confirmed that the majority of the narrow \\ot emission is not associated with the nucleus, and may rather be associated with the ``jet-like'' feature. This is in contrast to the H$\\alpha$ emission, whose peak intensity is coincident with that of the radio nucleus. On slightly larger scales, \\cite{mul94} also report the detection of an `S'-shaped distribution of ionized gas that extends $\\sim4''$ in either direction of the nucleus. Although this feature, together with the ``jet-like'' feature, extend along a similar position angle to the radio emission, neither appears to be spatially coincident with the radio morphology. Instead, the strongest optical line emission occurs where the radio emission is weak, and the optical ``jet-like'' feature is more curved than the radio jet. This suggests that, although possibly influenced by the radio-emitting feature, the optical line emission is not directly associated with it.  At X-ray energies, {\\it ROSAT} and {\\it BBXRT} observations of NGC 2110 (\\citealt{wea95}) showed the presence of spatially resolved X-ray emission that extends $\\sim4''$ North of the unresolved nucleus, possibly coincident with the optical emission-line `S'-shaped feature. They argued that at least a fraction of the soft excess observed in the X-ray spectrum of NGC 2110 below 1~keV (e.g.,~\\citealt{mus82}) originates in this extended region. However, the physical origin of the extended emission was unclear, and several models have been proposed to explain it. In one model, the impact of the radio jet into the ISM may induce shock-heating, producing the observed soft X-ray emission, together with the optical emission lines. \\cite{fer99} provide an in-depth discussion of photoionizing shock models, and demonstrated that if {\\it in situ} shock heating were to power the X-ray emission, shock velocities in excess of 500 km s$^{-1}$ would be required (c.f. the measured \\ot linewidth of 200 km s$^{-1}$ in this region --- e.g.,~\\citealt{wil85}). However, the shock-heating could occur in a region much closer to the nucleus, where the measured \\ot velocities are higher, and subsequently expand outwards. Another interpretation, discussed by \\cite{wea95}, is that the extended X-ray emission is electron-scattered nuclear radiation. Alternatively, the extended emission may result from pure photoionization of ambient gas by the nucleus. A final interpretation, given the common detection with \\Ch of resolved kpc-scale X-ray jets in low-power radio galaxies (e.g.,~\\citealt{wor01}), is that the resolved X-ray emission is from the VLA-detected radio ejecta. \\cite{don04} showed a series of zeroth-order \\Ch HETGS images of the nucleus of NGC~2110, which hinted at the presence of extended soft-band X-ray emission north of the nucleus, and we use the superb spatial resolution and good low-energy sensitivity of the \\Ch ACIS instrument to provide further constraints on the nature of this emission.  In this paper, we present a multiwavelength imaging study of the circumnuclear extended emission in NGC~2110. We make a detailed comparison between the \\Ch X-ray, and archival radio and optical observations of the extended emission north of the nucleus, and discuss models that might explain its origin. In addition, we present the first evidence for large-scale soft-band X-ray emission south of the nucleus. Finally, we compare our results for NGC~2110 with the prototypical Type 2 Seyfert galaxy NGC~1068, and argue that different physical processes could produce the extended circumnuclear X-ray emission in Seyfert galaxies. A detailed analysis of the \\Ch HETGS spectra of NGC~2110 will be presented in a future paper (Lee et al., in prep.), and the energetics arguments presented in this paper are in part a prelude to our analysis of the high energy-resolution \\Ch gratings data.  This paper is organized as follows. Section 2 contains a description of the data and a summary of their analysis. In Section 3, we present the \\Ch imaging observation and compare it to observations with {\\it HST} and the VLA. In Section 4, we discuss several models that may account for the small-scale northern extended emission, while in Section 5 we discuss the large-scale southern extended X-ray emission. In Section 6, we compare NGC~2110 with the prototypical Seyfert 2 galaxy NGC~1068. We end with our conclusions in Section 7. All results presented in this paper use a cosmology in which $\\Omega_{\\rm m, 0}$ = 0.3, $\\Omega_{\\rm \\Lambda, 0}$ = 0.7, and H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$. At the redshift of NGC 2110\\footnote{Taken from the SIMBAD database} ($z$ = 0.007579), 1 arcsec corresponds to 160 pc. This value is slightly different ($-131$ km s$^{-1}$) than the velocity (with respect to the CMB) derived from stellar Mg {\\sc i}b absorption lines (\\citealt{nel95}). However, none of our conclusions would be changed if we adopted this alternate value, and we note that previous publications on NGC~2110 have used a variety of redshift estimates. All spectral fits include absorption through our Galaxy using $N_{\\rm H, Gal}$ = 1.76$\\times$10\\(^{21}\\) cm$^{-2}$ (\\citealt{dic90}). Errors quoted in this paper are 90 per cent confidence for one parameter of interest (i.e., $\\chi^2_{\\rm min}$ + 2.7), unless otherwise stated. When distinguishing between different model fits to the data, we consider $F$-statistic results, and we adopt a threshold of 95\\%  for a significant improvement in the fitting statistic. ",
        "conclusions": "We have presented results from a \\Ch observation of the circumnuclear environment of the Seyfert 2 galaxy NGC~2110, and have combined them with VLA radio and {\\it HST} \\ot data, in order to constrain the physical processes present. To summarize our conclusions:  \\begin{enumerate}  \\item We detect resolved, soft-band X-ray emission $\\sim4''$ North of the nucleus. The X-ray emission is spatially coincident with \\ot optical line-emission, but not with the radio emission, and instead appears to lie north of the radio ejecta.  \\item We consider several different models that may account for the extended northern X-ray emission. We find that it is possible that the radio jet has driven a strong shock through a series of multi-phase clouds of gas. We demonstrate that nuclear radiation scattered by a population of electrons may account for the extended X-ray emission, if the nucleus emits its radiation anisotropically. Finally, we show that photoionizaton of ambient gas by the nucleus may be sufficient to be responsible for the extended emission; high signal-to-noise \\Ch grating observations can directly test this model. Though energetically all of these mechanisms are plausible, we consider shock emission from the radio jet interacting with the ISM to be the favored model because of the spatial properties of the multiwavelength emission.  \\item We detect large-scale ($\\sim30''$) extended X-ray emission South of the nucleus in NGC~2110, which we model as a collisionally ionized plasma of temperature $\\sim0.9$ keV. The origin of this emission is unclear, although its prominence on the Southern side is consistent with the overall geometry of the galaxy.  \\item A comparison of the spatial distributions of X-ray, radio, and \\ot emission of NGC~2110 and the prototypical Seyfert 2 galaxy NGC~1068 leads us to suggest that different physical processes could produce the extended circumnuclear X-ray environments in Seyfert galaxies.  \\end{enumerate}"
    },
    "0608/astro-ph0608133_arXiv.txt": {
        "abstract": "We present high angular resolution ($2\\rlap.{''}8 \\times 1\\rlap.{''}7$) SiO J=5$\\rightarrow$4; $v=0$ line observations of the OMC1S region in the Orion Nebula made using the Submillimeter Array (SMA). We detect for the first time a cluster of four compact bipolar and monopolar outflows that show high, moderate and low velocity gas and appear to be energized by millimeter and infrared sources associated with this region. The SiO molecular outflows are compact ($<$ 3500 AU), and in most of the cases, they are located very close to their exciting sources. We thus propose that the SiO thermal emission is tracing the youngest and most highly excited parts of the outflows which cannot be detected by other molecules. Moreover, since the ambient cloud is weak in the SiO line emission, these observations can reveal flows that in other molecular transitions will be confused with the ambient velocity cloud emission. Analysis of their positional-velocity diagrams show that some components of these outflows may be driven by wide-angle winds very close to the exciting object. Finally, we find that some of these SiO outflows seem to be the base of powerful Herbig-Haro jets and large-scale molecular flows that emanate from a few arcseconds around this zone. In particular, we find a strongly excited SiO bipolar outflow with a P.A. of $\\sim$ 100$^{\\circ}$, that is likely energized by the luminous ($\\sim$ 3 $\\times$ 10$^3$ L$_{\\odot}$) infrared protostar \"B\" and could be the base of the remarkable object HH269. ",
        "introduction": "The OMC1S region (Orion Molecular Cloud 1 South) located about 1$'$ south-east of the Trapezium region, is one of the three most important zones of star formation (BN-KL, The Trapezium and OMC1S) in the Orion Nebula. This region indeed seems to be the youngest and most active zone of star formation in the nebula, showing extremely well collimated molecular outflows, and at least a half-dozen extended and powerful HH flows emanating from it (Schmid-Burgk et al. 1990, Bally, O'Dell \\& McCaughrean 2000, Rosado et al. 2001, O'Dell \\& Doi 2003, Zapata et al. 2005).  The OMC1S region was first identified as a star formation region through measurements of NH$_{3}$ (J,K)= (1,1) and (2,2) lines (Ziurys et al. 1981) and through 400 $\\mu$m dust continuum images (Keene, Hildebrand \\& Whitcomb 1982). Later, NH$_3$ line observations showed very high gas kinetic temperature ($\\geq$ 100 K), possibly due to the presence of strong outflows associated with the region (Batrla et al. 1983; Mauersberger et al. 1986; Ziurys, Wilson \\& Mauersberger 1990, Wiseman \\& Ho 1998). Many molecular lines have been detected toward the OMC1 South region: SiO, SO$_2$, H$^{13}$CO$^+$, CH$_3$OH, CH$_3$CN, HC$_3$N, $^{34}$CS, H$_2$O (as maser emission), $^{12}$CO and $^{18}$CO (Batrla et al. 1983; Mundy et al. 1986; Ziurys, Wilson \\& Mauersberger 1990; Schmid-Burgk et al. 1990; McMullin et al. 1993, Gaume et al. 1998 and Rodr\\'\\i guez-Franco et al. 1999a). However, as all of the species were detected with single-dish observations of moderate angular resolution, their relation to the embedded sources was not clear.  In particular, the thermal emission of the SiO line has provided evidence for the strong outflow activity associated with this region. Ziurys, Wilson \\& Mauersberger (1990) found a low velocity (10 km s$^{-1}$) and extended bipolar SiO(J=5-4; v=0) outflow with a length of $\\sim$ 30${''}$ in the North(blueshifted)-South(redshifted) orientation, that is centered at the position R.A.[J2000]=$05^h~35^m~12\\rlap.{''}8$ and Dec.[J2000]= $-05^\\circ~24'~11''$. McMullin et al. (1993) found strong thermal SiO(J=2-1; v=0) emission using the Berkeley-Illinois-Maryland Association (BIMA) interferometer. This thermal emission was more spatially confined to the west of OMC1S showing very broad ($\\sim$ 15 km s$^{-1}$) wings, suggesting the presence of energetic outflow activity in the region.  The production of the silicon monoxide molecule (SiO) along outflows is mainly attributed to the destruction of the dust particles in strong shocks. In these violent shocks, the grain-grain collisions and thermal sputtering inject refractory elements (Si) into the gas phase, permitting a quick formation of SiO through reactions involving the OH molecule (Schilke et al. 1997; van Dishoeck and Blake 1998). Thus, the abundance of the SiO molecules can be enhanced by several orders of magnitude with respect to the quiescent ambient cloud values (e. g. Guilloteau et al. 1992).  Recent surveys of proper motions in the Orion Nebula (Bally et al. 2000; O'Dell \\& Doi 2003) revealed that at least six relatively large Herbig-Haro outflows (HH 202, HH 269, HH 529, HH 203/204, HH 530, and possibly HH 528) are originating from a region only a few arcseconds across in size located in OMC1S. O'Dell \\& Doi (2003) located this region at R.A.[J2000]=5:35:14.56, Dec.[J2000]=-5:23:54.0 with a positional error of $\\pm$1.5$''$, and they referred to it as the Optical Outflow Source (OOS). The relation of these half-dozen Herbig-Haro flows with the molecular outflows found toward OMC1S has remained uncertain due to the poor angular resolution of the molecular observations. Moreover, deflections of those optical flows by their parental molecular cloud can be another problem in the association of them with molecular flows or exciting objects. Some optical flows that have been reported in the literature to be deflected or bent are: HH 110 (Reipurth, Raga \\& Heathcote 1996), and HH34 (Eisl\\\"{o}ffel \\& Mundt 1997; Bally et al. 2006), as have the molecular outflows, NGC 2264G (Fich \\& Lada 1997), and  NGC1333 IRAS 4A (Choi 2005). This mechanism has been successfully modeled by Raga and Cant\\'o (1995) and Raga et al. (2002). Based on the HH 110 case, they modeled the collision of a supersonic, radiative jet with a dense cloud, and found that the spatial velocity difference observed between HH 270 and HH 110 agrees surprisingly well with the velocity decrease expected for a radiative jet/cloud collision at the deflected angle. Finally, the change in collimation from HH 270 to HH 110 also agrees well with the broadening of the post-collision jet predicted in the theoretical models.  In this paper, we present new sensitive and high angular resolution SiO J=5$\\rightarrow$4; $v=0$ SMA observations of the OMC1 South region that reveal a cluster of four hidden compact outflows that are present in this region, and which seem to be the base of some remarkable Herbig-Haro flows that also emanate from here. In $\\S$ 2 we describe our SiO observations. In $\\S$ 3 we report the results of the imaging in SiO thermal emission. Finally, a summary is given in $\\S$ 4. ",
        "conclusions": "We note that most of the SiO outflows are compact ($<$ 3500 AU  $\\times$ $<$ 1500 AU), located close to their proposed exciting objects, and display a high dispersion of velocities.  It is puzzling that in our previous CO observations we did not detect these three monopolar and bipolar SiO flows. A possible explanation for the non-detection of these new SiO outflows in the CO data would be that two of them are low velocity outflows. The 137-408 monopolar south outflow (0 to 27 km s$^{-1}$ for the southern component, and 30 to 50 km s$^{-1}$ for the northern component) and the 137-347 monopolar south-east outflow (5 to 35 km  s$^{-1}$), may be confused with extended CO J=2$\\rightarrow$1 emission at ambient velocities (-25 to 21 km s$^{-1}$), see Zapata et al. (2005). However, for the other two east-west SiO outflows (136-356 and 137-408) it is not clear why they were not detected as they have extremely high velocities ($>$ 70 km s$^{-1}$), and are apparently massive outflows. In particular, for the strong SiO outflow associated with 136-356 we estimate, from comparison with the data of Zapata et al. (2005), that an enhancement of SiO of order 50 with respect to the adopted value of [SiO/H$_2$]$\\simeq$1 $\\times$ 10$^{-7}$ is needed in order to explain the non detection of this SiO outflow in CO.  We also note that the blueshifted velocity component of the highly collimated CO outflow shows a faint extension between the second and third molecular \"bullets\", to the west of these components (see Figure 1). If one searches at the position of this weak CO feature in the SiO velocity-integrated intensity image, we find that this component coincides very well in position with the blueshifted side of the 136-356 bipolar outflow. Therefore, in the CO velocity-integrated intensity image, we may be seeing weak CO emission associated with the blueshifted component of the 136-356 SiO bipolar outflow. In such a way, this extension can be the key to explain the nature of those two new outflows. Probably, the SiO emission is tracing high density entrained gas with a broad range of velocities and closer to their exciting objects, associated perhaps with their primary jets (see Hirano et al. 2006). The CO emission, is possibly tracing low density entrained gas, with a less broad range of velocities as has been observed by Palau et al. (2006) in the HH211 flow. Following this premise, in these SiO observations, we have imaged unresolved emission from the jet itself, as observed by Hirano et al. (2005) and Palau et al (2006). However, for this mechanism to work, we need strong chemical mechanisms that would strongly favor the detection of each molecule in the two components of different density.  At this point, it is not well understood yet why the SiO is a selective tracer. In some cases it seems to be associated with the most excited parts of the flows (136-356 and 137-408 outflows; this paper, and HH211; Hirano et al. 2006; Palau et al. 2006) while in other cases it is associated with internal shocks in the outflows, (L1157; Gueth, Guilloteau and Bachiller 1998, L1448; Dutrey et al. 1997, and the 136-359 CO outflow; this paper). One possibility for unifying these two mechanisms would be if the SiO molecule is formed/excited in shocks that are propagating down the jet (\"internal working surfaces\"), as proposed by Arce et al. (2006).  In Table 3 we show a list of tentative associations between the HH objects, the molecular outflows, and their proposed exciting sources located in the OMC1S region."
    },
    "0608/astro-ph0608075_arXiv.txt": {
        "abstract": "{}{% The development of a general code for 3D \\lya\\ radiation transfer in galaxies to understand the diversity of \\lya\\ line profiles observed in star forming galaxies and related objects.} {% Using a Monte Carlo technique we have developed a 3D \\lya\\ radiation transfer code allowing for prescribed arbitrary hydrogen density, ionisation, temperature structures, and dust distributions, and arbitrary velocity fields and UV photon sources.} {% As a first test and application we have examined  the \\lya\\ line profiles predicted for several simple geometrical configurations and their dependence on the main input parameters. Overall, we find line profiles reaching from doubly peaked symmetric emission to symmetric Voigt (absorption) in static configurations with increasing dust content, and asymmetric red- (blue-) shifted emission lines with a blue (red) counterpart ranging from absorption to emission (with increasing line/continuum strength) in expanding (infalling) media. In particular we find the following results which are of interest for the interpretation of \\lya\\ profiles from galaxies. {{\\em 1)}} Standard \\lya\\ absorption line fitting of global spectra of galaxies may lead to an underestimate of the true hydrogen column density in certain geometrical conditions. {{\\em 2)}} Normal (inverted) P-Cygni like \\lya\\ profiles can be obtained in expanding (infalling) media from objects without any intrinsic \\lya\\ emission, as a natural consequence of radiation transfer effects. {{\\em 3)}} The formation and the detailed shape of \\lya\\ profiles resulting from expanding shells has been thoroughly revised. In particular we find that, for sufficiently large column densities (\\nh\\ $\\ga 10^{20}$ cm$^{-2}$), the position of the main \\lya\\ emission peak is quite generally redshifted by approximately twice the expansion velocity. This is in excellent agreement with the observations of $z \\sim 3$ LBGs, which show that \\lya\\ is redshifted by  $\\sim 2 \\vexp$, where $\\vexp$ is the expansion velocity measured from the interstellar absorption lines blueshifted with respect to the stellar redshift. This finding indicates also that large scale, fairly symmetric shell structures must be a good description for the outflows in LBGs. }{} ",
        "introduction": "\\label{s_intro}  The \\lya\\ line plays an important role in a variety of astrophysical problems, especially as a diagnostic tool to observe and study the high redshift universe. It is a simple redshift indicator for distant galaxies, a frequently used star formation rate diagnostic at high $z$, as well as an important tool probing the ionisation state of the intergalactic medium and hence the reionisation epoch. The \\lya\\ emission line is e.g.\\ a strong feature observed in nearby star forming galaxies, distant Lyman break galaxies (LBGs), sub-mm galaxies,  emission line selected galaxies (LAE, for Lyman-$\\alpha$ emitters), and in the enigmatic so called Lyman-$\\alpha$ blobs (LABs) whose nature remains debated (e.g.\\ Steidel \\etal\\ 2000, Dijkstra \\etal\\, 2005b).  Since the early suggestion of strong \\lya\\ emission from young high redshift galaxies by Partridge \\& Peebles (1967) and until the late 1990s, only few \\lya\\ emitters have been found (cf.\\ Djorgovski \\& Thompson 1992). This lack of \\lya\\ emission detection has triggered a variety of studies discussing the possible physical effects (mostly metallicity, dust, neutral gas kinematics, and geometry) which may significantly affect and suppress the \\lya\\ emission and the resonance line radiation transfer, thereby reducing the observed \\lya\\ intensity and destroying simple expected correlations, e.g.\\ between \\lya\\ intensity and metallicity,  \\lya\\ intensity and UV continuum flux and others (Meier \\& Terlevich 1981, Hartmann \\etal\\ 1988, Neufeld 1990, Charlot \\& Fall 1993, Valls-Gabaud 1993, Kunth \\etal\\ 1998, Tenorio-Tagle \\etal\\ 1999,  Mas-Hesse \\etal\\ 2003).  In the last few years, with the availability of deeper and wider surveys such as the Large Area Lyman Alpha (LALA) survey and the Subaru Deep Field survey, many emission galaxies have been detected (cf.\\ Hu \\etal\\ 1998, 2004; Kudritzki \\etal\\ 2000, Rhoads \\etal\\ 2000, Ouchi \\etal\\ 2003, Taniguchi \\etal\\ 2005). Although the majority of these distant \\lya\\ emitters shows rather simple asymmetric line profiles, the overall diversity of the observed \\lya\\ line shapes, both from star forming galaxies in the nearby universe and at high-$z$, is quite heterogeneous and complex. The observed line profiles include schematically pure Voigt absorption profiles, P-Cygni profiles, double peak profiles, pure (symmetric) emission line profiles, and combinations thereof (see e.g.\\ Kunth \\etal\\ 1998, Mas-Hesse \\etal\\ 2003, Shapley \\etal\\ 2003, M\\\"oller \\etal\\ 2004, Venemans \\etal\\ 2005, Wilman \\etal\\ 2005,  Noll \\etal\\ 2004, Tapken \\etal\\ 2004, Tapken 2005).  Although in principle the main physical processes shaping the \\lya\\ line are known, in practice the inferences drawn so far from \\lya\\ observations rely mostly on rather simple measurements (e.g.\\ line flux) or on oversimplified  Voigt-profile fits, which often have no strong physical motivations. For example, for a \\lya\\ line profile formed purely within a galaxy (i.e.\\ neglecting subsequent alterations from the intergalactic medium and/or intervening clouds), it is physically inconsistent to fit one or several Voigt profiles without making strong implicit assumptions on the geometry of the neutral gas. In such a case, due to the resonance scattering nature of \\lya\\, it is in fact unlikely that the resulting emergent line profile is actually a Voigt profile and a detailed radiation transfer calculation has thus to be carried out to predict the proper shape of the emergent resonance line profile. In general, quantitative simulations for appropriate geometries and gas kinematics, which take properly into account the main physical processes of \\lya\\ line formation and radiation transfer are therefore needed for a better understanding of the variety of observed \\lya\\ line profile morphologies. The physical properties of the \\lya\\ emission mechanism and the ones of the ambient gas hosting the \\lya\\ emitter, which both shape the observed \\lya\\ line, must be properly modeled to investigate the correspondent impact on the emergent line profile and their possible degeneracies. This in turn will provide useful hints to guide the interpretation of the observed profiles and to gain insight on the physical properties of the \\lya\\ emitters and of their environment.  Analytic solutions for the \\lya\\ radiation transfer problem have been derived for simple geometries. Neufeld (1990) has extensively studied the case of static plane parallel slabs, yielding important insight on the line formation mechanism and providing solutions for configuration including dust, Bowen fluorescence and \\lya\\ pumping of H$_2$ Lyman band lines. The case of a static, uniform sphere has been recently studied by Dijkstra \\etal\\ (2005a). Loeb \\& Rybicky (1999) and Rybicky \\& Loeb (1999) have derived solutions for \\lya\\ scattering in a Hubble flow. However, more general geometries and velocity fields do not allow for an analytic solution and require alternative approaches.  Over the last few years, several groups have developed new numerical algorithms, mostly based on Monte Carlo techniques (Spaans 1996, Ahn \\etal\\ 2001, 2002; 2003, Ahn 2004; Zheng \\& Miralda-Escud\\'e 2002; Richling \\etal\\ 2001; Richling 2003; Kobayashi \\& Kamaya 2004; Cantalupo \\etal\\ 2005; Dijkstra \\etal\\ 2005ab; Hansen \\& Oh 2006; Tasitsiomi 2006). Some of these codes have been specifically designed for and can be reliably applied only to particular configurations: relatively low column densities (Richling \\etal\\ 2001; Richling 2003), Hubble flows (Kobayashi \\& Kamaya 2004), 1D geometry (Ahn \\etal), spherically symmetric configurations Dijkstra \\etal\\ 2005ab); others are strongly tailored towards cosmological simulations (Cantalupo \\etal\\ 2005, Tasitsiomi 2006) and can deal with clumpy/inhomogeneous media (Spaans, 1996; Richling, 2003; Hansen \\& Oh 2006). In addition, the effect of dust absorption, which is one of the most important factor affecting \\lya\\ transmission, is treated only in some of the codes above.  In any case, none of these studies has attempted to explain systematically the observed variety of \\lya\\ line profile morphologies. Furthermore, none of the above schemes has so far attempted a detailed modeling of individual galaxies, taking into account the available observational constraints,  i.e.\\ constraints on the stellar populations, the ionised and neutral interstellar medium, on dust extinction and including their spatial distribution and kinematics. All this lack in the actual theoretical modeling needs to be filled up, in order to extract some information from the huge reservoir contained in the available observational data. With these objectives in mind, and with the main aim of improving our understanding of \\lya\\ in both nearby and distant starburst galaxies, we have developed a general-purpose 3D \\lya\\ radiation transfer code applicable to arbitrary geometries and velocity fields.  In the present paper we provide a description of the code and test its validity against known solutions and results from other codes reported in the literature. Exploring different geometries, dust-free and dusty media, and different input spectra (e.g.\\ line emission, or continuum + line), we examine the resulting line profiles and their dependence on various physical parameters. Our immediate goals are to obtain an overview over the possible \\lya\\ line profile morphologies, and to gain physical insight into the processes governing them. Applications to observed galaxies and other simulations will be presented later.   The remainder of the paper is structured as follows. A description of the radiation transfer code is given in Sect.\\ \\ref{s_code}. Tests of the code and results for simple geometrical configurations (slabs, infalling/expanding halos) are presented in Sect.\\ \\ref{s_test}. In Sect.\\ \\ref{s_slab} we comment on the formation of damped (Voigt) \\lya\\ profiles and related profiles. Spherically expanding dust-free or dusty shells are re-examined in Sect.\\ \\ref{s_shell}. An overview of the predicted \\lya\\ line profile morphologies and qualitative comparisons with observations is given in Sect.\\ \\ref{s_discuss}. Our main conclusion are summarised in  Sect.\\ \\ref{s_conclude}. ",
        "conclusions": "\\label{s_conclude} We have developed a new general 3D \\lya\\ radiation transfer code allowing for arbitrary hydrogen density, ionisation \\& temperature structures, dust distributions, and velocity fields using a Monte Carlo technique (Sect.\\ \\ref{s_code}). Currently the code works on a cartesian grid and it can handle \\lya\\ transfer problems in objects with arbitrary hydrogen column densities. The main radiation--matter interaction processes included in the code are \\lya\\ scattering, dust scattering, and dust absorption.  UV/\\lya\\ photon sources with an arbitrary spatial distribution and arbitrary spectra can be treated. The main direct observables predicted by our simulations are ``global'' (i.e.\\ integrated over all viewing angles) \\lya\\ line profiles, spectra for different viewing directions and opening angles, and surface brightness maps at different wavelengths (or integrated over \\lya).  The code has been tested successfully and compared against analytical results and results from other simulations (e.g.\\ Neufeld 1990, Ahn \\etal\\ 2001, Zheng \\& Miralda-Escud\\'e 2002, Richling \\etal\\ 2003, Dijkstra \\& Haiman 2005a) for various geometrical configurations including plane parallel slabs, disks, and expanding or infalling halos, and for computations with and without dust. Overall an excellent agreement is found, except in cases where assumptions made in the analytical calculation break down (see Sect.\\ \\ref{s_test}).  With the aim of understanding and interpreting ultimately the variety of \\lya\\ profiles observed in nearby and distant starbursts, we have examined the \\lya\\ line profiles predicted for several simple geometrical configurations and their dependence on the main parameters such as the \\hi\\ column density, temperature, velocity, and dust content. We have considered slabs with internal sources, disks, expanding and infalling halos (Sect.\\ \\ref{s_test}), externally illuminated slabs (Sect.\\ \\ref{s_slab}), and expanding shells (Sect.\\ \\ref{s_shell}). For the source spectrum we have considered the limiting cases of pure monochromatic line emission, a pure continuum, or intermediate cases with different line strengths. The two latter cases have rarely been discussed in the literature so far.  Schematically, the following morphologies are found (cf.\\ Table \\ref{t_overview}): profiles reaching from doubly peaked symmetric emission to symmetric \"Voigt-like\" (absorption) profiles in static configurations with increasing dust content, and asymmetric red- (blue-) shifted emission lines with a blue (red) counterpart ranging from absorption to emission (with increasing line/continuum strength) in expanding (infalling) media. In principle, symmetric or nearly symmetric profiles are only obtained for static geometries or for small systematic velocities. In practice, however, this may be  altered by \\lya\\ scattering/absorption in the intervening IGM, especially at high  redshift.  Some specific results are worthwhile summarising here: \\begin{itemize} \\item Pure \\lya\\ absorption line profiles may be observed in ``integrated''/global spectra  of objects with static geometries. In this case the naturally arising double \\lya\\ emission peaks have been suppressed by dust absorption. Standard Voigt profile fitting of such profiles will significantly underestimate the true hydrogen column density (see Sect.\\ \\ref{s_slab}).  \\item It should be noted that normal (inverted) P-Cygni like \\lya\\ profiles can be obtained in expanding (infalling) media from objects without any intrinsic \\lya\\ emission, as a natural consequence of radiation transfer effects redistributing UV (continuum) photons from the \\lya\\ line center to the red (blue) wing (e.g.\\ Figs.\\ \\ref{movie_halo}, \\ref{movie_shell}). Adding dust in such cases will progressively transform the P-Cygni profile into a broad asymmetric absorption profile.  \\item Our 3D simulations of expanding shells lead us to revise the earlier results of Ahn and collaborators (Ahn \\etal\\ 2003, Ahn 2004) and to clarify the radiation transfer processes explaining the detailed shape of \\lya\\ profiles predicted from expanding shells. In particular we found that, for sufficiently large column densities (\\nh\\ $\\ga 10^{20}$ cm$^{-2}$), the position of the main \\lya\\ emission peak is quite generally redshifted by approximately twice the expansion velocity (Sect.\\ \\ref{s_shell}). This is in excellent agreement with the observations of $z \\sim 3$ LBGs from Shapley \\etal\\ (2003), which show that \\lya\\ is redshifted by  $\\sim 2 \\vexp$, where $\\vexp$ is the expansion velocity measured from the interstellar absorption lines blueshifted with respect to the stellar (i.e.\\ galaxian) redshift. This finding indicates also that large scale, fairly symmetric shell structures must be a good description for the outflows in LBGs. \\end{itemize}  As already clear from earlier investigations, radiation transfer effects even in the simple configurations just mentioned give rise to a complex morphology of \\lya\\ line profiles and some degeneraciesare found e.g.\\ between the main parameters such as the hydrogen column density \\nh, the Doppler parameter $b$, and also the intrinsic emission line width. However, codes as the one developed and presented here allow one to undertake detailed \\lya\\ line profile fitting to determine the physical properties of the  emitting gas and of the surrounding  \\hi\\ and dust and to quantify possible degeneracies. With such tools it is now possible to fully exploit the information encoded in \\lya\\ profiles without recurring to physically  unmotivated superpositions of emission and Voigt absorption line profiles, which -- except in specific geometric configurations -- generally do not properly account for radiation transfer effects. Such detailed studies of nearby and distant starforming galaxies will be presented in subsequent publications. Various other applications, combining also our {\\em MCLya} code with the 3D radiation transfer and photoionisation code CRASH (Maselli \\etal\\ 2003) are foreseen."
    },
    "0608/astro-ph0608596_arXiv.txt": {
        "abstract": "We consider a model in which massive stars form in a self-gravitating accretion disk around an active galactic nucleus. These stars may evolve and collapse to form compact objects on a time scale shorter than the accretion time, thus producing an important family of sources for LISA. Assuming the compact object formation/inspiral rate is proportional to the steady-state gas accretion rate, we use the observed extra-galactic X-ray luminosity function to estimate expected event rates and signal strengths. We find that these sources will produce a continuous low-frequency background detectable by LISA if more than $\\gtrsim 1\\%$ of the accreted matter is in the form of compact objects. For compact objects with $m \\gtrsim 10 M_\\odot$, the last stages of the inspiral events should be resolvable above a few mHz, at a rate of $\\sim 10-100$ per year. ",
        "introduction": "One of the major challenges that will face gravitational wave (GW) astronomers is the successful discrimination between instrumental noise, stochastic GW backgrounds, and individual, resolvable GW sources. For example, the population of galactic white dwarf binaries in close orbits will provide a major contribution to the LISA noise curve in the $0.1-1$ mHz band \\cite{BH}. At the same time, this confusion ``noise'' can also be treated as a signal, and its shape and amplitude will provide important information about the distribution and properties of white dwarf binaries in the galaxy.  It has recently been proposed that, in the self-gravitating accretion disks of active galactic nuclei (AGN), massive stars could form and evolve, eventually collapsing into compact objects and merging with the central black hole \\cite{Goodman:2003sf,Levin}. At different stages in the inspiral evolution, this population will contribute to the GW background confusion or alternatively produce individual, resolvable chirp signals. It is therefore a matter of theoretical and practical interest to understand the nature of such a population.  In this paper we derive a relationship between the observable electro-magnetic (EM) emission and the predicted GW emission from AGN, following the procedures and summarizing the results outlined in our longer paper \\cite{sigl}. In particular, we use the hard X-ray luminosity function of Ueda et al.~\\cite{uaom} to infer the accretion history of supermassive black holes (SMBHs) out to redshifts of $z\\sim 3$. Then we assume a few simple scaling factors, such as the average (Eddington-scaled) accretion rate and the efficiency of converting accretion energy to X-rays, and derive the GW spectrum that might be seen by LISA.  Depending on the specific model parameters, we find this background could be an important class of LISA sources, similar in strength and event rates to extreme mass-ratio inspirals from captured compact objects \\cite{barrack}. As in those sources, here too it is a matter of preference as to whether the background should be thought of as signal or noise. But for higher masses (perhaps as large as $m\\simeq 10^5 M_\\odot$ \\cite{Goodman:2003sf}), disk-embedded compact objects should produce individual, resolvable inspiral events with high signal-to-noise over a wide band of frequencies. ",
        "conclusions": ""
    },
    "0608/astro-ph0608360_arXiv.txt": {
        "abstract": "We review theoretical developments in studies of dense matter and its phase structure of relevance to  compact stars. Observational data on compact stars, which can constrain the properties of dense matter, are presented critically  and interpreted. ",
        "introduction": "Since their discovery almost 40 years ago, neutron stars have been recognized as promising laboratories for studying matter under extreme conditions. All the ambient conditions that characterize these objects tend to the extreme. The typical density inside a neutron star is comparable to the density inside nuclei, $\\rho_\\mathrm{nuclear}=2.5 \\times 10^{14}$ g cm$^{-3}$, which corresponds to a number density of baryons $n_B \\simeq 0.15 $ fm$^{-3}$. Under these conditions the nucleons are very degenerate, with a typical Fermi momentum $k_F \\sim (3 \\pi^2 n_B)^{1/3}$ in the range of $300$--$600$ MeV. In the interior, the interaction energy of the nucleons is several times larger than their Fermi degeneracy energy $E_F = k_F^2/2m \\sim 60$--$150$ MeV. The baryon chemical potential, which is defined as the energy needed to introduce one unit of baryon number, can be $500$--$1000$ MeV larger, at the center of the star, than the rest mass of the nucleon. There the internucleon distance becomes comparable to their intrinsic size of $1$ fm. Under such extreme conditions several new states of matter, containing new high-energy degrees of freedom, may be favored. Novel hadronic phases -- which contain, in addition to nucleons, a Bose condensate of pions or kaons, and/or hyperons -- have been proposed and studied in some detail. There is also the exciting possibility that the core may contain deconfined quark matter, much like the quark-gluon plasma being probed in terrestrial experiments such as the Relativistic Heavy Ion Collider (RHIC). Although model descriptions of these various phases are still in rudimentary form, in several cases qualitative differences between the exotic phases and nuclear matter have been identified. A quantitative understanding of these trends and, in particular, how they could affect observable aspects of compact star structure and evolution is an exciting area of research. A primary purpose of this review is to convey some of this excitement in a pedagogic style to make it accessible to graduate students. This review is not intended to be a comprehensive account of all recent developments. Instead, we have chosen topics that serve (in our view) as pedagogic tools that help illustrate key concepts of relevance to dense matter in compact stars.  \\begin{figure}[h] \\begin{center} \\includegraphics[width=0.8\\textwidth,angle=-0]{NStar.eps} \\caption{Inside a Neutron Star.} \\label{fig:Nstar} \\end{center} \\end{figure}  Figure \\ref{fig:Nstar} illustrates a theoretician's view of a neutron star. The environment just outside the star, characterized by large magnetic (and electric) fields and  temperatures, is extreme from plasma physics point of view. At the surface of the star we expect a very thin atmosphere composed of hydrogen and in some cases perhaps a mix of heavy elements, or even a condensed magnetic surface. This surface is of utmost observational importance because it allows us to determine the temperature, and possibly the radius, of the neutron star, an issue we describe in Section~\\ref{sec:obs-cooling}. An envelope exists just below the atmosphere where matter is not yet fully degenerate and, with a thickness of a few tens of meters, it acts as a thermal insulator between the hot interior and the surface. The outer $500$--$1000$ meters of the star, its crust, contains nuclei, forming a lattice immersed in a quantum liquid of neutrons, most probably in a superfluid state (Figure~\\ref{fig:Nstar}, insert A). Owing to the rotation of the star, the neutron superfluid forms vortices, which very likely pin on the nuclei and participate in the glitches. As we go deeper into the crust, when approaching nuclear density there is a first-order phase transition from the inhomogeneous regime of the crust to the homogeneous core. In this regime, termed the nuclear pasta (Figure~\\ref{fig:Nstar}, insert B), as density grows, nuclei are increasingly elongated (the spaghetti phase), then form bidimensional structures (lasagnas) -- with the space between them still filled by the superfluid neutrons -- untill the geometry inverts in the Swiss cheese phase, where bubbles of neutrons are immersed in the almost homogeneous neutron+proton liquid. The physical reasons for the peculiar aspect of this phase transition are elucidated in Section~\\ref{sec:pasta}. At a density of approximately $0.6 \\; \\rho_\\mathrm{nuclear}$, we enter the core where the neutron superfluid coexists with the proton fluid, which is probably a Type II superconductor that reacts to the neutron star magnetic field by confining it into fluxoids (Figure~\\ref{fig:Nstar}, insert C). A few kilometers below may be the inner core, marked in Figure~\\ref{fig:Nstar} by ``?''. This question mark is the focus of this review.   In Section~\\ref{sec:section2} we present a selected set of observational constraints. These include mass (Section~\\ref{sec:masses}) and radii (Section~\\ref{sec:obs-cooling}) measurements that probe the equation of state (EoS), i.e., the high-energy properties of dense matter. Measurements of temperatures and thermal luminosities of isolated neutron stars and accreting ones in binary systems are presented in Section~\\ref{sec:obs-cooling} and Section~\\ref{sec:SXRT}, respectively. These observations provide information about the structure of the stars, i.e., the EoS, but also about the low-energy response functions, such as specific heat, transport properties, and neutrino emission, which probe the low-energy (keV--MeV) excitation spectrum. In Section~\\ref{sec:section3} we discuss the essential physics of dense nuclear matter and describe the current state of model calculations of the nuclear EoS and low-energy response. In Section~\\ref{sec:section4} we describe theoretical models for phase transitions to novel states of matter such as Bose-condensed hadronic matter, hyperonic matter, and quark matter at supranuclear density. Section~\\ref{sec:section5} confronts theoretical models of neutron star structure and of evolution derived from the studies described in the two previous sections with the observational baggage acquired in Section~\\ref{sec:section2}. Finally, we offer some conclusions in Section~\\ref{sec:section6}. ",
        "conclusions": "\\label{sec:section6}  To infer the phase structure and properties of matter at extreme density through observations of compact stars is a compelling but difficult task. Despite the rather rudimentary state of quantitative theoretical calculations in describing dense matter, the observational data are already providing guidance to and hinting at possible tensions between different models. In particular the prospect for a firm mass measurement of a heavy neutron star (or a low-mass black hole) and accurate radius measurements are excellent. Such measurements will directly confirm or rule out the possibility for a strong phase transition to a soft EoS in the vicinity of nuclear density, and presently available values already seem to disfavor a soft EoS. Observations of cooling neutron stars and the realization that superfluidity and superconductivity play a dominant role in neutrino cooling of the core has sparked renewed interest in theoretical studies of pairing correlations in dense matter. Furthermore, the connection between seemingly different observables such as long-term neutron star cooling, quiescent luminosity observed in LMXBs, and superbursts are indicating new trends that point to a diversity in the thermal evolution of neutron stars, with growing evidence in favor of the occurence of some fast neutrino emitting process(es), at least in heavy, $>1.4 M_\\odot$, neutron stars. These developments call for a more global approach in the theory and simulations to confront different observables within a unified framework so that the sensitivity to the underlying dense matter physics can be properly assessed.  Theoretical progress in describing dense nuclear matter and new insights into using terrestrial experiments on nuclei to better constrain EoS models provide compelling reasons to expect a quantitative description of neutron-rich matter for the densities of interest. Throughout this review we highlighted important qualitative differences and possible similarities between nucleon matter and other exotic phases. Minimizing the theoretical uncertainty for the predictions of the nucleon matter is crucial to confront observables. For example, a precise determination of the density dependence of the nuclear symmetry energy can reduce the uncertainty simultaneously in radii and cooling rates predicted for the different nuclear EoSs.   The situation with exotic states of matter is less clear. Although softening and enhanced cooling appear to be common trends, there are important exceptions and quantitative estimates are often very model dependent. Nonetheless, as mentioned above, radius measurement smaller than 10 km or the discovery of low-mass black hole candidate with  mass $ \\simless 2 M_\\odot$  would be very suggestive of a phase transition. Better experimental or theoretical constraints on the hyperon-nucleon and kaon-nucleon interactions are necessary to ascertain if novel hadronic phases are relevant.  Quark matter remains an especially interesting possibility. Color superconductivity is relevant, especially for the calculation of neutrino processes and cooling. However, an understanding of how a finite strange quark mass affects the phase structure is needed because this can result in vastly different low energy properties."
    },
    "0608/astro-ph0608683_arXiv.txt": {
        "abstract": "{The detection of the rotational lines of CO in proto-galaxies in the early Universe provides one of the most promising ways of probing the fundamental physical properties of a galaxy, such as its size, dynamical mass, gas density, and temperature. While such observations are currently limited to the most luminous galaxies, the advent of ALMA will change the situation dramatically, resulting in the detection of numerous normal galaxies at high redshifts. } {We investigate the morphology and strength of the CO rotational line emission emerging from $z\\sim 3$ progenitor systems of normal field galaxies seen in the present-day Universe, examine how well ALMA will be able to detect such systems and how accurately CO line widths, gas morphologies and ultimately dynamical masses can be inferred. } {Maps and spectra of rotational CO line emission were calculated from a cosmological N-body/hydrodynamical TreeSPH simulation of a $z\\sim 3$ \"Lyman break galaxy\" of UV luminosity about one order of magnitude below $L^*$. To simulate a typical observation of our system with ALMA, we imposed characteristic noise, angular, and spectral resolution constraints. } { The CO line properties predicted by our simulation are in good agreement with the two Lyman break systems detected in CO to date. We find that while supernovae explosions from the ongoing star formation carve out large cavities in the molecular ISM, they do not generate large enough gas outflows to make a substantial imprint on the CO line profile. This implies that for most proto-galaxies -- except possibly the most extreme cases -- stellar feedback effects do not affect CO as a dynamical mass tracer. } { Detecting CO in sub-$L^*$ galaxies at $z\\simeq 3$ will push ALMA to the limits of its cababilities, and whether a source is detected or not may depend critically on its inclination angle. Both these effects (sensitivity and inclination) will severely impair the ability of ALMA to infer the gas kinematics and dynamical masses using line observations. } ",
        "introduction": "Since the first discovery of CO in IRAS\\,F10214$+$4724 at $z=2.29$ (Brown \\& Vanden Bout 1992), observations tracking the molecular gas (H$_2$) in distant galaxies have taken on an increasingly important role as a way to study galaxy formation and evolution (Solomon \\& Vanden Bout 2005). Prime reasons for this are that CO provides an estimate of the total molecular gas reservoir available for star formation and black hole accretion and can be used as a tracer of the dynamical mass of a galaxy (Bryant \\& Scoville 1996). The latter is particularly important during the formative stages of a galaxy, where extreme dust-obscuration, as well as outflows, can make dynamical mass estimates derived from optical/near-IR spectroscopy unreliable.  Today, only the most extreme objects at high-$z$, such as QSOs and luminous submillimeter-selected galaxies (SMGs -- Blain et al.\\ 2002), are detectable in CO. In contrast, the more numerous, but also less massive and presumably less gas-rich, Lyman break galaxies  (LBGs -- Giavalisco 2002) are generally too faint to be detected in CO with current instruments (cf.\\ Baker et al.\\ 2004a). With the advent of the Atacama Large Millimeter Array (ALMA), this situation may change, making it important that realistic numerical models are developed in order to interpret the data. In spite of this, the amount of numerical/theoretical work on simulating the expected CO emission from primordial galaxies has been limited to a few pioneering studies (Silk \\& Spaans 1997; Combes, Maoli \\& Omont 1999). While these initial efforts were mainly focused on the detectability of high-$z$ CO lines with future mm-facilities, using 'idealized' generic starburst galaxies as templates, more recently, however, Narayanan et al.\\ (2006) simulated the CO emission from two merging gas-rich disk galaxies in detail, and studied the effects of a massive black hole on the gas dynamics of the system as revealed by the CO line emission.\\\\ \\indent In this letter we present CO intensity maps and line profiles of a cosmological self-consistent simulation of a merger system, representative of the LBG population at $z\\sim 3$. We explore the feedback effects from supernovae (SNe) on the molecular gas, and in particular whether they affect dynamical mass estimates based on CO. The simulation is subjected to the sensitivity and resolution constraints imposed by ALMA, and we explore the effects this have on our results. We adopt a flat cosmology with $\\Omega_m=0.27$, $\\Omega_\\Lambda=0.73$ and $H_0=71$\\,km\\,${\\mbox{s}^{-1}}$\\,Mpc$^{-1}$ (Spergel et al.\\ 2003). ",
        "conclusions": "\\label{section:results} \\subsection{Intensity maps and line profiles} The resulting CO 4--3 intensity maps of the inner 6\\,kpc$\\times$6\\,kpc (corresponding to $0.4\\arcsecs \\times 0.4\\arcsecs$ at $z=3.26$) viewed along the $x, y$ and $z$ lines of sight are shown in Fig.\\ \\ref{figure:co-maps}a-c. The CO 4--3 line is the lowest transition which will be accessible to ALMA (band 3) at $z=3.26$. It is seen that the bulk of the CO emission comes from the two merging disk systems, reflecting the fact that the molecular gas is primarily situated in these two systems, and reaches its highest densities there ($n(\\mbox{H}_2)\\sim 10^{6}\\,\\mbox{cm}^{-3}$). However, faint CO emission is also seen to originate from filamentary structures extending out from the main system. These are merger-induced tidal streams of diffuse ($n(\\mbox{H}_2)\\sim 10^{1-2}\\,\\mbox{cm}^{-3}$) molecular gas, hence the low CO surface brightness. The diffuse gas accounts for about 1 and 10 per-cent of the total CO 4--3 and 1--0 luminosities, respectively. We find a global CO 4--3/1--0 line luminosity ratio of $\\sim 0.6$ from our simulation, which is marginally subthermal. This, however, is merely an 'aperture effect', due to the fact that the extended, diffuse gas contributes more strongly to the CO 1--0 than the 4--3 line emission. The CO emission shows significant substructure within each of the merging galaxies. In the face-on view of the central galaxy (Fig.\\ \\ref{figure:co-maps}a), the ISM is seen to consist of multiple cavities of low CO emission, surrounded by filamentary regions where the CO emission is strong. These structures are created by SNe explosions, which carve out 'holes' in the ISM, while at the same time creating regions of dense gas as the blastwave ploughs through the surrounding medium, compressing it in the process. Thus in our simulation, SNe appear to have a significant impact on the distribution of molecular gas within a galaxy, which is consistent with the notion that star formation sites are born in the wake of SNe.  The CO spectra integrated over the entire region are shown in black in Fig.\\ \\ref{figure:co-maps}d-f. The first thing to notice is the strong dependence of the line shape on the viewing angle. As expected, the spectrum in Fig.\\ \\ref{figure:co-maps}d is dominated by a very narrow ($\\Delta V_{\\mbox{\\tiny{FWHM}}}\\simeq 50\\,\\mbox{km}\\,\\mbox{s}^{-1}$) and smooth single profile, which corresponds to the main disk viewed face-on. In contrast, the spectra in Fig.\\ \\ref{figure:co-maps}e-f have double-horn profiles, which is normally taken to indicate orbital motion in a disk or a merger. While such double-horn profiles observed towards local ULIRGs have been shown to be due to a compact circumnuclear disk of molecular gas (e.g.\\ Downes \\& Solomon 1998), the issue is still controversial for high-$z$ objects (Ivison et al.\\ 2001; Downes \\& Solomon 2003) In our case, it is the merger-nature of the system which is responsible for the overall double-horn profile. Although, we note that each disk system (whose individual spectra are shown in red(A) and yellow(B)), when viewed edge-on, also give rise to a double-peaked profile, albeit a less distinct one. The above illustrates that great care needs to be taken when interpreting high-$z$ observations, and underlines the need for high-resolution observations capable of resolving the emission. Finally, we note that the fact that the line profile in Fig.\\ \\ref{figure:co-maps}d is smooth, strongly suggest that while the molecular gas can be perturbed by SNe, the velocities so attained and/or the amount of gas affected are insufficient to make a significant imprint on the CO line profile. Of course this does not rule out the possibility that in galaxies with more prodigal star formation rates, SNe, and possibly even stellar winds, could have a significant effect on the observed CO line profile. Also, we note that Narayanan et al.\\ (2006) found that feedback from an AGN in massive merger systems can produce outflows of molecular gas, which are substantial enough to be detectable in CO.  \\begin{figure*} \\centering \\includegraphics[width=0.500\\hsize]{8013fig1.eps}~~~~\\includegraphics[width=0.475\\hsize]{8013fig2.eps} \\caption{{\\bf a-c):} Velocity-integrated CO 4--3 intensity maps (in units $\\mu$Jy\\,km\\,s$^{-1}$) of our simulation as seen from three different lines of sight. Each panel is 6\\,kpc on a side, with a pixel size of 10\\,pc. {\\bf d-f):} The corresponding CO spectra integrated over the entire field (black curve), the A component (red) and its companion B (yellow). {\\bf g-i):} CO 4--3 maps (in units $\\mu\\mbox{Jy}\\,\\mbox{beam}^{-1}$) of our simulated system in a-c) after 24\\,hrs of integration with ALMA at $0.1\\arcsecs$ resolution. The maps were obtained by averaging the channel-maps from -200\\,km\\,s$^{-1}$ to 200\\,km\\,s$^{-1}$. The contours are $-2, 2, 3, 4\\,\\sigma$, where $\\sigma = 10\\,\\mu$Jy\\,beam$^{-1}$, where negative contours are shown as dashed curves. {\\bf j-l):} The corresponding spectra (binned to 20\\,km\\,s$^{-1}$) extracted from the centroid position of each source.} \\label{figure:co-maps} \\end{figure*}   \\subsection{Comparison with observations} Does the predicted CO properties of the simulated proto-galaxy resemble those of observed $z\\sim3$ LBGs? Only two LBGs have been detected in CO to date: the gravitationally lensed systems MS\\,1512$-$cB58 ($z=2.73$) and J2135$-$0101 ($z=3.07$) (Baker et al.\\ 2004; Coppin et al.\\ 2007). The two LBGs have star formation rates in the range $\\sim 20-60\\,\\Msolar\\,$yr$^{-1}$, i.e.\\ comparable to our simulated system. The peak CO 3--2 line intensities towards cB58 and J2135 are $\\sim (2\\,$mJy$)/\\mu$ and $\\sim (3\\,$mJy$)/\\mu$, where the lensing amplification factors are $\\mu \\simeq  31.8$ and $\\simeq 8$, respectively. Thus, their intrinsic CO line fluxes -- even if the CO 4--3/3--2 brightness temperature ratios differ from unity (Papadopoulos et al.\\ 2000) -- are within the range predicted by our simulation. While the intrinsic CO brightness of J2135 is $\\sim 7\\times$ that of cB58, Fig.\\ \\ref{figure:co-maps}d-f shows that such differences can in principle come about merely from orientation effects. Integrating over the entire line, however, we find that the total CO luminosity of our simulation is relatively constant (within $\\sim 30$ per-cent) with respect to inclination angle, suggesting that CO is a robust tracer of the gas mass. The CO line widths of cB58 and J213 are very similar ($\\Delta V_{\\mbox{\\tiny{FWHM}}}\\simeq 170-190\\,\\mbox{km}\\,\\mbox{s}^{-1}$), and in excellent agreement with the edge-on CO line widths of components A and B. This suggest that both systems are viewed edge-on (J2135 even has hint of a double-horn profile) and that the difference in the CO properties of these two $z\\sim 3$ LBGs is not due to inclination effects, but due to an intrinsic diffference in the gas content of the two systems.      \\subsection{CO as a dynamical mass tracer} Mapping the CO emission from distant galaxies is one of the best ways of determining their dynamical masses, although due to the inability of current interferometers to robustly resolve the CO emission in high-$z$ objects, most attempts at deriving dynamical masses have resulted in little more than upper limits. Here we aim to quantify just how well one can determine the dynamical mass using CO.  For each line of sight we have calculated the dynamical masses of A and B using the disk-formula of Bryant \\& Scoville (1996). Comparing with their total masses ($M^A_{tot}=1.0\\times 10^{10}\\,\\Msolar$ and $M^B_{tot}=0.5\\times 10^{10}\\,\\Msolar$), calculated as the sum of the total baryonic mass, i.e.\\  stars and gas (atomic and molecular), and DM, we find that in the best case scenario when the disks are seen edge-on, the dynamical masses can be inferred from the CO emission to within 20 per-cent. The virial mass of the overall merger-system is given by $M_{vir} = 2RV^2/G$, where $R$ and $V^2$ are the mass-weighted distance between the two systems and the line-of-sight velocity variance. The actual total mass of the entire system ($M_{tot}=2.0\\times 10^{10}\\,\\Msolar$) is found by integrating up all the mass (baryonic and dark) within a sphere centered between the two systems and with a radius corresponding to half their distance. We find that when the two systems are well-separated (Fig.\\ \\ref{figure:co-maps}b-c), the total mass can be inferred to within 20 per-cent, whereas the discrepancy increases to 50 per-cent when the two systems overlap as in Fig.\\ \\ref{figure:co-maps}a.  Observations of CO in distant galaxies have to deal with limited sensitivity and spatial/spectral resolution, all of which can affect dynamical mass estimates. How well can we expect to do such observations using ALMA? To answer this question we simulated a 24\\,hr integration with ALMA of our system (Fig.\\ \\ref{figure:co-maps}g-l). We adopted the projected rms noise of ALMA at $\\sim 108\\,$GHz given by the online ALMA sensitivity calculator with default settings. The noise was added to each 'velocity-slice' in the CO 'data-cube', and convolved with a {\\sc fwhm}~$=0.1\\arcsecs$ Gaussian beam -- a resolution obtainable with ALMA in its extended configuration. From Fig.\\ \\ref{figure:co-maps}g-l it is seen that only in the case where the system is viewed face-on and is at its brightest, do we see a line profile and a detection ($4\\sigma$) in the velocity-averaged maps. Thus if our simulated system is representative of LBGs at $z\\gs 3$, then ALMA observations of CO line emission from such systems will a) depend heavily on the inclination angle under which they are viewed, thereby potentially bias detections towards face-on systems with relatively small apparent CO linewidths, and b) essentially be pure detection experiments since ALMA will not be sensitive enough to constrain CO sizes and line shapes. The implication of a) and b) is that it may prove difficult to infer gas kinematics and dynamical masses of sub-$L^*$ galaxies (such as the one simulated in this paper) at high-$z$ using ALMA."
    },
    "0608/astro-ph0608156_arXiv.txt": {
        "abstract": "{ We present a new numerical algorithm for the solution of coupled collisional and collisionless systems, based on the block structured adaptive mesh and time refinement strategy (AMR). We describe the issues associated with the discretization of the system equations and the synchronization of the numerical solution on the hierarchy of grid levels.  We implement a code based on a higher order, conservative and directionally unsplit Godunov's method for hydrodynamics; a symmetric, time centered modified symplectic scheme for collisionless component; and a multilevel, multigrid relaxation algorithm for the elliptic equation coupling the two components. Numerical results that illustrate the accuracy of the code and the relative merit of various implemented schemes are also presented.} ",
        "introduction": "Astrophysical systems are typically complex and highly nonlinear, providing the ground for the occurrence, either singly or concomitantly, of numerous physical processes. These include, among others, hydro- and magnetohydro-dynamics, gravity, radiation and many-body interactions. They operate on a wide range of spatial and temporal scales and it is often desirable to fully cover these ranges for a thorough understanding of the problem. Thus the problems are demanding both in  terms of physics algorithms and dynamic range (resolution). While the development of high order numerical schemes certainly improves the quality of numerical solutions and the availability of ever more powerful computers has allowed the performance of larger calculations, special techniques are required in order to achieve very large dynamic ranges.  The Adaptive Mesh Refinement (AMR) technique offers a powerful solution for this purpose~\\cite{bergeroliger84,bergercolella89}. While there are difficulties associated with its implementation, the use of AMR in astrophysics and cosmology has grown significantly to include studies of nucleosynthesis in Supernovae explosions~\\cite{timmesetal00,kifonidisetal00,almgrenETAL:2006a,almgrenETAL:2006b}, of the multiphase interstellar medium and radiative shock hydrodynamics~\\cite{kmc94,wafo98}, the problem of star formation out of the collapse of protostellar clouds~\\cite{kleinetal03,trueloveetal97,trueloveetal98} and the formation of the first stars as well as the large scale structure in the universe~\\cite{abn02,norman04b}. The use of AMR technique in the above examples has often been instrumental in either revealing new properties of the investigated system (such as instabilities) or pointing out mistaken views based on limited resolution calculations.  We will consider two aspects of application of AMR to astrophysical problems. The first is the extension to incorporate self-gravity. This capability leads to new algorithmic difficulties due to the elliptic nature of the problem.  In particular the solution has to be computed simultaneously on all levels of refinement and continuity of both the solution and its normal derivative has to be enforced at coarse-fine level interfaces ~\\cite{maca96}.  However, this introduces non-trivial complications when time refinement is also employed: the coarser levels are advanced first and finer levels are advanced with the assumption that boundary conditions at fine-coarse interfaces are provided by the coarse level solutions interpolated in time. However, the full multilevel elliptic solution can only be computed when all levels are synchronized and thus is not available when the coarser levels are ahead of the finer levels.  Thus the first implementations of a full multilevel elliptic solver for self-gravity~\\cite{trueloveetal97,trueloveetal98,rickeretal05} do not use refinement in time. And when employing time refinement the multilevel elliptic equation has been solved as a set of independent boundary value problems, one for each level, not a fully multilevel solution~\\cite{kravtsov99,knebeetal01,teyssier02,quilis04}.  The second issue we will address is the application of AMR to hybrid systems, that is a self-gravitating gas coupled to a particle representation of a collisionless matter described by Vlasov-Poisson equations~\\cite{hoea81} .  This is relevant to several problems in astrophysics, particularly for modeling the formation and evolution of structure in the universe.  In this case the AMR technique is combined with Particle-Mesh methods to compute the right-hand side to the Poisson's equation due to the particles mass.  In order to take advantage of the higher resolution of the finer grids, it is desirable to advance the particles with the force compute on, and the timesteps of, the finest grids that cover their spatial position. This introduces a complication for the elliptic solver similar to the one described above for the case of self-gravitating gas dynamics, in the sense that the particles contribution to the right-hand side of Poisson's equation needs to be accounted for even when the levels of refinement are not synchronized.  Various AMR codes for hybrid systems have been developed over the past several years~\\cite{bryannorman97,kravtsov99,knebeetal01,teyssier02,quilis04,rickeretal05}. The refinement strategy in~\\cite{kravtsov99,teyssier02,rickeretal05} is based on splitting of individual cells and in~\\cite{knebeetal01} only the collisionless component is evolved.  Virtually all schemes use Strang splitting~\\cite{strang68} for the multidimensional version of the hydrodynamics and a modified leap-frog method for the integration of the equation of motion of the particles.  This has also been done for problems in which there is no coupling to a collisional fluid component, as occurs in computations of collisionless plasmas~\\cite{vayETAL:2002}.  In our approach we use the block structured scheme for adaptive mesh and time refinement proposed in~\\cite{bergercolella89} as a starting point and extend it include gravity and collisionless particle dynamics.  We use an unsplit Godunov's method for hydrodynamics \\cite{millercolella02}; a symmetric, time centered modified symplectic scheme based on the kick-drift-kick sequence for the collisionless component; and a multilevel, multigrid relaxation algorithm for the elliptic equation coupling the two systems.  We introduce two new procedures to solve synchronization issues described above that arise with the elliptic solver when the coarse and finer levels are not synchronized. We use a method analogous to those developed for AMR for incompressible flows \\cite{almgrenETAL:1998,martinColella:2000} to compute a lagged estimate of the correction of the elliptic matching conditions at boundaries between refinement levels at times when the levels are not synchronized.  We also present a detailed discussion of refinement in time in the presence of collisionless particles, including methods for associating particles with refinement levels, and a particle aggregation operation to cost-effectively estimate the density distribution on coarse levels due to particles evolved on finer levels without compromising the code accuracy. We provide a detailed description of the formal discretization of the system of equations and the issues involved with the synchronization of the numerical solution in the presence of refinement in time, at a level of detail which we feel is lacking in the current literature.  The paper is organized as follows.  First, the evolution equations and single level algorithms are outlined in Section~\\ref{sleq.sec}. In Section~\\ref{amrgf.sec} we provide a formal definition of the employed AMR volume discretization, variables and operators and then describe in detail the general AMR algorithm for hybrid system. Finally, in section~\\ref{tests.sec} we test the accuracy of the code by comparing its results against for set of standard solutions. ",
        "conclusions": "We have presented a new code based on AMR technique for systems comprising collisional and collisionless components coupled through a long range force. We have thus extended the scheme in~\\cite{bergercolella89} to include collisionless particle dynamics and gravity arising from the mass distribution of the two components. For the hydrodynamics we use a slightly modified directionally unsplit Godunov's method based on Ref.~\\cite{colella90}.  As for the collisionless component we have implemented various time centered modified symplectic schemes based on both the kick-drift-kick and drift-kick-drift sequence. Our implementations of these schemes appear to perform comparably.  We have also used several types of stencil to calculate the force from the potential.  We find that while the staggered schemes appear more accurate when the number (density) of particles is at least as large as the number of grid cells, it produces spurious results when the particles are sparse on the grid. Cell centered stencils thus seem more reliable, especially when a five point cell centered discretization of the Laplacian operator is used.  Due to the time refinement character of the AMR technique the solution on different levels is advanced with different timesteps. Synchronization issues then arise as the multilevel solution to the elliptic equation needs to be solved simultaneously on all levels.  In particular, the density field represented by the particles evolved on finer levels may not be available on coarser levels when they are not synchronized.  Similarly, one cannot account for the effects of the mass distribution on finer levels on the multilevel solution of the potential, unless all levels are synchronized.  Among other features of the code, we have thus introduced an {\\it aggregation} procedure to cost-effectively represent on the coarser levels the particles on finer levels without compromising the code accuracy and performance. We have also implemented a procedure for estimating (when the coarse and finer levels are not synchronized) the effects on the coarse potential produced by the matching conditions at refinement boundaries between coarse and fine solution, as they would arise in a full multilevel calculation.  We performed several standard tests which illustrate the code accuracy as well as the advantages of the AMR technique for the study of both self gravitating hyperbolic systems, collisionless system and hybrid systems.        \\begin{acknowledgment} FM is thankful to D. Serafini, D. Martin, B. van Straalen and D. Graves for useful discussions,  to the Lawrence Berkeley National Laboratory, for its hospitality, to the Institute of Informatics, ETH Z\\\"urich, for technical support. FM also acknowledges partial support by the European Community through contract HPRN-CT2000-00126 RG29185 and by the Swiss Institute of Technology through a Zwicky Prize Fellowship. \\end{acknowledgment}"
    },
    "0608/astro-ph0608226_arXiv.txt": {
        "abstract": "% {The X-ray source 2S0114+650$=$LSI$+$65 010 is a  binary system containing a B-type primary and a low mass companion believed to be a neutron star.  The system has  three reported periodicities:  the orbital period, P$_{orb}\\sim$11.6 days,   X-ray flaring with P$_{flare}\\sim$2.7 hours  and a ``superorbital\"  X-ray periodicity P$_{super}=$30.7 days.} {The objective of this paper is to show that the puzzling periodicities in the system may be explained in the context of  scenarios in which tidal interactions drive oscillations in the B-supergiant star. } {We   calculate the solution of the equations of motion for one layer of small surface elements distributed along the equator of the star, as they respond to the forces due to gas pressure, centrifugal, coriolis, viscous forces, and the gravitational forces of both stars, which provides variability timescales that can be compared with those observed for 2S0114$+$650.  In addition, we use  observational data obtained at the  Observatorio Astron\\'omico Nacional en San Pedro M\\'artir (OAN/SPM) between 1993-2004 to determine which  periodicities may be present in the optical region. } {The models for circular orbits predict ``superorbital\" periods while the eccentric orbit models predict strong variations on orbital timescales, associated with periastron passage.  Both also predict oscillations on timescales of $\\sim$2 hrs. We suggest that the tidal oscillations lead to a structured stellar wind which, when fed to the neutron star, produces the X-ray modulations. The connection between the stellar oscillations and the modulation of the mass ejection may lie in the shear energy dissipation generated by the tangential motions that are produced by the tidal effects, particularly in the tidal bulge region. From an observational standpoint, we find indications for variability in the He I 5875 \\AA\\ line on $\\sim$2 hrs timescale and, possibly, the ``superorbital\" timescale.  However, the line profile variability exceeds that which is predicted by the tidal interaction model and can be understood in terms of variable emission that is superposed on the  photospheric absorption.  This emission appears to be  associated with the B-supergiant's stellar wind rather than the vicinity of the companion. } {The model calculations lead us to conclude that the B-supergiant may be the origin of the periodicities observed in the X-ray data,  through a combination of a localized structured wind that is fed to the collapsed object  and, possibly, by production of X-ray emission on its own surface. This  scenario  weakens the case for 2S0114+650 containing a magnetar descendent.} ",
        "introduction": "The X-ray source 2S0114+650 (Dower \\& Kelley 1977) is a  binary system containing a luminous B-type star and a low mass companion believed to be a neutron star.  Unlike many such systems, however, the short-period X-ray pulsations associated with a rapidly rotating neutron star are apparently absent.  Koenigsberger et al. (1983) reported $\\sim$15 minute pulsations which, however, were not  found in later observations except those of Ginga (Yamauchi et al. 1990). On the other hand, X-ray flaring activity was shown to be periodic (Finley et al. 1992;  Hall et al. 2000) with P$_{flare}\\sim$2.78 hrs, a period that in more recent INTEGRAL data has decreased to $\\sim$2.67 hrs (Bonning \\& Falanga 2005). Finley et al. (1992) suggested that this period could actually be due to the pulsar's rotation which, however, would make it an extremely slow pulsar.  Li \\& van den Heuvel (1999) showed that the  neutron star in the 2S0114+650 system could  have reduced its rotation rate within the lifetime of the B-supergiant companion only if it had a very large ($>10^{14}$ Gauss) magnetic field at the time of its birth. Thus, the new-born pulsar would have been a {\\it magnetar}.  An alternative scenario for explaining the  2.7 hour periodicity that was also suggested by Finley et al. involves  pulsations of the optical component by which a structured stellar wind with alternating high and low density regions is produced. The   density variations of material being accreted by the collapsed companion would lead to the observed periodic flaring activity.  However,  it is not evident: a) whether a B-supergiant may oscillate with this short period;  nor b) how the stellar oscillations  can produce the variable mass-loss;  nor c) how the ejected mass can be  distributed in the precisely structured stellar wind that is needed for the periodic modulation of the accretion rate onto the neutron star. If a plausible scenario involving pulsations of the B-star could be constructed, the {\\it magnetar} hypothesis would be weakened.  A second  X-ray variability timescale in 2S0114+650 is the orbital period.  Corbet et al. (1999) suggested that this variability may be due to eclipses of the X-ray source,  but Mukherjee \\& Paul (2006) find that the change in the ASCA X-ray luminosity can be ascribed to a local change in the density of the stellar wind from which the neutron star accretes.  The latter would occur  either if the orbit is eccentric  or if the mass-loss characteristics of the B-star are orbital-phase dependent.  Interestingly, the description given by  den Hartog et al. (2006) of the broad maximum in the RXTE-ASM and INTEGRAL data as consisting mostly of flares supports the idea of a non-steady mass-loss rate from the B-star, with larger amplitude variability at phases when the average X-ray flux is maximum.  This behavior is one that might be expected in eccentric orbits.  On the other hand, however, a third periodic X-ray variability timescale has been recently reported by Farrell et al. (2006): P$_{super}\\sim$30.7 days, which they suggest may be caused by a warped accretion disk around the collapsed object.  However, the absence of He II 4686 \\AA\\ emission (Reig et al. 1996; Koenigsberger et al. 2003) argues against the presence of a significant accretion disk.   Non-synchronously rotating stars in binary systems are well-known for exhibiting numerous periodicities due to the oscillations that are driven by the tidal forces (Press \\& Teukolsky 1977; Zahn 1977; Savonije et al. 1995; Kumar et al. 1995). In particular, Moreno et al. (2005) found ``superorbital\" periodicities in models of binary systems with circular orbits, in addition to shorter period variability. Thus, the question naturally arises of whether   we can produce a consistent picture for explaining  the puzzling periodicities of 2S0114+650 through scenarios involving tidal interactions.  In Section 2 we present the observational data;  in Section 3 we summarize the tidal interaction model that we use and  describe its predictions for an eccentric and a circular orbit;  in Section 4  we describe  the behavior of the He I 5875 \\AA\\ line; and  Section 5 contains  our conclusions.    \\begin{table*} \\caption{Radial velocity and residual emission He I 5875 \\AA} \\label{table:1} \\centering \\begin{tabular}{lllllllrlr} \\hline\\hline Epoch& Instrum.& $<MJD>$ & N & S/N & $\\phi_{orb}$& $\\phi_{sup}$ & s.d.(ISM)& RV ($\\pm$s.d.)& $-$EW  \\\\ &         &       & &   &     &        & km/s     & km/s   &    \\AA        \\\\ \\hline 1993Oct& REOSC   & 49269.54& 5  & 26(8)  &  0.804&0.78  & 1.1 & $-$41.5 (1.8) & 0.35\\\\ & --- & 49270.33& 5  & 29(7)  &  0.881&0.80   & 1.5  &$-$37.7 (2.3)& 0.39  \\\\ & --- & 49273.39& 14 & 20(7)  &  0.145&0.90   & 2.6 & $-$41.7 (3.9)& 0.45 \\\\ 1995Nov& UH coud\\'e  & 50049.38& 13 &36(5)   &  0.052&0.11   & 0.7 &$-$65.8 (4.6) & 0.61 \\\\ & --- & 50050.34&  9 &33(9)   &  0.135&0.14   &  0.4 & $-$63.6 (3.5)&0.65   \\\\ 2001Jan& REOSC   & 51931.18& 11 &84(17)  &  0.304&0.22   & 2.1& $-$53.9 (3.8) & 0.00 \\\\ 2001Oct& REOSC   & 52195.36& 12 &68(11)  &  0.082&0.80   & 3.4 & $-$36.9 (4.2)& 1.10  \\\\ 2003Jan& MES     & 52657.40& 12 & 38(5)  &  0.920&0.81   &1.3 & $-$50.3 (3.3) & 1.09 \\\\ 2004Oct& MES     & 53307.33& 2  & 60     & 0.957&0.92   & ---  & $-$23.3 & 0.60       \\\\ & --- & 53308.30& 2  & 50     & 0.041 &0.95  & --- & $-$36.5 &  0.99      \\\\ & --- & 53309.27& 1  & 53     & 0.125 &0.98  & ---   & $-$44.0 & 0.97       \\\\ 2004Nov& REOSC   & 53311.34& 15 & 60(11) & 0.304 &0.05  &  2.4 & $-$62.3 (5.2)& 0.70  \\\\ & --- & 53314.30& 3  & 55     & 0.559 &0.14  & ---    &$-$87.6 &  0.87     \\\\ & --- & 53315.20& 2  & 62     & 0.637 &0.17  & ---    &$-$77.4 &  0.59     \\\\ & --- & 53316.30& 2  & 47     & 0.731 &0.21  & ---    &$-$75.0 &  0.81     \\\\ 2004Nov& MES     & 53328.16& 1  & 69     & 0.754&0.59    & ---    & $-$47.5& 0.75         \\\\ & --- & 53329.43& 1  & 74     & 0.864 &0.64  & ---    & $-$42.3& 0.66        \\\\ & --- & 53330.29& 1  & 53     & 0.938 &0.66  & ---    & $-$38.8& 0.57        \\\\ &  ---& 53333.40& 2  & 70     & 0.206 &0.77  & ---    & $-$63.2& 0.62        \\\\ & --- & 53334.41& 1  & 60     & 0.293 &0.80  & ---    & $-$70.7& 0.21        \\\\ & --- & 53335.40& 1  & 66     & 0.378 &0.83  & ---    & $-$80.2& 0.49        \\\\ 2004Dec& MES     & 53354.23& 16 & 58(9)  & 0.002 &0.44  & 1.5 & $-$30.1 (1.3) & 0.84    \\\\ & --- & 53355.20& 17 & 53(6)  & 0.085 &0.47  & 1.0 & $-$43.6 (1.5)& 0.90  \\\\ & --- & 53356.20& 15 & 53(13) & 0.172 &0.51  & 1.1 & $-$46.0 (1.5)& 0.77  \\\\ & --- & 53357.20& 18 & 54(8)  & 0.258 &0.54  & 1.1 & $-$59.1 (2.1)& 0.95  \\\\ & --- & 53358.20&  8 & 59(9)  & 0.344 &0.57  & 0.3 & $-$60.8 (1.6)& 1.15  \\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "Tidal interactions in non-synchronously rotating binary systems  produce strong perturbations on the external layers of the stars, particularly when the orbital separations are relatively small and the stellar radii are large. The binary system LSI$+$65 010=2S0114$+$650 is precisely such a system.   The exploratory calculations performed for this paper are aimed at establishing the plausibility that the X-ray variability patterns of LSI+65 010$=$2S0114+650 may be attributable to these tidal interactions.  We show that  the  oscillations on the surface of the B-supergiant that are predicted in our calculations have timescales that are consistent with the periodic timescales found in X-ray data.  We propose that these oscillations may modulate the B-star's wind, leading to a fluctuating accretion rate onto the neutron star, and hence, X-ray variability. Figure 14 is a cartoon depicting the  scenario for oscillations with P$\\sim$2-3 hours which are assumed to produce enhanced-density shells.  This idea, originally proposed by Finley et al. (1992) was discarded when the optical component of 2S0114$+$650 was found to be a B-supergiant instead of a $\\beta$ Cep-type star, since such short timescale pulsations are not predicted for single supergiant stars.   We rescue Finley et al.'s hypothesis by proposing that the presence of a companion can lead to the appearance of such short periods.  The precise mechanisms by which the  oscillations may translate into a structured stellar wind are not clear, but we speculate that an answer may lie in the combination of the radial pulsation components and the shear energy dissipation due to the tangential motions of the outer stellar layers sliding over the  more slowly rotating inner layers.  These two combined sources (the radial pulsation component and the shear energy dissipation), when added to the radiation field that normally drives the stellar wind, might cause the inhomogeneities to appear on the same timescale as that of the largest amplitude oscillations. In order to establish a connection between the tidal oscillations and the modulation of  mass-ejection in LSI$+$65 010, it is necessary to analyze spectral lines that are less susceptible to contamination from emission, together with observations of H$\\alpha$.  It is important to note that such observations have been successfully performed in other stars (Fullerton, Gies \\& Bolton 1996; Prinja et al. 2004; Kaufer et al. 2006) showing a link between photospheric activity  and perturbations in the dense inner regions of stellar winds.  It is interesting to note that the largest oscillation  amplitudes  induced by the tidal forces occur near the primary tidal bulge.  Hence, the periodic modulation of the stellar wind may only be present in a small solid angle oriented in the general direction of the neutron star. The localized nature of the   oscillations would explain the lack of a clear $\\sim$2.7 hr periodicity in the  He I 5875 \\AA\\  data, since the irregular variations that are present over the whole B-star surface most likely drown  out the small-scale periodic effects.  Thus, an interesting conclusion is that the X-ray emission is a much finer diagnostic of the local wind structure than the optical line profiles, the latter reflecting the integral of the absorption/emission over the stellar disk.  The tidal interaction model predicts a large difference between the pulsation energy dissipation rates, $\\dot{E}$, near  periastron and apastron for an eccentric orbit.  This implies a significant variation in surface activity as a function of orbital phase, assuming that the surface activity is associated with $\\dot{E}$.  X-ray  variability on the orbital timescale is indeed present, but the line profiles in our optical data do not display strong changes between periastron and apastron phases, weakening the case for an eccentric orbit. In a circular orbit, $\\dot{E}$ remains constant over orbital phase and thus no modulation of the stellar wind with P$_{orb}$  is expected.  If the orbit of LSI$+$65 010 is circular, the observed X-ray variability is likely due to eclipse effects, rather than to differences in the accretion rates onto the companion.  A circular orbit would also be consistent with the existence of a ``superorbital\" period.  Our final consideration concerns the possibility that part of the X-ray emission may actually originate in the B-supergiant.   In particular, the tangential component of the perturbation produces a differentially rotating  structure in the surface layers that one might speculate could lead to magnetic field generation.  Admittedly, the timescales over which a particular region of the star develops this structure  may be too short for magnetic field amplification mechanisms for massive stars, such as proposed by Spruit (2002), to operate. However, it would be interesting to  explore the  mechanisms that could convert tidal energy into magnetic field energy, as has been done  for low-mass stars (for example, Zaqarashvili et al. 2002).  Furthermore, observational evidence for  surface magnetic field structures in OB stars now exists (Henrichs, Neiner, Geers 2003), so it would be interesting also to explore the effects of the tidal oscillations on assumed pre-existing  magnetic field structures and determine the X-ray emitting characteristics due to these perturbations."
    },
    "0608/hep-th0608207_arXiv.txt": {
        "abstract": "{In this paper, we investigate the noncommutative KKLMMT D3/$\\overline{\\rm D}$3 brane inflation scenario in detail. Incorporation of the brane inflation scenario and the noncommutative inflation scenario can nicely explain the large negative running of the spectral index as indicated by WMAP three-year data and can significantly release the fine-tuning for the parameter $\\beta$. Using the WMAP three year results (blue-tilted spectral index with large negative running), we explore the parameter space and give the constraints and predictions for the inflationary parameters and cosmological observables in this scenario. We show that this scenario predicts a quite large tensor/scalar ratio and what is more, a too large cosmic string tension (assuming that the string coupling $g_s$ is in its likely range from 0.1 to 1) to be compatible with the present observational bound. A more detailed analysis reveals that this model has some inconsistencies according to the fit to WMAP three year results.}       \\begin{document} ",
        "introduction": "The inflation paradigm provides a compelling explanation for the homogeneity, isotropy and flatness of the universe by positing an early epoch of accelerated expansion \\cite{Guth}. This accelerated period of expansion also predicts an almost scale-invariant density perturbation power spectrum, which has received strong observational support from the measurement of the temperature fluctuation in the cosmic microwave background radiation \\cite{cobe,wmap1}. It is expected for inflation that it should emerge naturally from some fundamental theory of microphysics. String theory is the most promising candidate hitherto for a fundamental theory, so it is rather natural to look for an explicit realization of the inflationary scenario within the framework of the superstring theory. An attractive scenario for inflation based upon string theory is the brane-antibrane inflation \\cite{Dvali:1998pa,Burgess:2001fx} which proposes that inflation might arise from the interaction potential between a D3-brane and an anti-D3-brane which are parallel and widely separated in a six dimensional compact manifold where the six compactified dimensions are dynamical stabilized \\cite{Giddings:2001yu,kklt}. The anti-D3-brane is held fixed at the bottom of a warped throat (due to an attractive force) while the D3-brane is mobile. The D3-brane experiences a small attractive force towards the anti-D3-brane. The distance between the branes plays the role of the inflaton field. This is the KKLMMT scenario \\cite{kklmmt}. Herein, inflation ends when the D3-brane and the anti-D3-brane collide and annihilate, initiating the hot big bang epoch. The annihilation of the D3-$\\overline{\\rm D}$3-branes allows the universe to settle down to the string vacuum state that describes our universe. During the process of the brane collision, cosmic strings are copiously produced \\cite{Jones:2002cv,Sarangi:2002yt}. For a well-worked inflationary scenario, the production of topological defects other than cosmic strings must be suppressed by many orders of magnitude. It should be notable that the D3/$\\overline{\\rm D}$3 brane inflation scenario achieves this property automatically.   For the D3/$\\overline{\\rm D}$3 brane inflation scenario, the potential that arises from the attraction between the brane and the antibrane in the throat generically takes the following form \\cite{Firouzjahi:2005dh} \\begin{equation} V(\\phi)={1\\over 2}\\beta H^2\\phi^2+2T_3h^4\\left (1-{M^4\\over\\phi^4}\\right), \\label{potential} \\end{equation} where \\begin{equation} M^4={27\\over 32\\pi^2}T_3h^4.\\label{m4} \\end{equation} Here, $T_3$ is the tension of the D3-brane, $h$ is the warp factor of the throat, $H$ is the Hubble parameter, and $\\phi$ is the canonical inflaton ($\\phi=\\sqrt{T_3}r$, where $r$ is the position of the D3-brane with respect to the bottom of the throat). In equation (\\ref{potential}), the first term ($\\beta$ term, where $\\beta$ characterizes the mass of the inflaton field) comes from the contributions of the K\\\"ahler potential and various interactions in the superpotential \\cite{kklmmt} as well as possible D-terms \\cite{Burgess:2003ic}, the second term ($T_3h^4$ term) is the effective cosmological constant coming from the presence of the $\\overline{\\rm D}$3-brane (with tension $T_3$) sitting at the bottom of the throat, and the last term ($\\phi^{-4}$ term) is the Coulombic-like attractive potential between the D3-brane and the $\\overline{\\rm D}$3-brane \\cite{kklmmt}. The parameter $\\beta$ is generically of order unity. However, to achieve slow roll, $\\beta$ has to be fine-tuned to be much smaller than 1. This case corresponds to the KKLMMT scenario, where the effect of the $\\beta$ term is negligible. Effort has been made for relaxing the fine-tuning of $\\beta$ \\cite{Firouzjahi:2005dh}. As is shown in reference \\cite{Firouzjahi:2005dh} that the fine-tuning of $\\beta$ can be relaxed to $\\beta\\lesssim 1/7$ when admitting $n_s\\leq 1.086$ (2 $\\sigma$ result) which comes from an analysis of the first year WMAP data combined with SDSS Lyman $\\alpha$ and other data \\cite{Seljak:2004xh}. What is remarkable is that slow roll inflation still works fine in this case, while the cosmic string tension increases rapidly. In this very simple generalization of the KKLMMT scenario, the cosmic string tension can easily saturate the observational bound. As the fine-tuning is relaxed, the primordial density perturbation spectrum moves from red-tilted to blue-tilted. However, the recent announced WMAP three-year data explicitly indicate that the scalar spectral index of the primordial fluctuations is red-tilted when assuming a power law spectrum, namely $n_s=0.951^{+0.015}_{-0.019}$ \\cite{wmap3}. This fact places significant and stringent constraints on the D3/$\\overline{\\rm D}$3 brane inflation scenario \\cite{Huang:2006ra}. Here we show in figure \\ref{fig:nbpl} the constraints on the brane inflation model from the WMAP result of the spectral index of the primordial density perturbation. Note that $\\beta<0$ is forbidden since in this case $\\beta$ term in the potential tends to push D3-brane out of the throat such that the inflation will not happen. However, we see clearly in this figure that according to the WMAP result $\\beta$ is very likely to be less than zero. For example, $n_s=0.951$ gives $\\beta=-0.007$ at $N=54$, where $N$ is the $e$-folds before the end of the inflation. It should be pointed out that a reasonable range for $e$-folds is $47<N<61$ (or $N=54\\pm 7$) \\cite{Alabidi:2005qi}. Thus the area of $\\beta>0$ and $47<N<61$ seems a rational region for the brane inflation. The curve for $n_s=0.966$ (1 $\\sigma$) has a very small part locating in this region as shown in figure \\ref{fig:nbpl}. For $n_s=0.966$, taking $N=47$ one derives $\\beta=6.2\\times 10^{-4}$, this is a fine-tuning. The curve for 2 $\\sigma$ result $n_s=0.981$ (note that we take 2 $\\sigma$ as twice of 1 $\\sigma$ here for simplicity) resides in the legal region. Here for this case we take $N=54$ and then get $\\beta=5.4\\times 10^{-3}$ which is also a fine-tuning value. On the whole, even though $\\beta>0$ can be achieved constrainedly, the fine-tuning of $\\beta$ is inevitable provided that the power spectrum of the primordial perturbation is red-tilted.  \\begin{figure}[htbp] \\begin{center} \\includegraphics[scale=1.2]{nbpl.eps} \\caption[]{\\small The constraint relationship between the parameter $\\beta$ and the $e$-folding number $N$ before the end of inflation, in the usual (i.e. commutative) D3-${\\rm \\overline{D}3}$ brane inflation, according to the WMAP results of the spectral index of the primordial density perturbation (within 2 standard deviations) which comes from assuming a power law spectrum. Note that the area of $\\beta>0$ and $47<N<61$ is a rational region for the brane inflation.}\\label{fig:nbpl} \\end{center} \\end{figure}   On the other hand, for the power spectrum of primordial fluctuations there exists the possibility that the spectral index varies slowly with scale. When allowing for a running spectral index, the WMAP three-year data favor a blue-tilted spectral index with a large negative running (and also a large tensor amplitude, characterized by $r$, the ratio of the tensor to scalar power spectrum) \\cite{wmap3}, \\begin{equation} n_s=1.21^{+0.13}_{-0.16},~~~~\\alpha_s=dn_s/d\\ln k=-0.102^{+0.050}_{-0.043},~~~~r\\leqslant 1.5~~~{\\rm at~95\\%~CL}.\\label{wmap3} \\end{equation} However, the D3-$\\overline{\\rm D}$3 brane inflation scenario can not produce a so large absolute value of the running spectral index (the typical value it attains is $|\\alpha_s|\\sim 10^{-4}$) though the blue-tilted spectrum can release the fine-tuning of $\\beta$ effectively. Actually, it is rather hard for most inflationary models to produce such a large running of the spectral index within the framework of the usual slow-roll mechanism \\cite{hardrun}. However, it has been shown that when considering the space-time noncommutative effects in the very early universe due to the possible effects of quantum gravity, the inflation models can nicely produce a large running spectral index \\cite{noncomu,ncnew}.\\footnote{For extensive studies on the noncommutative inflation models see \\cite{ncold}, and for other discussions on the running spectral index see \\cite{run1,run2}.} Recently, a noncommutative brane inflation scenario has been proposed to accommodate the large negative running of the spectral index in the WMAP three-year results and to release the fine-tuning of $\\beta$ due to the blue-tilted spectrum \\cite{Huang:2006zu}. In this paper, we will investigate this noncommutative brane inflation model in detail. We shall show the ranges of the mass parameter $\\beta$ as well as the noncommutative parameter $\\mu$ allowed by WMAP three-year results of the spectral index and its running. Within the allowed ranges of $\\beta$ and $\\mu$, we evaluate the tensor-to-scalar ratio and the cosmic string tension. We shall show that such a noncommutative brane inflation will predict a rather large tensor-to-scalar ratio (implying a large gravity wave amplitude) and what is more, a too large cosmic string tension (assuming the value of the string coupling $g_s$ is in the range 0.1 to 1) to be compatible with the present observational results of the anisotropy of the cosmic microwave background radiation. Furthermore, we analyze the various parameters and physical quantities/cosmological observables in this noncommutative brane inflation model by fitting to the three-year WMAP results and show that some inconsistencies occur in this model. We shall discuss about these inconsistencies and give some comments. ",
        "conclusions": "In this paper, we have investigated the noncommutative D3/$\\overline{\\rm D}$3 brane inflation scenario in detail. The usual D3/$\\overline{\\rm D}$3 brane inflation (KKLMMT) model suffers from the fine-tuning problem, especially after the announcement of the WMAP three-year data, the fine-tuning of $\\beta$ becomes more severe (due to the red-tilted spectrum), and only a narrow range of parameter space for this model is still allowed. On the other hand, when allowing for a running spectral index, the WMAP three-year data favor a blue-tilted spectral index with a large negative running. However, the KKLMMT scenario can not explain the running spectral index unless accepting a noncommutative generalization. The noncommutative brane inflation model proposed in reference \\cite{Huang:2006zu} can not only nicely explain the large negative running of spectral index indicated by WMAP data but also effectively eliminate the fine-tuning problem in the KKLMMT model. Though having these advantages, nevertheless, such a noncommutative KKLMMT model predicts an unacceptable large cosmic string tension (if assuming $g_s=1$) beyond the observational bound by two orders of magnitude. A more detailed analysis indicates that there are some inconsistencies in this model. The most severe problem is that the value of $\\beta(\\phi_N/M_{\\rm Pl})^2$ derived by fitting to WMAP three year results is too big to insure the validity of some approximations in the model and furthermore, the occurrence of $\\beta(\\phi_N/M_{\\rm Pl})^2>6$ breaks down the Friedmann equation. This implies that the way of generalizing the brane inflation to the noncommutative space-time in this model is not delicate enough. Still, the idea of combining the brane inflation scenario with the noncommutative inflation framework is interesting and worth advocating. Hence, it is necessary to find a graceful way of incorporating these two scenarios and construct a self-consistent model of noncommutative brane inflation."
    },
    "0608/astro-ph0608295_arXiv.txt": {
        "abstract": "\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}}  We investigate the gas content and baryonic Tully-Fisher relationship for extremely low luminosity dwarf galaxies in the absolute magnitude range $-13.5>M_r-5\\log_{10} h_{70} >-16$.  The sample is selected from the Sloan Digital Sky Survey and consists of \\ngal\\ galaxies for which we have obtained follow-up HI observations using the Arecibo Observatory and Green Bank Telescope.  This represents the largest homogeneous sample of dwarf galaxies at low luminosities with well-measured HI and optical properties.  The sample spans a range of environments, from dense groups to truly isolated galaxies. The average neutral gas fraction is $\\langle f_{\\rm gas}\\rangle=$\\,\\fgas, significantly exceeding that of typical gas-rich galaxies at higher luminosities.  Dwarf galaxies are therefore less efficient at turning gas into stars over their lifetimes.  The strong environmental dependence of the gas fraction distribution demonstrates that while internal processes can reduce the gas fractions to roughly $f_{\\rm gas} = 0.4$, external processes are required to fully remove gas from a dwarf galaxy.  The average rotational velocity of our sample is $\\langle v_{\\rm rot}\\rangle=\\langle$W$20_{i,t}/2\\rangle=50$ \\kms\\ based on HI line-widths.  In this luminosity range, the optical Tully-Fisher relationship has significantly more scatter compared to the baryonic relationship.  Including more massive galaxies from the literature, we fit a baryonic Tully-Fisher slope of $M_{\\rm baryon} \\propto v_{\\rm rot}^{3.70\\pm 0.15}$.  This slope compares well with CDM models that assume an equal baryon to dark matter ratio at all masses.  While gas stripping or other processes may modify the baryon to dark matter ratio for dwarfs in the densest environments, the majority of dwarf galaxies in our sample have not preferentially lost significant baryonic mass relative to more massive galaxies. ",
        "introduction": "\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}}  The properties of low mass dwarf galaxies, and their differences from luminous galaxies, provide important clues to understanding both cosmology and galaxy formation.  As probes of galaxy formation, dwarf galaxies should be extremely sensitive to processes such as star formation feedback, ram pressure stripping, and tidal stripping. These processes are often invoked to explain properties of relatively luminous galaxies and should even more dramatically affect dwarf galaxies, which have smaller masses and shallower potential wells.  As probes of cosmology, the number of dwarf galaxies has provided a challenge for cosmological models.  The current Cold Dark Matter model with a cosmological constant ($\\Lambda$CDM; \\citealt{spergel06a}) makes robust predictions for the number of dark matter halos as a function of mass \\citep{sheth01a}. However, extending this model to predict the number density of observed galaxies requires including physical effects such as gas cooling, star-formation, supernova feedback, etc., that are too complex to follow exactly in numerical predictions.  It is clear that the number of galaxies as a function of luminosity is much shallower than the prediction for the number as a function of halo mass \\citep{benson03a}. The most extreme example of this problem is the ``substructure problem'' in the Local Group \\citep{klypin99a,moore99a}.  If we assume galaxies are associated one-to-one with dark matter halos and that $\\Lambda$CDM is correct, the mass-to-light ratios of galaxies therefore must decline as a function of host halo mass.  That is, the efficiency of forming stars and/or retaining gas must be much smaller for low mass galaxies.  A number of physical models have been proposed to explain high mass-to-light ratios in low mass galaxies, including photoevaporation of gas during reionization \\citep{shapiro04a,mayer04a}, suppression of gas accretion during reionization \\citep{gnedin00a,bullock00a,benson03a}, feedback from supernovae \\citep{dekel03a,nagashima04a}, and ram pressure stripping of gas \\citep{mori00a}.  Some of these processes only occur for dwarf galaxies near a luminous galaxy and not for isolated ones \\citep{kravtsov04a}.  The relative importance of these processes in the evolution of dwarf galaxies has significant implications for galaxies of all sizes at both low and high redshift.  For example, the importance of galactic winds in low mass galaxies will strongly affect our understanding of the mass-metallicity relation \\citep{tremonti04a,lee06a} and the enrichment of the intergalactic medium at high redshift \\citep{scannapieco05a}.  By comparing the gas content and Tully-Fisher relationships of dwarfs as a function of environment, it should be possible to determine the importance of ``external'' environmental effects like stripping relative to ``internal'' effects such as feedback.  Historically, extremely low luminosity nearby galaxies have been difficult to find because their colors and apparent sizes are similar to those of more luminous, and far more numerous, background galaxies. By far the best known set of dwarfs is that in the Local Group \\citep{mateo98a,grebel03a,dolphin05a}, where the even the lowest mass dwarfs can be detected as over-densities in stellar counts \\citep{willman05a, zucker04a}.  The dwarf galaxy population has also been well studied in nearby clusters, such as Virgo and Fornax, where the number of dwarfs is large relative the background population \\citep{binggeli85b,ferguson91a} or in other environments where dwarfs are expected, eg.~as satellites to bright galaxies \\citep{zaritsky97a,gutierrez02a}.  Assembling samples of dwarfs in more distant and less dense environments has been challenging.  Using criteria such as surface brightness and irregularity, several authors have assembled large catalogs of nearby dwarf galaxies \\citep{impey96a, schombert01a, hunter04a,karachentsev04a}.  Addressing the above issues requires a large sample of galaxies across a range of environments, with minimal selection effects.  Deep, wide-field optical spectroscopic surveys, such as the Sloan Digital Sky Survey (SDSS), provide a large, homogeneous sample of dwarf galaxies covering a wide range of galactic properties and environments.  Concurrently, blind HI surveys such as the HI Parkes All Sky Survey (HIPASS; \\citealt{meyer04a}) and the upcoming Arecibo Legacy Fast ALFA Survey (ALFALFA; \\citealt{giovanelli05a}) will provide similarly large samples, with a very complementary set of selection effects.   Unlike present-day massive disk galaxies, which have used up much of their gas reservoirs \\citep{read05a}, the gaseous component of dwarf galaxies cannot be ignored \\citep{schombert01a,vanzee01a}.  Indeed, we show that at the lowest luminosities dwarf galaxies are {\\it mostly} neutral gas.  A galaxy's gas content reflects the efficiency of star formation modified by processes which add or removal gas from the system such as gas infall or ram pressure stripping.  The goal of this paper is to combine the optical properties, dynamics and neutral gas content of low mass dwarf galaxies in order to better understand the processes important in dwarf galaxy evolution.   We present neutral hydrogen (HI) observations for \\ngal\\ low luminosity dwarf galaxies selected with uniform and well-understood criteria from the SDSS.  We describe the catalog from which these dwarfs were selected in \\S\\,\\ref{ssec_sample}, estimate the environment of each object in \\S\\,\\ref{ssec_environ} and discuss the HI observations in \\S\\,\\ref{ssec_redux_HI}.  We investigate correlations between the gas fraction, color, star formation rate and environment of this sample in \\S\\,\\ref{ssec_fgas} and the Tully-Fisher relationship in \\S\\,\\ref{ssec_tf}.  When necessary, we assume a Hubble constant of H$_{0} = 70 h_{70}$\\,Mpc \\kms\\ where $h_{70} = 1$. ",
        "conclusions": "We present results from an HI survey of low luminosity dwarf galaxies in the absolute magnitude range $-13.5 > M_r - 5\\log_{10} h_{70} > -16$ selected from the SDSS.  These data represent the largest homogeneous sample of low luminosity dwarf galaxies with well-measured HI and optical properties.  The sample spans a range of total baryonic masses between M$_{\\rm baryonic} = 10^8-10^9$\\Msun with measured HI line-widths between $\\langle v_{\\rm rot}\\rangle=\\langle$W$20_{i,t}/2\\rangle=20 - 80$ \\kms.  The dwarf galaxies in this sample are found in a wide range of environments, from dense groups to truly isolated galaxies.  For the majority of our dwarf galaxy sample, the total galactic baryonic mass is dominated by the gas mass rather than stellar mass. The average gas fraction for our sample is $\\langle f_{\\rm gas}\\rangle=$\\,\\fgas, with gas fractions as high as \\maxfgas\\%. This significantly exceeds that of gas-rich galaxies at higher luminosities.  Thus, dwarf galaxies are inefficient at turning gas into stars compared to more luminous galaxies.  We find that gas-rich galaxies have bluer $g-r$ colors compared to gas-poor objects implying that gas-rich dwarf galaxies have had more recent star formation.  The majority of galaxies in our sample are currently forming stars based on observed H$\\alpha$ emission in the SDSS spectra.  We find that the {\\it current} rate of star formation per unit area (based on the total $u$-band flux) is not correlated with gas fraction, confirming previous results suggesting that the star formation rates in dwarf galaxies are stochastic \\citep{vanzee01a,hunter04a}.  Thus, dwarf galaxies form stars with a rate that varies over million year time scales, but is correlated with the gas fraction over longer time periods.   Using the range of environments within our sample, we examine whether internal or external processes control the gas fractions of dwarf galaxies.  Trends between galaxy properties, such as color or morphology, and environment have been well studied for galaxies at brighter luminosities \\citep{norberg02a,blanton05b}, but have been previously limited to the Local Group and nearby groups at dwarf luminosities \\citep{grebel03a}.  Our sample allows us to quantify these trends for dwarf galaxies over a much wider range of environments.  We find that dwarf galaxies with gas fractions $f_{\\rm gas} < 0.4$, comprising $15\\%$ of this sample, are exclusively found within 0.5\\,Mpc of a luminous neighbor galaxy.  Galaxies with gas fractions $f_{\\rm gas} > 0.4$ are found across the full range of environments contained within this sample.  Thus, while internal galaxies processes can reduce the gas fractions, external processes are required to fully remove gas from a dwarf galaxy.   The optical Tully-Fisher relationship within our dwarf galaxy sample demonstrates considerable scatter.  As previous work has shown for brighter samples (\\citealt{bell01b, verheijen01b, gurovich04a, mcgaugh05a}), the baryonic Tully-Fisher relationship within our sample is much tighter. When we consider the Tully--Fisher relationship across all absolute magnitudes, we find a logarithmic slope of $b = 0.27\\pm 0.01$ (or $1/b=3.70 \\pm 0.15$). These results are in accord with CDM predictions of $1/b \\sim 3.5$ assuming no preferential loss of baryons as a function of mass (\\citealt{bullock01a}). Thus, our results suggest that processes which preferentially remove gas from dwarf galaxies are not important for the majority of galaxies in our sample with maximum rotation velocities $v_{\\rm max} \\sim 20 - 80$ \\kms\\ (roughly corresponding to total dynamical masses of M$_{\\rm total}=10^9-10^{10}$\\Msun).   If galaxies are associated one-to-one with $\\Lambda$CDM dark matter halos, the mass-to-light ratios of galaxies must decline as a function of host halo mass \\citep{benson03a}.  The most well-known example of this problem is the Local Group ``substructure problem'' \\citep{klypin99a,moore99a}.  To explain this trend in mass-to-light ratio, investigators often invoke processes such as ram pressure stripping or supernova feedback that either removes gas or heats it sufficiently to prevent star-formation in low mass galaxies (e.g., \\citealt{dekel86a,dekel03a,mori00a,bullock00a}).  Most semi-analytic models require significant baryonic mass loss \\citep{cole00a,benson03a,croton06a}, predicting baryon-to-dark matter ratios for dwarf galaxies well below the ``Universal'' baryon fraction \\citep{read05a}.  Comparing the metallicity of galaxies over a wide range of baryonic masses, \\citet{tremonti04a} found that dwarfs are metal depleted relative to massive galaxies and argue in favor of galactic winds efficiently removing enriched material from low mass galaxies.  However, at the masses of the galaxies in our sample, detailed models of the physics of gas blow-out and blow-away tend to disfavor significant gas mass loss, predicting loss of only a few percent \\citep{maclow99a,ferrara00a}.  Our sample cannot rule out small amounts baryon mass loss from these galaxies, nor scenarios in which gas remains bound to the galaxy but is blown away from the region participating in star formation.  This may explain the low observed metallicities of dwarf galaxies --- enriched gas is removed from regions of active star formation, resulting in low observed metallicities, but either remains bound or represents only a few percent of the total baryonic mass.  Indeed, recent cosmological simulations which include realistic feedback and chemical evolution suggest that the metallicity, mass-to-light ratios and other galaxies properties can be explained without including supernova energy feedback \\citep{tassis06a}.  Our results favor such dwarf galaxy formation models in which processes such as supernovae winds are not responsible for significant gas or baryon mass in low mass galaxies."
    },
    "0608/astro-ph0608540_arXiv.txt": {
        "abstract": "Millisecond X-ray pulsars consist of a rapidly-spinning neutron star accreting from a low-mass stellar companion, and are the long-sought evolutionary progenitors of millisecond radio pulsars, as well as promising candidate sources for gravitational radiation. The population of these sources has grown significantly over the last three years, with the discovery of six new examples to bring the total sample to seven. Three sources are ultracompact binaries with H-depleted donors and orbital periods of $\\approx40$~min, like the 185~Hz pulsar XTE~J0929$-$314. Three more have orbital periods of 2~hr or longer, similar to IGR~J00291+5934, first detected in outburst by {\\it INTEGRAL} in December 2004. The neutron star in this 2.46~hr binary has the most rapid spin of the accreting pulsars at 599~Hz. The most recently-discovered pulsar, HETE~J1900.1$-$2455 (377~Hz), has an intermediate orbital period of 83.3~min, and has been active for more than 1~yr, much longer than the typical transient outburst. Pulsations were detected only in the first few months of the outburst; this source has since resembled a faint, persistent non-pulsing low-mass X-ray binary, typical of the broader low-mass X-ray binary population. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608630_arXiv.txt": {
        "abstract": "We present 24 \\um images of three protoplanetary disks being photoevaporated around high mass O type stars. These objects have ``cometary'' structure where the dust pulled away from the disk by the photoevaporating flow is forced away from the O star by photon pressure on the dust and heating and ionization of the gas. Models of the 24 $\\mu$m and 8 $\\mu$m brightness profiles agree with this hypothesis. These models show that the mass-loss rate needed to sustain such a configuration is in agreement with or somewhat less than the theoretical predictions for the photoevaporation process. ",
        "introduction": "The discovery of ``proplyds'' in Orion \\citep[e.g.][]{Odel93,Odel94} let us image directly protoplanetary disks and study their structure. It also led to a number of theoretical studies of the photoevaporation of disks by external radiation fields from neighboring hot stars \\citep{John98,Rich98,Rich00,Mats03,Holl04,Thro05}. Many of these studies predict behavior that agrees well with the proplyd properties. The proplyds have been observed intensely, for example through detection of associated molecular hydrogen \\citep{Chen98} and high resolution measurements in the mid-infrared \\citep{Smith05}.  Deeper infrared observations can provide additional constraints for proplyd models. At 24 \\mm, the resolution of {\\it Spitzer} is sufficient to identify and perhaps resolve such objects at distances up to a few kpc. At longer wavelengths, the resolution of {\\it Spitzer} is too low to resolve these objects, and they will be lost in the confusion of the surrounding emission by interstellar material.  The predicted characteristics of an evaporating disk in the $\\sim$2 to 50 AU 24 $\\mu$m-emitting zone differ substantially among the studies (\\citet{John98, Rich00, Holl04, Mats03}. However, the calculations do generally agree that the process occurs on a time scale of $\\sim 1 - 3 \\times 10^5$ years. As a result, one might expect to observe a wide range of consequences of disk evaporation. For stars that have dwelt close to an O star since their formation, we might find severely truncated disks and possibly reduced output in the 24 \\um band. However, if the stellar velocities within a young stellar cluster are random, a star might enter the intense radiation field of an O star and begin the erosion of its disk well after its formation. A system that has just entered an O star radiation field is likely to be anomalously bright at 24 \\um because dust grains that would normally be shielded from the ambient ultraviolet photons in the optically thick disk are being torn from it by the evaporating gas and exposed to this field. \\citet{Rich98} calculate that the nominal disk emission at 24 \\um may be boosted more than an order of magnitude for some stages in the evaporation process. This paper describes three systems that are in this dramatic process of early disk evaporation. ",
        "conclusions": "We report the detection of three photoevaporating disks with cometary tails around high-mass stars in {\\it Spitzer}/24 $\\mu$m images. Their observed brightnesses and morphologies agree well with a model where the dust is entrained in the evaporating gas and pulled from the disk. Once exposed to the O star radiation field, a structure like a comet tail grows, with dust and gas streaming away from the O star under the influence of its radiation pressure on the dust and heating and ionization of the gas. We have roughly estimated the evaporation mass loss rate from our observations and find an upper limit in reasonable agreement with or somewhat lower than theoretical expectations. Further observations of these systems should be able to improve our understanding of the evaporation process substantially."
    },
    "0608/astro-ph0608406_arXiv.txt": {
        "abstract": "We present new observations leading to an improved black hole mass estimate for the Seyfert $1$ galaxy NGC~4593 as part of a reverberation-mapping campaign conducted at the MDM Observatory. Cross-correlation analysis of the H$\\beta$ emission-line light curve with the optical continuum light curve reveals an emission-line time delay of $\\tau_{\\rm cent}=3.73 \\pm 0.75$ days.  By combining this time delay with the H$\\beta$ line width, we derive a central black hole mass of $M_{\\rm BH} = (9.8 \\pm 2.1) \\times 10^{6} M_{\\odot}$, an improvement in precision of a factor of several over past results. ",
        "introduction": "Over the past two decades, reverberation mapping \\citep{Blandford82,Peterson93} has emerged as the most widely applicable method for directly measuring the masses of supermassive black holes (SMBH) in active galactic nuclei (AGNs).  The appeal of this method arises from its physical simplicity, despite the complicated physics within the central regions of these galaxies.  Moreover, unlike other methods for measuring the masses of SMBHs, it does not require high angular resolution.  On the other hand, this method is challenging because it requires a large series of observations, well spaced in time and taken over a long duration (at least three times the length of the longest physical timescale of interest).  The reverberation-mapping method depends on continuum flux variations from the central source, presumably an accretion disk around a SMBH. Continuum photons are absorbed and then re-emitted as line photons in the broad line-emitting region (BLR).  Variations in the broad-line flux have a time lag relative to the continuum, corresponding to the light-travel time across the BLR.  Through measurements of this time lag, the size of the region can be determined, which then leads, through virial arguments and a measurement of the broad-line width, to the black hole mass, $M_{\\rm BH}$.  To date, black hole masses have been determined for nearly 40 AGNs using reverberation mapping \\citep[see the compilation by][]{Peterson04}. While this demonstrates how far the field has progressed, \\citet{Peterson04} make it clear that substantial work is still needed to decrease uncertainties in time lag, and thus $M_{\\rm BH}$, measurements in order to better test and calibrate the various black hole mass-dependent scaling relationships, such as the $M_{\\rm BH}-\\sigma_{\\star}$ relationship, where $\\sigma_{\\star}$ is the bulge velocity dispersion.  This need for better reverberation-based mass estimates is the driving force behind the campaign we describe below.  One requirement for a successful reverberation mapping program is good temporal sampling \\citep{Horne04}.  The major goal of the current observing program is to improve the sampling rate for reverberation-mapped AGNs that were undersampled in previous campaigns. The mass of the black hole in NGC~4593 was previously measured by reverberation mapping \\citep{Dietrich94,Onken03}, but the H$\\beta$ lag determination of $\\tau_{\\rm cent}=3.1$ days was unfortunately much smaller than the average interval between observations of 15.8 days. This resulted in uncertainties ($+7.1,-5.1$ days) much larger than the time delay itself, and as such, were consistent with no lag at all.  By revisiting this object with our program, we were able to achieve the sampling necessary to drive down the previous high uncertainties in the time lag and black hole mass.  Here, we present new lag and $M_{\\rm BH}$ determinations from reverberation mapping of NGC~4593.  The observational uncertainties from this campaign represent a factor of several improvement over previous results. ",
        "conclusions": "Results presented here for the emission-line lag and black hole mass of NGC~4593 of $\\tau_{\\rm cent}=3.73 \\pm 0.75$ days and $M_{\\rm BH} = (9.8 \\pm 2.1) \\times 10^{6} M_{\\odot}$, for these respective quantities, represent a factor of several improvement over past measurements. Previous measurements for this object from LAG data exhibited much larger uncertainties due in part to the average sampling interval, which due mostly to bad weather, was much larger (by about a factor of 4) than the emission-line lag of this object.  In addition, for rather short observing campaigns, some of the success of reverberation mapping comes from serendipitously observing a large variability event during the campaign.  The best results come from observing a large increase and then decrease (or vice versa) in the flux, which we observed in this campaign but was unfortunately not seen to the same degree in the previous campaign of NGC~4593.  This represents yet another example exhibiting the need for longer observing campaigns.  The results reported here were obtained as part of a larger spectroscopic monitoring campaign whose primary goal was to improve the emission-line lag and black hole mass measurements for some of the nearest, apparently brightest AGNs.  The proximity of these AGNs makes them especially important not only as calibrators of the AGN $M_{\\rm BH}$--$\\sigma_*$ and BLR radius--luminosity relationships \\citep{Bentz06a}, but also as candidates for measurement of their black hole masses through other means that depend on high angular resolution. Key elements of this program were (1) scheduled daily observations of each of these relatively low-luminosity AGNs, (2) supporting observations at multiple sites to mitigate the effects of gaps in the time series caused by weather, and (3) high-quality homogeneous data that permitted relative spectrophotometric flux calibration at better than the 2\\% level. Our 42-night program yielded improved H$\\beta$ lags and black hole masses with a factor of several improvement in precision relative to previous campaigns for three of the six AGNs in our sample, NGC 4151 \\citep{Bentz06b}, NGC 5548, which was observed in the lowest-luminosity state yet recorded (Bentz et al., in preparation), and NGC 4593, as described here.  Three other AGNs in our program, NGC 3227, NGC 3516, and NGC 4051 were insufficiently variable during this relatively brief campaign, underscoring the point that longer duration programs are necessary for detection of variability signatures that are favorable for reverberation analysis."
    },
    "0608/astro-ph0608412_arXiv.txt": {
        "abstract": "\\noindent The presence of charge promotes the propensity of a rotating black hole to support the propagation of electromagnetic plane waves with negative phase velocity (NPV)  in its ergosphere, whether the Kerr--Newman or the Kerr--Sen metric descriptions of spacetime  are considered.   Striking  differences in the NPV characteristics for the Kerr--Newman and Kerr--Sen metrics emerge from numerical studies, particularly close to the outer event horizon when the magnitude of the charge is large. ",
        "introduction": "Electromagnetic plane waves with negative phase velocity (NPV) have been the centre of much attention lately \\c{LMW03}. They give rise to many exotic phenomenons~---~the most prominent being  negative refraction \\c{Rama}~---~which are not associated with their positive phase velocity (PPV) counterparts. Much of the interest in NPV propagation has stemmed from the realization of negatively refracting metamaterials \\c{SSS}. In addition,  the prospects of NPV propagation in special \\c{ML04a} and general \\c{MLS05} relativistic scenarios, and the concomitant implications for observational astronomy, have  prompted study and speculation within the past few years \\c{Current_Sci}.  It is now well established theoretically that, for certain spacetime metrics, NPV propagation can be supported in free space \\c{LMS05,Schwarzs}. In particular,   NPV can arise in the ergosphere of a rotating black hole described by the Kerr metric \\c{Tom}. The NPV--supporting regions are concentrated towards the equator of the ergosphere, and this concentration becomes exaggerated as the angular velocity of the black hole is increased. Here we address the question: how does the presence of charge effect the propensity of a rotating black hole to support NPV propagation?  The most commonly encountered spacetime description of a charged rotating black hole is provided by the Kerr--Newman metric \\c{exact}. This solution to the Einstein field equations was derived in the 1960s \\c{Kerr,KerrN}. In 1992, Sen developed an alternative metric solution describing a charged rotating black hole, based on the action of a heterotic string \\c{KerrS}. Since then, various aspects of the structure and properties of this Kerr--Sen metric have been established \\c{okai,Blaga,wu}. In the following sections, we consider both the Kerr--Newman and the Kerr--Sen metrics in our investigation of the NPV characteristics of a charged rotating black hole. ",
        "conclusions": "Our investigation has revealed that the propensity for a rotating black hole to support NPV propagation is highly sensitive to the presence of charge in the black hole. Specifically, \\begin{itemize} \\item the  overall incidence of NPV  increases as  $\\qrbh$ increases, for both the Kerr--Newman and Kerr--Sen metrics;  \\item as $\\qrbh$ increases, differences in the NPV characteristics associated with the  Kerr--Newman   and  Kerr--Sen metrics emerge; and  \\item the greatest differences between the NPV characteristics of the Kerr--Newman and the Kerr--Sen metrics arise close to the  outer event horizons of their respective ergospheres. \\end{itemize} These findings suggest the possibility of a NPV--based experiment to distinguish between the Kerr--Newman and Kerr--Sen metric descriptions of a charged rotating black hole.   \\vspace{20mm}  \\noindent{\\bf Acknowledgements:} BMR is supported by a UK EPSRC (grant GR/S60631/01) and a US NSF Graduate Research Fellowship. TGM and AL thank Dr. Sandi Setiawan (University of Edinburgh) for suggesting this study and Prof. Malcolm MacCallum (Queen Mary, University of London) for his assistance with  \\S\\ref{KS}."
    },
    "0608/gr-qc0608039_arXiv.txt": {
        "abstract": "\\noindent \\noindent We study the dipolar magnetic field configuration in dependence on brane tension and present solutions of Maxwell equations in the internal and external background spacetime of a magnetized spherical star in a Randall-Sundrum II type braneworld. The star is modelled as sphere consisting of perfect highly magnetized fluid with infinite conductivity and frozen-in dipolar magnetic field. With respect to solutions for magnetic fields found in the Schwarzschild spacetime brane tension introduces enhancing corrections both to the interior and the exterior magnetic field. These corrections could be relevant for the magnetic fields of magnetized compact objects as pulsars and magnetars and may provide the observational evidence for the brane tension through the modification of formula for magneto-dipolar emission which gives amplification of electromagnetic energy loss up to few orders depending on the value of the brane tension. ",
        "introduction": "It is well known that magnetic fields play an important role in the life history of majority astrophysical objects especially of compact relativistic stars which posses surface magnetic fields of $10^{12}G$. Magnetic fields of magnetars~\\cite{dt92,td95} can reach up to $10^{15}\\div 10^{16}G$ and in the deep interior of compact stars magnetic field strength may be estimated up to $10^{18}G$. The strength of compact star's magnetic field is one of the main quantities determining their observability, for example as pulsars through magneto-dipolar radiation. Therefore it is extremely important to study the effect of the different phenomena on evolution and behaviour of stellar interior and exterior magnetic fields.  Recently obtained solutions for relativistic stars on branes~\\cite{mm05}--\\cite{nkd00} have created an interest in the study of the influence and implications of the bulk tension on the various astrophysical processes, for example to the gravitational lensing~\\cite{k06,kp05a,kp05b} and to the motion of test particles~\\cite{i05}. To the best of our knowledge the effect of brane tension on magnetic field configuration in compact stars has not yet been studied. Since magnetic field determines the reach observational phenomenology of compact stars we will study here the consequences of the brane effects on stellar magnetic fields.  In the Newtonian framework the exterior electromagnetic fields of magnetized and rotating sphere are given in the classical paper of Deutsch~\\cite{d55} and interior fields are studied by many authors, for example, in the paper~\\cite{rt73}. In the general-relativistic approach the study of the  magnetic field structure outside magnetized  compact gravitational objects has been initiated by the pioneering work of Ginzburg and Ozernoy~\\cite{go64} and have been further extended by number of authors~\\cite{ac70}--\\cite{japan}, while in some papers~\\cite{getal98}--\\cite{zr} the work has been completed by considering magnetic fields interior relativistic star for the different models of stellar matter. General-relativistic treatment for the structure of external and internal stellar magnetic fields including numerical results has shown that the magnetic field is amplified by the monopolar part of gravitational field depending on the compactness of the relativistic star.  Here we will consider static spherically symmetric stars in the braneworld endowed with strong magnetic fields. The magnetic field structure will be assumed to be dipolar and axisymmetric and the effect of the gravitational field of the star in the braneworld on the magnetic field structure is considered without feedback, amounting to the astrophysical evidence that the magnetic field energy is not strong enough to affect the spacetime geometry. We find an exact analytical interior solution to Maxwell equations for the magnetic field when the stellar matter has unrealistic equation of state of stiff matter. When stellar matter has constant density the interior fields are found as numerical solutions of Maxwell equations. External magnetic fields are also found numerically. We show that both external and interior magnetic fields will be essentially modified by five-dimensional gravity effects.  In the paper we describe our model assumptions, present the set of the Maxwell equations to be solved together with boundary and initial conditions for the models under consideration and obtain analytical and numerical results for the stellar magnetic field. The paper is organized as follows. In section \\ref{meq} we provide a description of spherical compact objects in the braneworld and basic Maxwell's equations in the spacetime of these objects. In section \\ref{ss} we make main assumptions, discuss the boundary and initial conditions for the magnetic fields and also discuss interior magnetic fields for the different equations of state. In subsection~\\ref{internan} we find exact analytical solution for the stellar matter with the stiff equation of state. In subsection~\\ref{srst_es} we integrate external Maxwell equations from asymptotical infinity to the surface of star and find numerical solutions for magnetic field outside the braneworld star. In subsection~\\ref{internum} we numerically integrate the equation for magnetic field from the surface of the star till some inner radius.  As astrophysical application of the obtained results we look for the modification of luminosity of electromagnetic magneto-dipolar radiation from the rotating star due the braneworld effects in section~\\ref{application}.  Throughout, we use a space-like signature $(-,+,+,+)$ and a system of units in which $G = 1 = c$ (However, for those expressions with an astrophysical application we have written the speed of light explicitly.). Greek indices are taken to run from 0 to 3 and Latin indices from 1 to 3; covariant derivatives are denoted with a semi-colon and partial derivatives with a comma. ",
        "conclusions": "\\label{conclusion}  We have studied the stationary magnetic field of isolated relativistic compact star in the braneworld assuming that their magnetic fields are confined to the stellar crust of ideal perfect matter and have been working with the braneworld effects on the stellar magnetic field, accompanied by proper boundary conditions. In other words we generalized the general relativistic approach in the sense that we took into account the effect of additional braneworld tension on electromagnetic fields.  First we have solved interior Maxwell equations analytically and found exact solution for interior magnetic field inside the stellar stiff matter with unrealistic equation of state.  Then we have found numerical calculations which take into account the effect of brane tension on the structure of magnetic field outside the star and configuration of interior magnetic field for the stellar matter with constant density. Comparing the behaviour of the magnetic field when brane tension effects are included with the case when they are neglected, one can see enhancement of magnetic field especially near the surface of the relativistic star for external field and in the inner boundary taken for the interior field. This effect is stronger the more bigger the braneworld parameter $Q$ is.  Numerically calculated magnetic field structure for the external magnetic field, without brane tension effects incorporated, totally match the known analytical solution for magnetic field when 'branetension charge' $Q$ is equal to zero.  The numerical calculations made confirm that there are two effects of brane tension on magneto-dipolar emission. One is due to amplification of surface magnetic field by brane tension. Other is due to the presence of function $Q/r^2$ in the red-shift factor $\\sqrt{1-2M/r+Q/r^2}$ defined at the stellar surface in the expression for power of magnetodipolar radiation. We have found that the effect of branetension on magnetic fields of compact stars can be very important, and the expression for magnitodipolar luminosity of rotating braneworld magnetized star gives enhancement up to two orders.  While the previous papers extensively studied the influence of braneworld effects into the different astrophysical processes and cosmological problems, we conclude here the incorporation of the brane tension effect into magnetic field structure of the relativistic stars, could give an additional key for astrophysical evidence of the parameter $Q$."
    },
    "0608/astro-ph0608248_arXiv.txt": {
        "abstract": "{\\sf The present paper describes the first results of an observational program intended to refine and extend the existing $v\\sin i$ measurements of metal--poor stars, with an emphasis on field evolved stars.The survey was carried out with the FEROS and CORALIE spectrometers. For the $v\\sin i$ measurements, obtained from spectral synthesis, we estimate an uncertainty of about 2.0 \\kms. Precise rotational velocities \\ensuremath{v\\sin i} are presented for a large sample of 100 metal--poor stars, most of them evolving off the main--sequence. For the large majority of the stars composing the present sample, rotational velocities have been measured for the first time.} ",
        "introduction": "For the past two decades a large observational effort has been carried out at the Geneva Observatory, Switzerland, and at the Department of Physics at Natal, Brazil, to determine precise rotational velocity for a large stellar sample at different evolutionary stages and populations. The major aim behind such an effort is the understanding of the stellar angular momentum evolution. For population I evolved stars of luminosity classes IV, III, II, and Ib, covering the spectral range F to K, a large set of precise $v\\sin i$ measurements was already published (De Medeiros and Mayor 1999; De Medeiros et al. 2002, 2004). To carry on with these studies, we turn now to the measurement of $v\\sin i$ for field metal--poor stars. Since the pioneering works by Peterson (1983, 1985a,b), there have been an increasing number of observational studies on the rotational velocity of population II stars ( Peterson et al. 1995; Cohen and McCarthy 1997; Behr et al. 2000a,b; Carney et al. 2003; Recio--Blanco et al. 2002, 2004), dedicated mostly to field stars. Nevertheless, in spite of these solid observational efforts, the data presently available in the literature are not yet large enough in quantity for a statistically robust analysis of the rotational behavior of metal--poor stars in different evolutionary stages.  The present work brings the first results of an observational effort dedicated to the study of the rotational velocity of metal--poor stars, with an emphasis on field evolved stars. The paper is structured as follow: in Sect. 2 the main characteristics of the sample and the observational procedure are described. The data reduction, spectral synthesis, and the projected rotational velocity $v\\sin i$ measurements are presented in Sect. 3. ",
        "conclusions": ""
    },
    "0608/astro-ph0608281_arXiv.txt": {
        "abstract": "Although the surface of a magnetar is a source of bright thermal X-rays, its spectrum contains substantial non-thermal components.  The X-ray emission is pulsed, with pulsed fractions that can be as high as $\\sim 70\\,\\%$. Several properties of magnetars indicate the presence of persistent, static currents flowing across the stellar surface and closing within the magnetosphere. The charges supporting these currents supply a significant optical depth to resonant cyclotron scattering in the 1-100~keV band. Here we describe a Monte Carlo approach to calculating the redistribution of thermal seed photons in frequency and angle by multiple resonant scattering in the magnetosphere. The calculation includes the full angular dependence of the cyclotron scattering cross section, the relativistic Doppler effect due to the motion of the charges, and allows for an arbitrary particle velocity distribution and magnetic field geometry.  We construct synthetic spectra and pulse profiles for arbitrary orientations of the spin axis, magnetic axis, and line of sight, using a self-similar, twisted dipole field geometry, and assuming that the seed photons are supplied by single-temperature black body emission from the stellar surface.  Pulse profiles and 1-10~keV spectra typical of AXPs are easily produced by this model, with pulsed fractions of $\\sim 50\\%$.  However, this model cannot reproduce the hard, rising energy spectra that are observed from SGRs during periods of activity, without overproducing the thermal emission peak.  This suggests that the 1-100 keV emission of SGRs has a common origin with the hard X-ray emission detected from some AXPs above $\\sim 20$~keV. ",
        "introduction": "Magnetars have many remarkable properties that appear to derive from the decay of an ultrastrong magnetic field.  Measurements of spindown torques and bright X-ray outbursts have revealed dipole fields as strong as $\\sim 10^{15}$ G, and imply the existence of internal fields at least an order of magnitude stronger in some objects \\citep{td96,wt06}.  More recent measurements strongly suggest that the external magnetic field can support electric currents that are several orders of magnitude larger than the Goldreich-Julian current that flows along the open field lines. In some magnetars, the injection of these closed-field currents into the magnetosphere appears to be associated with Soft Gamma Repeater burst activity; but in others, the current appears to be sustained in the absence of bright SGR outbursts, in the class of magnetars known as the Anomalous X-ray Pulsars.  The basic properties of a twisted neutron star magnetosphere were outlined by \\citet{tlk02}. A detailed one-dimensional plasma model of the current in the closed magnetosphere has been constructed by \\citet{bt06}.  The presence of such a global twist has several observational effects, which mirror the behavior of magnetars in the following ways.  \\begin{enumerate} \\item \\emph{Correlation between spindown torque and X-ray hardness.} The presence of a twist in the external magnetic field causes it to flare out slightly.  As a consequence, the open-field flux is larger than estimated from the dipole model, and the spin-down torque is increased.  The charges flowing along the twisted closed field lines also supply a significant optical depth to resonant cyclotron scattering.  This optical depth is independent of radius (resonant frequency) when the magnetosphere is self-similar.  The thermal photons emitted from the surface of the star are multiply scattered. Thus the observed correlation between spin-down torque and hardness of the power-law component of the 1-10 keV X-ray spectrum \\citep{marsden01} can be explained by adjusting a single parameter:  the angle through which the magnetospheric field lines are twisted \\citep{tlk02}.  \\item \\emph{Persistent changes in X-ray pulse profiles.} Following the 1998 giant flare, the pulse profile of SGR~1900$+$14 changed to a single, nearly sinusoidal peak, from a more complicated profile with multiple peaks that was observed before the flare \\citep{woods01}. The change in the pulse profile persisted even though the quiescent X-ray flux returned to its pre-burst level.  Something similar happened to SGR 1806$-$20 after its 2004 giant flare, except that instead of simplifying, the pulse profile became more complex \\citep{woods06}. The AXP 1E~2259$+$586 also showed a systematic change in pulse profile following the 2002 outburst, although the duration of the change was shorter lived than in the SGRs \\citep{woods04}. It appears that the external magnetic field of a magnetar is reconfigured following a period of burst activity \\citep{woods01,tlk02}.  \\item \\emph{Persistent changes in $\\dot P$}.  The spindown torque of SGR~1900$+$14 and SGR~1806$-$20 can show large-amplitude modulations, increasing over the previous long term trend by at least a factor $\\sim 4$ (\\citealt{woods06}, and references therein).  The increase sets in gradually several months after a period of burst activity, and then persists for years.  Related but less dramatic effects are seen in the torque behavior of some AXPs \\citep{gavriil04}.  \\item \\emph{Hard X-ray emission.}  Emission above $20$~keV has been measured from both SGRs and AXPs (\\citealt{gotz06}, and references therein).  In the AXPs, the spectrum sometimes shows a strong upward break between 10 and 50 keV, with the output at 100 keV being 10 times stronger than the thermal emission from the neutron star surface.  It appears that the bolometric output of these objects is dominated by the magnetosphere. The additional high-energy component could be explained by bremsstrahlung emission from the transition layer near the neutron star surface, or possibly by a pair cascade at $\\sim 100$ km from the neutron star \\citep{tb05}.  \\end{enumerate}  In this paper, we confront the magnetar model with some of the observational evidence listed above.  We focus on the part of the non-thermal X-ray spectrum that lies directly above the black body peak, and also on the X-ray pulse behavior in this band (items 1 and 2 above). The most robust way of testing the twisted magnetosphere hypothesis is to produce synthetic analogs of observable quantities, using the self-similar model developed by \\citet{tlk02}. The radiation transport through the magnetosphere is sensitive to several geometrical effects, which include the angular emission pattern of cyclotron scattering by moving charges, the magnetic field geometry, orientation with respect to the observer, and the particle velocity distribution.  A first strategy is to ignore all the geometrical complications, focusing on the essential physics entering the multiple scattering process in a one-dimensional geometry. The problem of radiation transport through the accretion column of an X-ray pulsar has been treated in this way, using both Feautrier methods to solve the radiation transport equation \\citep{meszaros85a,meszaros85b} and Monte Carlo techniques \\citep{pravdo81, wang88}. The more general problem of the transfer of X-ray photons through a single resonant surface has been analyzed by \\citet{zheleznyakov96} and \\citet{lyutikov06} by restricting the transfer of photons to one spatial dimension and assuming a sub-relativistic and thermal distribution of scattering charges.  We do not expect narrow spectral features to form as the result of the transfer of thermal  X-ray photons through the quiescent, twisted magnetosphere. The cyclotron optical depth is regulated to a value close to unity over a broad range of frequencies.  By contrast, the optical depth is much higher in a pulsar accretion column, and is concentrated over a narrower range of radius and magnetic field. (In addition, the equivalent width of the absorption feature that forms in the cold atmosphere of a passively cooling magnetar is suppressed by polarization mode exchange \\citealt{holai04}.) The optical depth varies significantly with angle in the twisted dipole magnetic field, and the field lines impart a non-trivial angular structure to the particle velocity distribution.  This means that a numerical simulation in three dimensions is required to model the formation of a high-energy tail to the spectrum, and the amplitude and shape of the X-ray pulses that are seen by a distant observer.  In this paper, we report the development of a Monte Carlo code that follows the transfer of X-ray photons through the stellar magnetosphere, with resonant cyclotron scattering as the source of opacity. The code treats the scattering problem in three dimensions, for an arbitrary magnetic field geometry, and an arbitrary velocity distribution of the scattering charges.  We account for the relativistic Doppler effect and polarization exchange during scattering, but neglect the effects of recoil (which are important only at high energies where the spectra of AXPs and SGRs are not well measured, and where additional emission components are sometimes observed).  The output spectra and pulse profiles are obtained for an arbitrary line of sight, and for an arbitrary orientation of the spin axis and magnetic axis of the neutron star.  The paper is structured as follows.  In \\S~2 we review the physical input to the model, and in \\S~3 we describe the Monte Carlo code.  Results are presented and discussed in detail in \\S~4.  A comparison with the existing data and general conclusions follow in \\S~5. ",
        "conclusions": "We have developed a Monte Carlo code to study the reprocessing of thermal X-rays by resonant cyclotron scattering in neutron star magnetospheres using the model of \\citet{tlk02}. The code is fully three dimensional, and allows for an arbitrary velocity distribution of the scattering particles and an arbitrary magnetic field geometry.  Angle-dependent X-ray spectra and pulse profiles can be calculated for an arbitrary orientation between the magnetic axis and the spin axis, and between the spin axis and the line of sight to the star.  Our aim here has been to show which types of X-ray spectra and pulse profiles follow generically from the model, and which may point to different emission mechanisms.  We can reproduce most of the properties of the 1-10 keV emission of the AXPs, but not of the SGRs, by assuming a broad and mildly relativistic velocity distribution, and moderate twist angles ($\\Delta\\phi_{\\rm N-S} \\sim 0.3-1$ radian).   In particular, the SGRs in their most active states have harder X-ray spectra than can be reproduced by the model \\citep{woods06}. Our results demonstrate that changes in the strength of the twist and in the dynamics of the charged particles (which are coupled strongly to the radiation field) can lead to substantial changes in flux, hardness, pulse shape, and pulsed fraction.  \\subsection{Relative Strength of the Thermal and Non-thermal Components of the X-ray Spectrum} \\label{s:sumone}  The X-ray spectra of most magnetars display a prominent thermal component.  The corresponding blackbody area is less than, or sometimes comparable to, the expected surface area of a neutron star (e.g. \\citealt{wt06}).   In sources where the blackbody component is dominant at $\\sim 1-3$ keV, the high-energy photon index is typically $2$ or softer.  SGR 1900$+$14 presents an example of an SGR where a thermal component and a hard power-law component of the spectrum are simultaneously present \\citep{esposito06}.  Model spectra with a mildly relativistic particle distribution ($k_{\\rm B}T_0 \\simeq 0.5\\,m_ec^2, \\beta_0 \\simeq 0.75$; eq. [\\ref{eq:boltzone}]) can easily reproduce the 1-10 keV spectra of sources like 4U 0142$+$614, 1RXS J170849$-$4009 and 1E 1841$-$045. Some of our model spectra show a noticeable downward break at $\\sim 30$-$50\\,k_{\\rm B}T_{\\rm bb} \\simeq 10$-30 keV (e.g. Figs. 4f, 5f).  Such a break may be present in the combined XMM and Integral spectra of SGR 1900$+$14 \\citep{gotz06}, and its presence is not excluded in other sources such as the AXPs 4U 0142$+$614 and 1E 1841$-$045 when one takes into account the presence of a separate rising high-energy spectral component above 20 keV \\citep{kuiper06, gotz06}.  The quiescent SGRs have relatively weak blackbody components to their spectra, and during periods of extreme activity a pure power-law fit to the spectrum can have a photon index of $\\sim -1.6$ \\citep{mereghetti05, woods06}.  (Even harder photon indices are obtained from combined blackbody-powerlaw fits to the spectrum, but then the continuation of the 2-10 keV fit to higher energies disagrees with the Integral flux measurements.) We could obtain spectra as hard as this by choosing a broad power-law momentum distribution for the magnetospheric charges (Fig. 6), but the output spectrum then displays a strong blackbody component.  In the fit of \\citet{esposito06} for SGR 1900$+$14, the amplitude of the black body component is about three times that of the power-law component at the black body peak, whereas in the spectra of Fig. 6b, it is 6-8 times larger.  We expect this conclusion to hold for any model for the high-energy continuum which relies on the upscattering of the thermal seed photons.  Balancing the energy input to the coronal charges with the loss by upscattering, one finds that $\\tau (\\langle\\gamma^2\\rangle-1) \\sim 1$ (as in corona models based on non-magnetic Compton scattering).  When $\\langle\\gamma^2\\rangle^{1/2} \\gg 1$, it is generally not possible to upscatter a large fraction of the thermal photons because the optical depth $\\tau$ of the corona is small.   The results obtained with a broad relativistic particle distribution function should also be treated with caution: the charges are subject to a strong cyclotron drag force at a radius of $\\sim 50$-100 km where the cyclotron energy is 1-10 keV, and so it is probably not self-consistent to assume that they have the same distribution function there as they do in the inner parts of the magnetosphere.  These considerations point to the conclusion that the hard non-thermal spectral states of the SGRs (with photon indices $< 2$) have a similar origin to the {\\it excess} high energy emission that is observed from several AXPs above 20 keV.  In the magnetar model, this high-energy emission must involve a different emission process, such as bremsstrahlung from a hot transition layer between the relativistic corona and the colder neutron star atmosphere; or synchrotron emission triggered by pair cascades at a distance of $\\sim 100$ km, where the electron cyclotron energy is $\\sim 1$ keV \\citep{tb05}.  Indeed, there are reasons to expect that some active magnetars (e.g. the SGRs) can have significantly reduced X-ray emission due to deep internal cooling, even when the time-averaged rate of magnetic field decay remains high.   Large increases in the neutrino emissivity are expected when the core neutrons undergo a superfluid transition, which leads to an order-of-magnitude drop in blackbody X-ray flux \\citep{arras04}.  The greater intermittency of the magnetic dissipation in a Soft Gamma Repeater (as compared with the AXPs) could then be tied to the greater rigidity of the neutron star crust at lower temperatures. Given the high absorbing columns $N_H \\sim 2-6\\times 10^{22}$ cm$^{-2}$ that are inferred from the soft X-ray spectra of the two active SGRs, it would then be possible to hide the surface thermal component from detection. (The combination of a low thermal luminosity and strong absorption does not remove the inconsistency with the model spectra in Fig. 6c-e: the models then would predict that the non-thermal continuum is much dimmer than $\\sim 10^{35}$ ergs s$^{-1}$ at a photon energy of a few keV.)  \\subsection{Low Energy X-ray Spectra}\\label{s:lowen}  Our results show that the rescattering of thermal photons by magnetospheric charges does not lead to any enhanced emission at energies below the peak of the seed X-ray distribution. Our new fitting formula, eq.~(\\ref{eq:fit}), takes this low energy behavior of resonant scattering into account.  It should be kept in mind that the softest power-law components in measured spectra are typically needed to fit a {\\it low-energy} excess to a single-temperature black body. In our model most of the emission below the thermal peak comes from unscattered photons: our rotationally-averaged spectra have a Rayleigh-Jeans shape at low energies, which lies just below the input blackbody distribution.  We therefore suggest that a fit to an AXP X-ray spectrum involving a small $N_H$ and a very soft power law component is physically inconsistent, and that in such a case the fitting formula (\\ref{eq:fit}) is strongly preferred. Of course, the softest AXP spectra can also be fitted by a superposition of two or more black body components (e.g. \\citealt{halpern05}).  The formation of a high-energy tail to the black-body spectrum is closely tied, in this model, to an increased spindown rate (due to the flaring out of the poloidal field lines as they are twisted).  The AXP 1E 2259$+$586 provides an example of a magnetar which has demonstrated SGR-like burst activity \\citep{woods04}, but which also has the lowest spindown rate of any magnetar candidate, as well as very soft 1-10 keV spectrum, and only an upper bound on the 100 keV emission that has been detected from several other AXPs \\citep{kuiper06}.  We would therefore expect that the 1-10 keV spectrum of 1E 2259$+$586 is most accurately modelled by a variable surface temperature, rather than by a combination of thermal and non-thermal components.  Indeed, a recent reanalysis of the X-ray spectra of several AXPs has revealed a broader spectrum than a pure single-temperature black body near the thermal peak \\citep{durant06a}.  This suggests that the thermal X-ray flux is distributed inhomogeneously across the surface of the neutron star.  This is expected given the inhomogeneous nature of the yielding behavior in the crust of a magnetar, and the anisotropy of the thermal conductivity in the presence of a strong magnetic field \\citep{td96,let02,perez06,geppert05}.  \\subsection{Pulse Profiles and Pulsed Fractions}\\label{s:sumtwo}  Our calculations generally reproduce X-ray pulse profiles that are characteristic of the AXPs: single or multiple peaks, with sinusoidal shapes and sometimes the presence of narrow sub-pulses. The strong angular dependence of the optical depth to resonant scattering explains why phase resolved spectra show variations both in their flux level, as well as in hardness (e.g. \\citealt{tiengo05,goehler05}). We find that the X-ray spectrum is generally hardest at pulse maximum, in agreement with some measurements \\citep{mereghetti04, tiengo05}.  The number of sub-pulses depends essentially on the orientation of the spin and magnetic axes of the neutron star, which are not constrained by any independent measurements.  The pulse profiles of the actively bursting SGRs tend to be more complicated, which may reflect presence of high-order multipoles in the current flowing near the stellar surface.  As discussed in \\S \\ref{s:sumone}, their very hard X-ray spectra are more consistent with a separate emission process, possibly localized close to the surface of the neutron star.  Pulsed fractions as high as $\\sim 40-50\\%$ near the black body peak are easily obtained from the anisotropic scattering of thermal photons in the magnetosphere (Tables 1-3). Significantly higher pulsed fractions are possible in the power-law tail of the spectrum, due to the increased anisotropy of the multiply-scattered photons.  It should be emphasized that this result is obtained assuming a {\\it spherically symmetric} source of seed photons, i.e. a constant temperature across the surface of the neutron star.  Only in one object (the AXP 1E 1048.1$-$5957) is the pulsed fraction significantly higher (70-80\\%) in the 1-10~keV range \\citep{tiengo05}.  It should be possible to obtain higher pulsed fractions if the seed photons are emitted at localized hotspots, as suggested by the small black body emitting area \\citep{durant06b}.   Our calculations are therefore complementary to those of \\citet{pod00}, \\citet{opk01}, and \\citet{ozel02}, who studied the pulse profiles and pulsed fractions that result from isolated hotspots on rotating neutron stars, and compared their calculations with the observed 1-10 keV pulse profiles of the AXPs.  \\subsection{Comparison with Previous Work}  Semi-analytical work on resonant cyclotron scattering in the atmospheres of neutron stars goes back to the 1980s, the original focus being the transfer of radiation through the atmosphere of the star, or through an accretion column \\citep{yahel79,wasserman80,nagel80,meszaros85a,meszaros85b}.  In these calculations, the geometrical thickness of the scattering material was generally taken to be small compared with the gradient scale of the magnetic field. The effect of magnetic field inhomogeneity was analyzed by \\citet{zheleznyakov96}, who found solutions to the one-dimensional radiative transfer equation using the Schwarshild-Shuster method. \\citet{lyutikov06} generalized this method to include the combined effects of the particle motion and magnetic field inhomogeneity.  In the limit of small particle dispersion $\\beta c$, they found that the radiation spectrum that is transferred through an optically thick medium is upscattered in energy by a factor $1+2\\beta$.  This model was shown to provide a good fit to the 1-10 keV spectrum of 1E 1048.1$-$5937, the AXP source in which the thermal component of the spectrum peaks at the highest frequency.  However, this non-relativistic and one-dimensional approximation cannot provide a realistic description of the more extended high-energy spectra that are observed in some magnetars, including the slope of a power-law tail to the input thermal spectrum, the relative amplitude of the power-law and thermal components, or the detailed shape and frequency-dependence of the pulse profiles.  Monte Carlo codes for resonant cyclotron scattering through the surface layers of a neutron star were developed around the same time (e.g. \\citealt{pravdo81,wang88}), and also applied to the problem of line formation in gamma-ray bursts \\citep{wang93}.  The last two studies remain the most realistic to date, as they include the effects of thermal broadening, vacuum and plasma polarization, as well as electron recoil.  All of these effects but the last one are included in our study (plasma makes a negligible contribution to the dielectric properties of the magnetosphere at frequencies of $10^{17}-10^{19}$ Hz).  However, these studies are all one-dimensional and focus on the regime of very large optical depth and a weak gradient in the magnetic field, and so the results cannot be compared directly with ours.  \\subsection{Further Developments}  We now discuss how some of the key simplifications of the model will need to be relaxed in future work.  \\begin{enumerate} \\label{s:radpres} \\item \\emph{Anisotropic radiation pressure.} The most important simplification in the treatment of the transfer of radiation through the magnetosphere is the decoupling of the particle dynamics from the radiation force that is imparted through cyclotron scattering.  An electric field is induced by the twisted magnetic field lines, which lifts charged particles off the stellar surface, and accelerates them to high speeds (\\citealt{tlk02}, \\citealt{bt06}). Where the Lorentz-boosted cyclotron energy lies in the range 1-10 keV (or 1-100 keV for an active SGR), the particle moving against the direction in which the radiation is flowing feel a strong drag force, and a strong electric field is required to compensate \\citep{tb05}.  The charged particle dynamics and the radiation transfer are therefore strongly coupled.  \\item\\emph{Particle recoil.} The non-inclusion of particle recoil after scattering leads to significant errors in the shape of the output spectra above $\\sim 100$ keV. The spectra obtained with broad relativistic particle distributions involve essentially the single scattering of photons of a frequency $\\sim \\omega_{\\rm peak} = 2.82\\,k_{\\rm B}T_{\\rm bb}/\\hbar$. Recoil is important only when $\\gamma \\ga m_ec^2/2\\hbar\\omega_{\\rm peak}$ and $\\hbar\\omega \\sim \\gamma^2\\omega_{\\rm peak} \\ga 10$ MeV, where we do not tabulate spectra.  It should be emphasized that this remains a technical issue insofar as additional emission mechanisms dominate at $\\sim 30$-100 keV (see \\S \\ref{s:sumone}).  \\item \\emph{Seed photon distribution.} In all our simulations we have assumed that the seed photons are emitted radially, from points uniformly distributed over the stellar surface, and with a single polarization. There are several reasons for this choice.  First, the radius of first scattering by electrons (or positrons) is $R_{\\rm res} \\sim 10\\,R_{\\rm NS}$; second, we are not including the effects temperature inhomogeneities on the neutron star surface, focusing instead on in the pulsations that result from the transport of X-ray photons through a twisted magnetosphere. The effects of light bending must be included as soon as this second assumption is relaxed.   \\item \\emph{Polarization.}   Continuum spectra and pulsed profiles do not provide a clear way of discriminating between thermal emission from the surface due to deep cooling (E-mode emission) and due to particle bombardment and stopping by Coulomb collisions (O-mode emission). In both cases, the surface temperature will in practice be anisotropic (due to the channeling of the cooling heat flux along the magnetic field, and due to variations in the strength of the twist across the surface). In the second case, one expects the intensity of the O-mode photons to be more strongly beamed along the local direction of the surface magnetic field (e.g. \\citealt{basko75}), but we have found that pulsed fractions of 40-50\\% percent are achievable solely by radiation transport through the magnetosphere.  Polarization measurements of AXPs and SGRs would say much about the physics of the beam-heated atmosphere of a magnetar, and the distribution of magnetospheric currents, but they require polarimeters that are sensitive at flux levels of $\\sim 10^{-11}$ ergs cm$^{-2}$ s$^{-1}$ in the 1-10 keV band.  \\item \\emph{Magnetic field geometry.}  The magnetospheric model employed by \\citet{tlk02} is self-similar:  the pitch angle of the external field lines is independent of distance from the neutron star. Clearly different X-ray spectra and pulse profiles will result if the current is distributed in a different way throughout the magnetosphere. Observations which show a delayed increase in the spindown torque following periods of X-ray burst activity (\\citealt{woods02}) suggest that the twist is not distributed homogeneously throughout the magnetosphere in the immediate aftermath of a crustal yielding event. Quadrupole and octopole components of the magnetic field have a large effect on the pulse profile only if ions are the dominant scattering species, or if the non-thermal emission is localized close to the neutron star surface (as it is with bremsstrahlung emission; \\citealt{tb05}).  \\end{enumerate}"
    },
    "0608/astro-ph0608004_arXiv.txt": {
        "abstract": "We present new high-resolution observations of Sagittarius A* at wavelengths of 17.4 to 23.8 cm with the Very Large Array in A configuration with the Pie Town Very Long Baseline Array antenna.  We use the measured sizes to calibrate the interstellar scattering law and find that the major axis size of the scattering law is smaller by $\\sim 6\\%$ than previous estimates.  Using the new scattering law, we are able to determine the intrinsic size of Sgr A* at wavelengths from 0.35 cm to 6 cm using existing results from the VLBA.  The new law increases the intrinsic size at 0.7 cm by $\\sim 20\\%$ and $< 5\\%$ at 0.35 cm.  The intrinsic size is $13^{+7}_{-3}$ Schwarzschild radii at 0.35 cm and is proportional to $\\lambda^{\\gamma}$, where $\\gamma$ is in the range 1.3 to 1.7. ",
        "introduction": "Imaging the radio emitting region surrounding the massive black hole in the Galactic Center, Sagittarius A*, has been a goal since its discovery \\citep{1974ApJ...194..265B}. Turbulent electrons along the line of sight to Sgr A*, however, scatter radio wavelength photons and produce an image that is an elliptical Gaussian with a major axis size of $\\sim 0.5$ arcsec at 20 cm and a $\\lambda^2$ dependence \\citep{1978ApJ...222L...9B}.  Separating the effects of the small intrinsic source from the effects of scattering has required observations at short wavelengths, careful calibration, and the use of closure amplitude techniques, which reduce sensitivity but remove uncertainty due to calibration error \\citep[e.g.,][]{2004Sci...304..704,2005Natur.438...62S}.  These efforts have recently resulted in the first robust detections of intrinsic structure in Sgr A* at wavelengths of 1.3 cm, 0.7 cm and 0.35 cm. The intrinsic source has a size that scales with $\\lambda^{1.1}$ or $\\lambda^{1.6}$ to a minimum of $\\sim 10$ Schwarzschild radii at 0.35 cm (assuming $M_{bh}=4 \\times 10^6 M_\\sun$ and $d=8$ kpc ; \\citet{2005ApJ...620..744G,2003ApJ...597L.121E}).  These detections of the intrinsic size of Sgr A* have a number of consequences. The brightness temperature of $10^{10}$ K strongly excludes advection dominated accretion flows \\citep[ADAFs;][]{1998ApJ...492..554N} and Bondi-Hoyle accretion models \\citep{1994ApJ...426..577M}.  These size measurements, however, cannot differentiate between jet models \\citep{1993A&A...278L...1F}, generic radiatively inefficient accretion flows \\citep{2000ApJ...539..809Q}, and hybrids of these models \\citep{2002A&A...383..854Y}.  This limitation is principally due to the limited sensitivity in the minor axis size of the scattering ellipse.  Coupled with measurements of the proper motion of Sgr A* \\citep{2004ApJ...616..872R}, the assumption that the black hole is smaller than the emission region implies a lower limit to the mass density of the black hole $\\sim 10^5 M_\\sun {\\rm AU^{-3}}$, which  strongly excludes alternative models for dark mass objects.  The scattering medium itself is a system of intense interest \\citep{1998ApJ...505..715L,2001ApJ...558..127B,2006ApJ...640L.159G}. The $\\lambda^2$ size dependence of Sgr A* is a strict consequence of the strong scattering and the short projected baselines at the distances of the scattering medium \\citep{1989MNRAS.238..963N}.   We present here detailed measurements of the scattering properties at wavelengths that range from 17.4 to 23.8 cm using the Very Large Array and the Pie Town Very Long Baseline Array antenna (\\S 2).  The addition of the PT antenna to the VLA A configuration improves the East-West resolution by a factor of two. The resulting scattering law is smaller by $\\sim 6\\%$ than previous estimates. The independent estimate of the scattering law enables us to measure the intrinsic size of Sgr A* at a wavelength as long as 6 cm (\\S 3).  We discuss these results in (\\S 4).  We wish to clarify here the numerous axes and steps involved in translating observations of Sgr A* into a measurement of the intrinsic size.  VLA observations of Sgr A* are obtained with a synthesized beam that is extended North-South.  Deconvolution of the observed image with the synthesized beam gives the apparent, scatter-broadened image of Sgr A*.  This image is predominantly a two-dimensional Gaussian with the major axis oriented $\\sim 80$ degrees East of North.  Throughout this paper when we refer to the major and minor axes, we refer to the orientation of the scattering Gaussian.  Finally, to obtain the intrinsic image, we deconvolve the apparent image with a model of the scatter-broadened image, which is determined from long wavelength apparent sizes. ",
        "conclusions": "We have measured the intrinsic size of Sgr A* from 0.35 to 3.6 cm. At short wavelengths, the result is consistent with the conclusions of recent efforts by % \\citet{2004Sci...304..704} and \\citet{2005Natur.438...62S}. The size of Sgr A* at 0.35 cm is 13.3$^{+6.7}_{-3.1} R_s$, where $R_s=1.2 \\times 10^{12}$ cm is the Schwarzschild radius  for $M_{bh} = 4 \\times 10^6 M_\\sun$ and the Galactic Center distance $d=8$ kpc.  This compact size confirms tight restrictions on accretion models and black hole alternatives previously claimed and stated in \\S 1.  The wavelength dependence of the source size agrees with that found by % \\citet{2004Sci...304..704} and is steeper than that found by \\citet{2005Natur.438...62S}, who found $\\gamma\\approx 1.1$.  The steeper dependence indicates that the brightness temperature decreases as $\\lambda^{-1}$, assuming that the size in the second dimension is proportional to the major axis size. The peak brightness temperature at 0.35 cm is $\\sim 10^{10}$ K for a flux density of 1 Jy. The power-law dependence of the size as a function of wavelength indicates a stratified, smoothly varying emission region.  Detailed jet models for Sgr A* predict $\\gamma \\approx 1$ \\citep{2000A&A...362..113F}. Generalized jet models, however, allow a range of $\\gamma$, depending on the details of the magnetic field and particle energy density distributions % \\citep[e.g.,][]{1981ApJ...243..700K}. A jet with $B \\propto r^{-1}$, electron density decreasing as $r^{-1}$, and optically thin power-law index of 1 will show a size $\\propto \\lambda^{1.4}$.  \\citet{2006ApJ...642L..45Y} model the size of Sgr A* for a radiatively inefficient accretion flow.  They fit sizes at 0.35 cm and 0.7 cm that are fit with a power law of index $\\gamma=1.1$.    Variations in the nonthermal electron distribution or deviations from equipartition, however, could alter $\\gamma$ in their model.  The critical remaining observational goals for understanding the image of the radio emission of Sgr A* are a measurement of the two-dimensional size and detection of structural variability.  The simple one-dimensional deconvolution that we have performed here only gives schematic information on the size of Sgr A*.  With a more accurate two-dimensional scattering model, future analysis will directly compare the observed image with non-Gaussian source models convolved with the imaging and scattering constraints.  Astrometric observations may indicate a shift in the centroid of the image with frequency. A heterogeneous jet will exhibit such a shift due to changing location of the optically thick surface of the source \\citep{1981ApJ...243..700K}.  At mm and submm wavelengths, the gravity of the black hole will distort the image \\citep{2000ApJ...528L..13F,2005MNRAS.363..353B}. Detailed knowledge of the shape of the longer wavelength image will permit a more accurate characterization of light-bending effects in the actual image.  Ultimately, these images will provide one of the strongest tests of the existence and characterization of black holes."
    },
    "0608/astro-ph0608232_arXiv.txt": {
        "abstract": "Flares in Sagittarius A* are produced by hot plasmas within a few Schwarzschild radii of the supermassive black hole at the Galactic center. The recent detection of a correlation between the spectral index and flux during a near infrared (NIR) flare provides a means to conduct detailed investigations of the plasma heating and radiation processes. We study the evolution of the electron distribution function under the influence of a turbulent magnetic field in a hot collisionless plasma.  The magnetic field, presumably generated through instabilities in the accretion flow, can both heat the plasma via resonant wave-particle coupling and cool the electrons via radiation.  The electron distribution can generally be approximated as relativistic Maxwellian. To account for the observed correlation, we find that the magnetic field needs to be anti-correlated with the electron ``temperature''. NIR and X-ray light curves are produced for a cooling and a heating phase. The model predicts simultaneous flare activity in the NIR and X-ray bands, which can be compared with observations. These results can be applied to MHD simulations to study the radiative characteristics of collisionless plasmas, especially accretion flows in low-luminosity AGNs. ",
        "introduction": "It is generally accepted that the near-infrared (NIR) and X-ray flares in Sagittarius A*, the compact radio source associated with a mass $M\\sim 3-4\\times 10^6\\;M_\\odot$ supermassive black hole at the Galactic Center (Sch\\\"{o}del et al. 2002; Ghez et al. 2004; the Schwarzschild radius $r_S =10^{12}(M/3.4\\times 10^6\\;M_\\odot$) cm), are produced by relativistic electrons within a few Schwarzschild radii of the black hole via synchrotron and synchrotron self-Comptonization (SSC) processes, respectively (Baganoff et al. 2001, 2003; Genzel et al. 2003; Belanger et al. 2006; Eckart et al.  2004, 2006; Yusef-Zadeh et al. 2006a). Studying the details of the electron acceleration therefore plays an important role in revealing the properties of the flaring plasma and its interaction with the surroundings and the space-time created by the black hole, which may lead to a measurement of the spin with comprehensive flare observations and general relativistic MHD and ray-tracing simulations of the black hole accretion and the corresponding radiation transfer.  The electrons are likely accelerated by turbulent plasma waves generated in an accretion torus via instabilities, which induce the release of gravitational energy of the plasma (Balbus \\& Hawley 1991; Tagger \\& Melia 2006). In the simplest case, in which there is no large scale magnetic field, the turbulent magnetic field $B$ can play the dual role of both heating the electrons (presumably via resonant wave particle interactions) and cooling them (via SSC radiation), since the heating and cooling rates have different energy dependencies. Earlier, we showed that the {\\it steady-state} electron spectrum can be approximated as a relativistic Maxwellian distribution with the ``temperature'' $\\gamma_cm_ec^2$ depending on the plasma density $n$ and the coherent length of the magnetic field, which should be comparable to the size of the flare region $R$. Simultaneous spectroscopic observations in the NIR and X-ray bands (with four measured quantities: fluxes and spectral indexes in both bands) can then be used to determine the four basic model parameters, namely $B$, $R$, $\\gamma_c m_ec^2$ and $n$, and to test the model (Liu et al. 2006a, 2006b).  These investigations, however, also indicated that the steady-state solutions are not always applicable, especially for NIR flares with very soft spectra. A time-dependent approach to the evolution of the electron distribution under the influence of the turbulent magnetic field is required to study the flares in quantitative details. In light of the recent discovery of a correlation between the flux $F_\\nu\\propto\\nu^{\\alpha-1}$ and the spectral index $\\alpha$ in the NIR band (Ghez et al. 2005; Gillessen et al. 2006; Krabbe et al. 2006; Hornstein et al. 2006), we carry out such an investigation. We show that relativistic Maxwellian distributions in general give a good description of the electron spectra with the corresponding ``temperature'' evolution determined by the initial temperature, the magnetic field and its coherent length and the gas density. It is therefore possible to couple these results with MHD simulations to study the radiative characteristics of collisionless plasmas self-consistently.  Several scenarios have been suggested to explain the observed correlation between flux and spectral index (Gillessen et al. 2006). We find here that the dynamical processes of plasma heating and cooling are the most likely mechanism and the observation requires that the magnetic field be anti-correlated with the electron temperature. Since the evolution of the magnetic field depends on the turbulence generation mechanism, which needs to be addressed via MHD simulations, we calculate the evolution of the electron distribution function and the corresponding radiation spectrum for a cooling and a heating phase, where the magnetic field is set to be inversely proportional to the mean electron energy. The model naturally recovers the observed correlation between the NIR flux and spectral index and predicts the X-ray emission characteristics accompanying the NIR flares. Simultaneous flare observations in the NIR and X-ray bands can readily test the model and uncover the underlying physical processes producing these flares.  The kinetic equation of the electron acceleration and its time-dependent solutions are described in \\S\\ \\ref{acc}. In \\S\\ \\ref{apl} the model is applied to flares in Sagittarius A* by calculating the radiation from a flaring plasma with prescribed magnetic field evolutions. The implication and limitation of these results are discussed in \\S\\ \\ref{dis}. ",
        "conclusions": "\\label{dis}  Flares from the direction of the supermassive black hole in the Galactic center are produced within a few Schwarzschild radii. Studying the details of the flare properties plays an important role in understanding the physics of black hole accretion. To infer the underlying physical processes with simultaneous spectroscopic observations over a broad frequency range, the plasma process of electron acceleration has to be addressed. Given the closeness of the flare variation timescale, the dynamical time at the last stable orbit of the black hole, and the synchrotron cooling time of electrons producing NIR emission during the flares, a time dependent study of the particle acceleration is necessary. Although there are already attempts to carry out this study analytically (Becker et al. 2006; Schlickeiser 1984), numerical solutions are required when specific events are considered.  In this paper we study the time dependent solutions of electron heating and cooling by a turbulent magnetic field. When the amplitudes of large scale magnetic fields are much smaller than the turbulent component, there are four basic model parameters, namely the magnetic field and its coherent length, the gas density and the initial distribution of electrons. Electrons gain energy from the plasma waves and loss energy via radiation, and the energy dependences of the heating and cooling rates are quite different. The evolution of the electron distribution is determined by these parameters. We show that for the optically thin collisionless plasma in Sagittarius A*, synchrotron cooling dominates and relativistic Maxwellian distributions usually give a good description of the electron spectra. A formula is given to describe the dependence of the temperature evolution on the model parameters, which can be coupled with MHD simulations to study the radiation properties of the plasma.  We demonstrate that it is not trivial to produce the observed correlation between the spectral index and flux in the NIR band. Doppler effects and adiabatic processes suggested previously for the correlation (Gillessen et al. 2006; Yusef-Zadeh et al. 2006b) can not reproduce the observed results unless one has a very unusual electron spectrum. If the magnetic field does not change dramatically during the flare, only the cooling model of electrons heated instantaneously at the flare onset fits the observed correlation marginally. All other models predict very weak dependence of the spectral index on the flux density in the relatively hard spectral states.  We argue that the dynamical processes of electrons heating and cooling give the most reasonable explanation for the observations. To fit the correlation, the magnetic field needs to be anti-correlated with the electron temperature to reduce the flux density in the hard spectral states, bringing the above theoretical predictions more in line with the observations. It is not clear what mechanism may be responsible for such an anti-correlation. One possibility is that the dependence of the heating rate on the magnetic field is not as strong as that of the cooling rate. For example, the heating rate may be determined by the sound velocity, which can be much higher than the Alfv\\'{e}n velocity. Moreover, the cooling due to inverse Comptonization will dominate when the electron temperature and column depth are very high, which is certainly the case in the early cooling phase\\footnote{We note that the X-ray luminosity in the early cooling phase is higher than the synchrotron luminosity, suggesting a Compton catastrophe. This is due to the choice of the model parameters to embrace a large dynamical range. Including the cooling effect due to inverse Comptonization will make the electron temperature decrease rapidly to a value, where the Comptonization luminosity becomes lower than the synchrotron luminosity.}. All these effects can give rise to the desired anti-correlation.  The required anti-correlation between the magnetic field intensity and the electron mean energy suggests efficient energy exchanges between the field and electrons. While magnetic reconnection can cause transfer of magnetic field energy into electrons (corresponding to the heating phase), it is not clear what mechanism are responsible for the increase of magnetic field in the cooling phase. If the flare region is in pressure equilibrium with a relatively stable background, the magnetic field pressure may increase as the electron pressure decreases due to cooling.  We note that the magnetic field pressure increases from 1.27 to $1.27 \\times 10^4$ ergs cm$^{-3}$ and the electron pressure decreases from $6.1\\times 10^5$ to $6.1\\times 10^3$ ergs cm$^{-3}$ in the cooling phase of the flare discussed in \\S\\ \\ref{vmag}. For the heating phase, the magnetic field pressure decreases from $3.2\\times 10^3$ to 0.32 ergs cm$^{-3}$, while the electron pressure increases from 2.1 to $2.1\\times 10^2$ ergs cm$^{-3}$. Although the pressures are far from equipartition in the initial and final states, the observed correlation is produced when the two are evolving toward an equipartition. These suggest that the flares are triggered by a process, which drives the energy densities of the magnetic field and electrons far from equipartition. The relaxation of the system toward energy equipartition produces the observed correlation. It is therefore critical to produce and/or identify features far from energy equipartition in MHD simulations to uncover the physical processes driving the flares.  It is also possible that the flare region has complicated structure. The results of the electron spectral evolution presented here, which are also applicable to the relatively quiescent-state of the accretion flow, then need to be coupled with MHD simulations to study the radiation characteristics of the accretion flow. The observed correlation between spectral index and flux suggests that, if the magnetic field does not change dramatically during the flare and through the flare region, more electrons are involved in producing the softer spectral state flux. Regions with small hot cores flashing in a extended lower temperature active envelope in MHD simulations may be identified with the flares."
    },
    "0608/astro-ph0608142_arXiv.txt": {
        "abstract": "Photometric transit surveys promise to complement the currently known sample of extra-solar planets by providing additional information on the planets and especially their radii. Here we present extra-solar planet (ESP) candidates from one such survey called, the Wide Angle Search for Planets (WASP) obtained with the SuperWASP wide-field imaging system.  Observations were taken with SuperWASP-North located in La Palma during the April to October 2004 observing season.  The data cover fields between 23hr and 03hr in RA at declinations above +12. This amounts to over $\\approx$400,000 stars with V magnitudes 8 to 13.5. For the stars brighter than 12.5, we achieve better than 1 percent photometric precision. Here we present 41 sources with low amplitude variability between $\\approx$ 1 and 10 mmag, from which we select 12 with periods between 1.2 and 4.4 days as the most promising extrasolar planet candidates. We discuss the properties of these ESP candidates, the expected fraction of transits recovered for our sample, and implications for the frequency and detection of hot-Jupiters. ",
        "introduction": "The discovery of the large debris disk around $\\beta$ Pic \\citep{ST84} and other young stellar sources with disks \\citep{J98, L97}, plus the subsequent discovery of numerous extra-solar planets (ESP) has led to a resurgence in the study of planet and solar system formation. Discoveries of ESPs have been dominated by radial velocity surveys \\citep{MQ95, MB98, MB00, Ud00}. Currently, there are over 170 known ESPs orbiting around over 150 host stars. Photometric transit surveys have identified about 10 new planets, starting with: the confirmation of the first system, HD 209458b \\citep{Hen00, Ch00}, followed by several discovered with the OGLE project and confirmed with radial velocity follow-up \\citep{Kon03, Kon04, Po04, Bo05a, Kon05}, to the discovery of TrES~1 \\citep{Al04}, the detection of the transits from the radial-velocity discovered HD~189733 \\citep{Bo05b}, and the recent discovery and follow-up of XO-1b \\citep{Mc06, W06}. These planets typically have periods less than 4 days, orbital distance $\\sim$0.05 AU, and Jupiter-like masses and radii, and hence have been called the {\\it hot-Jupiters}. Additionally, several very short period (P $<$ 1.5 d) systems have been discovered \\citep{Kon03,Bo04} and termed {\\it very-hot-Jupiters}.  Radial velocity (RV) surveys measure the Doppler velocity of individual stars at many different times, which is a very time consuming procedure. The photometric transit method can measure thousands of stars simultaneously. In addition to monitoring a large number of stars, photometric transit detections provide information on the planet's radius and inclination, not readily available from RV surveys. It is currently not know what fraction of stars contain planets. The fraction of planets that may be hot-Jupiters is also not known. About 20\\% of the currently known ESPs have periods less than 10 days \\citep{B06} and $\\approx$10-15\\% of solar-type stars surveyed have ESPs \\citep{LG03}. There have been several attempts to better define what percentage of stars have planets. These range from 1-2\\% \\citep{Br03} to 20\\% \\citep{LG03}, but the latter also speculate this may be a lower limit. The large number of stars covered by photometric transit surveys should provide better constraints on the fraction of stars with hot-Jupiters, and may shed light on their formation mechanisms.  Here we present results from one such photometric transit survey, called WASP (the Wide Angle Search for Planets). WASP consists of 2 separate telescopes, called SuperWASP North (SW-N) and SuperWASP South (SW-S).  The SuperWASP telescopes were designed to cover a large field-of-view with better than 1\\% photometric accuracy down to 12th magnitude. Each SW telescope consists of 8 separate cameras, each of which covers 61 sq degrees of sky, and the combined system surveys nearly 500 sq degrees. The results presented here are from the 2004 SW-N observing campaign taken with a 5 camera configuration. A description of the SW instruments and observing strategy are presented in \\S\\ 2. In \\S\\ 3 we present the data reduction and analysis techniques for the ESP search. Candidates with secure low amplitude variability that may be ESPs are presented in \\S\\ 4. In \\S\\ 5, we discuss possible false alarms, expected fraction of recovered transits, and indications for the frequency of hot-Jupiters, their properties and observing follow-up strategies to confirm the transits planetary nature. Lastly, in  \\S\\ 6 we summarize our findings. ",
        "conclusions": "We have presented 41 low amplitude variables from the SW-N 2004 observing season. From this list we identified 12 ESPs candidates with periods of 1.2 to 4.4 days. These have transit depths $\\approx$1--4\\% implying planetary radii of 1.0--1.5 R$_J$ for their derived spectral types. High resolution optical spectroscopy to measure radial velocities and obtain orbital solutions and masses for these systems are needed to confirm these candidates are true extrasolar planets. Such observations are being undertaken by our consortium. SW-N and SW-S both continue to operate for the 2006 observing season, and with a full 8 cameras complement, and are expected to obtain more than twice the amount of data as the 2004 campaign.   \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{SuperWASP Candidate Stellar and Planetary Parameters} \\begin{tabular}{lcccccccccr} \\hline \\hline Name & V$_{SW}$ &  V$-$K& J$-$K & T$_{eff}$ & Spec. Type  & R$_{\\star}$ & R$_{p}$ &  $\\eta$ & Comment    \\\\ 1SWASP&          &       &       &           &             & R$_{\\sun}$ & R$_J$  &   &   \\\\\\hline  J230808.34+333803.9 & 10.8 &  1.41 & 0.28 &  6051 & F9 & 1.14 & 1.43 & 0.85 & Ellip. large \\\\ J231240.14+321540.4 & 11.8 & 1.46 & 0.30 & 5990 & G0 & 1.11 & 1.66 & 0.85 & R$_{p}$ too large \\\\ J231302.08+262724.3 & 10.6 &  3.68 & 0.87 &  4170 & M0 & 0.63 & 0.79 & 1.19 & Giant? \\\\ {\\bf J231533.56+232637.5} & 11.6 &  1.72 & 0.35 &  5661 & G6 & 0.96 & 1.39 & 0.13 & ESPc \\\\ J231807.73+240522.1 & 12.0 & 2.31 & 0.54 & 5034 & K3 & 0.76 & 1.63 & 0.74 & R$_{p}$ too large \\\\ J232639.10+233219.2 & 11.5 &  1.26 & 0.22 &  6256 & F7 & 1.25 & 2.04 & 0.81 & R$_{p}$ too large \\\\ J232700.56+200609.1 & 12.5 &  1.55 & 0.29 &  5870 & G2 & 1.05 & 2.16 & 0.73 & R$_{p}$ too large \\\\ J233325.17+332632.1 & 11.1 &  1.74 & 0.41 &  5637 & G6 & 0.95 & 1.01 & 1.25 & Ellip. large \\\\ {\\bf  J234318.41+295556.5} & 10.7 &  2.31 & 0.55 &  5034 & K3 & 0.76 & 1.09 & 0.84 & ESPc \\\\ J000029.25+323300.9 & 12.0 &  2.24 & 0.57 &  5101 & K2 & 0.77 & 1.01 & 1.29 & Ellip. large \\\\ J000233.27+331516.8 & 11.3 &  1.04 & 0.29 &  6573 & F4 & 1.40 & 2.07 & 0.60 & R$_{p}$ too large \\\\ {\\bf J002040.07+315923.7} & 11.8 &  1.51 & 0.31 &  5921 & G1 & 1.08 & 1.04 & 1.06 & ESPc \\\\ J002728.02+284649.2 &  9.6 &  2.84 & 0.71 &  4602 & K5 & 0.69 & 0.65 & 1.72 & Giant?/Large $\\eta$ \\\\ J002819.22+335122.1 & 12.5 &  1.52 & 0.35 &  5908 & G1 & 1.07 & 2.39 & 0.98  & R$_{p}$ too large\\\\ {\\bf J003039.21+205719.1} & 11.0 &  1.31 & 0.55  &  6186 & F8 & 1.21 & 1.23 & 0.92 & ESPc \\\\ J004804.21+202258.8 & 10.7 &  2.08 & 0.50 &  5260 & K1 & 0.82 & 1.78 & 1.1 & Ellip. large \\\\ J005107.84+214352.7 & 11.1 &  1.51 & 0.30 &  5921 & G1 & 1.08 & 1.90 & 0.95  & R$_{p}$ too large\\\\ {\\bf J005225.90+203451.2} & 11.2 &  1.89 & 0.46 &  5465 & G9 & 0.88 & 1.25 & 0.71 & ESPc \\\\ J005250.45+221038.8 & 12.3 &  1.47 & 0.33 &  5973 & G0 & 1.10 & 2.38 & 0.74 & R$_{p}$ too large \\\\ J005818.26+204348.0 & 12.4 &  1.77 & 0.38 &  5602 & G7 & 0.93 & 2.09 & 0.63 & R$_{p}$ too large \\\\ {\\bf J010151.11+314254.7} & 11.1 &  1.04 & 0.17 &  6573 & F4 & 1.40 & 1.42 & 1.04 & ESPc \\\\ J010553.38+241358.3 & 11.1 &  1.49 & 0.33 &  5947 & G1 & 1.09 & 1.58 & 0.53 & R$_{p}$ large \\\\ J010706.32+313918.0 & 10.3 & 1.55 & 0.30 & 5870 & G2 & 1.05 & 1.91 & 0.6 & R$_{p}$ too large \\\\ {\\bf J013033.21+311447.0} & 10.6 &  1.64 & 0.34 &  5758 & G4 & 1.00 & 1.07 & 1.08 & ESPc \\\\ J013100.45+374745.2 & 11.1 & 1.35  & 0.30 & 6130 & F8 & 1.18 & 1.90 & 0.78 & R$_{p}$ large  \\\\ J013250.08+194332.7 & 11.4 &  1.91 & 0.42 &  5442 & G9 & 0.87 & 1.76 & 0.14 & R$_{p}$ too large \\\\ J013452.76+293626.5 & 10.1 &  2.68 & 0.64 &  4720 & K5 & 0.70 & 0.40 & 1.60 & Giant/Large $\\eta$ \\\\ J014400.22+344449.2 & 11.2 &  1.30 & 0.25 &  6200 & F8 & 1.22 & 1.72 & 0.88 & R$_{p}$ too large \\\\ {\\bf J015625.53+291432.5} & 10.3 &  2.30 & 0.54 &  5044 & K3 & 0.76 & 0.72 & 1.67 & ESPc?/Large $\\eta$ \\\\ {\\bf J015711.29+303447.7} & 10.4 &  1.19 & 0.34 &  6354 & F6 & 1.30 & 1.38 & 0.77 & ESPc \\\\ J015951.59+354455.4 & 11.0 &  3.07 & 0.73 &  4453 & K7 & 0.67 & 0.86 & 0.08 & Ellip. large \\\\ J020720.96+325526.5 & 12.4 &  2.27 & 0.65 &  5072 & K2 & 0.77 & 1.40 & 1.51 & Ellip. large \\\\ J021217.50+335319.2 & 11.2 &  2.62 & 0.67 &  4767 & K5 & 0.71 & 1.59 & 0.92 & R$_{p}$ large  \\\\ J022421.03+375419.9 & 10.7 &  1.70 & 0.30 &  5685 & G5 & 0.97 & 1.94 & 0.81 & R$_{p}$ too large \\\\ {\\bf J022651.05+373301.7} &  8.3 &  0.83 &0.11 &  6896 & F1 & 1.52 & 1.47 & 0.74 & ESPc \\\\ {\\bf J023445.65+251244.0} & 12.1 &  1.86 & 0.35 &  5498 & G8 & 0.89 & 1.17 & 1.17 & ESPc \\\\ J024206.53+364029.7 & 11.8 &  1.47 & 0.29 &  5973 & G0 & 1.10 & 2.05 & 0.89 & R$_{p}$ too large \\\\ J025419.14+324240.7 & 11.7 &  1.61 & 0.37 &  5795 & G3 & 1.02 & 1.97 & 0.86 & R$_{p}$ too large \\\\ {\\bf J025500.31+281134.0} & 11.8 &  2.00 & 0.45 &  5344 & K0 & 0.84 & 1.34 & 1.29 & ESPc \\\\ J025712.69+403140.0 & 10.6 &  1.13 & 0.17 &  6441 & F6 & 1.34 & 1.53 & 0.80 & Ellip. large \\\\ J025958.91+294434.6 & 12.2 &  1.57 & 0.21 &  5845 & G2 & 1.04 & 2.11 & 0.71 & R$_{p}$ too large \\\\  \\hline \\end{tabular} \\end{minipage} \\end{table*}"
    },
    "0608/astro-ph0608468_arXiv.txt": {
        "abstract": "The population of clearly identified anomalous X-ray pulsars has recently grown to seven, however, one candidate anomalous X-ray pulsar (AXP) still eludes re-confirmation.  Here, we present a set of seven \\textit{Chandra} ACIS-S observations of the transient pulsar \\psr, obtained during 2003. Our observations reveal a faint X-ray point source within the \\textit{ASCA} error circle of \\psr's discovery, which we designate \\cxou\\ and tentatively identify as the quiescent AXP.  Its spectrum is well described by an absorbed single-component blackbody ($kT\\sim2.0$~keV) or power law ($\\Gamma\\sim1.0$) that is steady in flux on timescales of at least months, but fainter than \\psr\\ was during its 1993 period of X-ray enhancement by at least a factor of 13.  Compared to the outburst spectrum of \\psr, \\cxou\\ is considerably harder: if truly the counterpart, then its spectral behaviour is contrary to that seen in the established transient AXP \\xte, which softened from $kT\\sim0.67$~keV to $\\sim0.18$~keV in quiescence.  This unexpected result prompts us to examine the possibility that we have observed an unrelated source, and we discuss the implications for AXPs, and magnetars in general. ",
        "introduction": "\\label{intro}  The 6.97-s X-ray pulsar \\psr\\ was discovered serendipitously in archival \\textit{ASCA} observations from 1993 \\citep{gv98,tkk+98} during a period of apparent outburst.  Its slow spin period, soft spectrum ($kT\\sim0.64$~keV) and positional coincidence with the newly discovered supernova remnant G29.6+0.1 \\citep{ggv99} were suggestive of the small but growing class of anomalous X-ray pulsars (AXPs, see review by V. Kaspi, this volume\\footnote{For a summary of AXP properties, see the online catalog http://www.physics.mcgill.ca/$\\sim$pulsar/magnetar/main.html.}). However, without an estimate of $\\dot{P}$, and thus $B$, the AXP identification could not be confirmed, so further attempts were made to re-detect the pulsar and pulsations.  Unfortunately, it was not seen in a 1997 observation from the \\textit{ASCA} Galactic Plane Survey \\citep{tkk+98}, and a 1999 pointed follow-up observation with \\textit{ASCA} found a possible counterpart, \\axj\\footnote{The position of \\axj\\ has been corrected since the publication of \\cite{vgtg00}; see \\cite{tkgg06} for the best position.}, that was almost 10 times fainter, too faint for a measurement of pulsations or a spectrum \\citep{vgtg00}.  \\textit{Chandra}, \\textit{XMM-Newton} and \\textit{BeppoSAX} observations during 2001-2003 revealed a point source coincident with \\axj\\ and similar in brightness, but with a slightly harder absorbed spectrum ($kT\\sim1.0$~keV) than that seen for \\psr\\ in 1993 \\citep{isc+04}.  Presented here are the results of a \\textit{Chandra X-ray Observatory} monitoring campaign, conducted in 2003, with the goal of characterizing the spectral and timing properties of \\psr\\ in a post-outburst state. ",
        "conclusions": "We have observed and analysed the \\textit{Chandra} point source \\cxou, which may be the transient X-ray pulsar and candidate AXP \\psr. If it is the counterpart, then either \\psr\\ is not actually an AXP, or AXPs are much more diverse in their spectral and flux characteristics during quiescence than previously thought.  If it is not the counterpart and \\psr\\ is an AXP, then the exhibited flux variability presents a challenge to our current understanding of AXPs as magnetars, and hints at a much larger population of faint AXPs that remain undetected."
    },
    "0608/gr-qc0608105_arXiv.txt": {
        "abstract": "We report on a $\\lessapprox 1\\%$ test of the Lense-Thirring effect with the Mars Global Surveyor  orbiter and on certain features of motion of Uranus, Neptune and Pluto which contradict the hypothesis that the Pioneer anomaly can be caused by some gravitational mechanism. ",
        "introduction": "Up to now the Lense-Thirring effect\\cite{1,2,3,4,5} has been only tested in the terrestrial gravitational field with the LAGEOS satellites\\cite{6,7,8,9,10}. Although the relativistic predictions are not in disagreement with the results of such tests, their realistic accuracy has always been controversial\\cite{8,11,12,13}. Recent advances in planetary ephemerides$^{14}$ have made meaningful to compare the relativistic predictions for the Lense-Thirring effect of the Sun on the inner planets of the Solar System\\cite{15,16} to the least-squares determined extra-rates of perihelion of such celestial bodies$^{14}$. There is no contradiction between them; although the errors are still large so that also a zero-effect cannot be ruled out, the hypothesis of the existence of the solar gravitomagnetic field  is in better agreement with the data$^{16}$. In April 2004 the GP-B spacecraft\\cite{17,18} has been launched to measure the Schiff precession$^{19}$ of the spins of four superconducting gyroscopes carried onboard: the expected accuracy is $\\approx 1\\%$. The space environment of Mars has recently yielded the opportunity of performing another test$^{20}$ of the Lense-Thirring effect. Almost six years of range and range--rate data of the Mars Global Surveyor (MGS) orbiter, together with three years of data from Odyssey, have been used in order to precisely determine many global properties of Mars$^{21}$. As a by-product, also the orbit of MGS has been very accurately reconstructed$^{21}$. The average of the RMS overlap differences of the out-of-plane part of the MGS orbit amounts to 1.613 m over an about 5-years time span (14 November 1999--14 January 2005). Neither the gravitomagnetic force was included in the dynamical models used in the data reduction, nor any empirical out-of-plane acceleration was fitted, so that the RMS overlap differences entirely account for the martian gravitomagnetic force. The average out-of-plane MGS Lense-Thirring shift over the same time span amounts just to 1.610 m: a discrepancy of $0.2\\%$. The error has been evaluated as$^{20}$ $0.5\\%$. Let us, finally, note that the MGS test is based on a data analysis done in a completely independent way with respect to the author of Ref. $20$, without having gravitomagnetism in mind at all. ",
        "conclusions": ""
    },
    "0608/astro-ph0608697_arXiv.txt": {
        "abstract": "The PAMELA experiment is a satellite-borne apparatus designed to study charged particles in the cosmic radiation with a particular focus on antiparticles. PAMELA is mounted on the Resurs DK1 satellite that was launched from the Baikonur cosmodrome on June 15$^{th}$ 2006. The PAMELA apparatus comprises a time-of-flight system, a magnetic spectrometer, a silicon-tungsten electromagnetic calorimeter, an anticoincidence system, a shower tail catcher scintillator and a neutron detector. The combination of these devices allows antiparticles to be reliably identified from a large background of other charged particles. This paper reviews the design, space qualification and on-ground performance of PAMELA. The in-orbit performance will be discussed in future publications. ",
        "introduction": "The PAMELA (a Payload for Antimatter Matter Exploration and Light-nuclei Astrophysics) experiment is a satellite-borne apparatus designed to study charged particles in the cosmic radiation with a particular focus on antiparticles (antiprotons and positrons). PAMELA is installed inside a pressurized container attached to a Russian Resurs DK1 earth-observation satellite that was launched into space by a Soyuz-U rocket on June 15$^{th}$~2006 from the Baikonur cosmodrome in Kazakhstan. The satellite orbit is elliptical and semi-polar, with an altitude varying between 350~km and 600~km, at an inclination of 70$^{\\circ}$. The mission is foreseen to last for at least three years.  The PAMELA mission is devoted to the investigation of dark matter, the baryon asymmetry in the Universe, cosmic ray generation and propagation in our galaxy and the solar system, and studies of solar modulation and the interaction of cosmic rays with the earth's magnetosphere. The primary scientific goal is the study of the antimatter component of the cosmic radiation, \\begin{itemize} \\item in order to search for evidence of dark matter particle (e.g. non-hadronic particles outside the Standard Model) annihilations by precisely measuring the antiparticle (antiproton and positron) energy spectra; \\item in order to search for antinuclei (in particular, anti-helium); \\item in order to test cosmic-ray propagation models through precise measurements of the antiparticle energy spectrum and precision studies of light nuclei and their isotopes. \\end{itemize} Concomitant goals include, \\begin{itemize} \\item a study of solar physics and solar modulation during the 24$^{th}$ solar minimum by investigating low energy particles in the cosmic radiation; \\item reconstructing the cosmic ray electron energy spectrum up to several TeV thereby allowing a possible contribution from local sources to be studied. \\end{itemize} Table~\\ref{t:science} shows the design goals for PAMELA performance. The various cosmic-ray components and energy ranges over which PAMELA will provide new results are presented. \\begin{table}   % \\caption{\\label{t:science} Design goals for PAMELA performance.} \\begin{center} \\begin{tabular}{|c|c|} \\hline {\\bf Cosmic-ray particle}  & {\\bf Energy range} \\\\ \\hline Antiprotons & 80~MeV - 190~GeV \\\\ \\hline Positrons & 50~MeV - 270~GeV \\\\ \\hline Electrons & 50~MeV - 400~GeV \\\\ \\hline Protons & 80~MeV - 700~GeV \\\\ \\hline Electrons+positrons & up to 2~TeV \\\\ \\hline Light nuclei (up to Z=6) & 100~MeV/n - 250~GeV/n \\\\ \\hline \\hline {\\bf Antinuclei} & {\\bf Sensitivity 95\\% C.L.} \\\\ \\hline Antihelium/helium ratio & of the order of 10$^{-7}$  \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  Antiparticle measurements are the main scientific goal of the experiment. The precise determination of the antiproton and positron energy spectra will provide important information concerning cosmic-ray propagation and solar modulation. For example, indications of charge dependent solar modulation effects have been already seen in the antiproton to proton ratio data~\\cite{asa02}. Antiparticles could also be produced from exotic sources such as primordial black holes~\\cite{mak96} or the annihilation of supersymmetric~\\cite{ber99a} or Kaluza-Klein~\\cite{hoo05,bri05} dark matter particles. Figures~\\ref{pbflu} and~\\ref{posratio} show the current status of cosmic-ray antiproton and positron energy spectrum measurements, respectively. Theoretical calculations for pure secondary production~\\cite{sim98,ber99b,pro82,mos98} and for pure primary production due to the annihilation of supersymmetric dark matter particles~\\cite{ull99,bal99} are also shown. Almost all data available so far have been obtained by balloon-borne experiments. PAMELA will be able to perform very precise measurements with high statistics ($\\sim$10$^4$ \\app and $\\sim$10$^5$ \\ps per year) and over a wider energy range than possible to date. The full boxes in figures~\\ref{pbflu} and \\ref{posratio} indicate the expected PAMELA performance in case of a pure secondary antiproton and positron components and the full circles show the expected performance in case of an additional primary component. The errors on the expected PAMELA data points only include statistical uncertainties. An average PAMELA orbit has been used to estimate the vertical geomagnetic cut-offs and, consequently, the expected number of antiproton and positron events at low energies~\\cite{lun04}.  \\begin{figure} \\begin{center} \\epsfig{file=pbflu.eps, width=9cm} \\caption{Recent experimental \\app spectra (BES\\-S00 and BES\\-S99~\\cite{asa02}, AMS~\\cite{agu02}, CAPRICE98~\\cite{boe01a}, BES\\-S95+97~\\cite{ori00}, MASS91~\\cite{bas99}, CAPRICE94~\\cite{boe97}, IMAX92~\\cite{mit96}) along with theoretical calculations for pure \\app secondary production (solid lines: \\cite{sim98}, dashed line: \\cite{ber99b}) and for pure \\app primary production (dotted line: \\cite{ull99}, assuming the annihilation of neutralinos of mass 964~GeV/c$^2$). The expected PAMELA performance, in case of a pure secondary component (full boxes) and of an additional primary component (full circles), are indicated. Only statistical errors are included in the expected PAMELA data.} \\label{pbflu} \\end{center} \\end{figure} \\begin{figure} \\begin{center} \\epsfig{file=posratio.eps, width=9cm} \\caption{The positron fraction as a function of energy measured by several experiments (\\cite{gol87,mul87,cle96} and MASS89~\\cite{gol94}, TS93~\\cite{gol96}, HEAT94+95~\\cite{bar98}, CAPRICE94~\\cite{boe00}, AMS~\\cite{alc00}, CAPRICE98~\\cite{boe01b}, HEAT00~\\cite{bea04}). The dashed \\cite{pro82} and the solid \\cite{mos98} lines are calculations of the secondary positron fraction. The dotted line is a possible contribution from annihilation of neutralinos of mass 336~GeV/c$^2$ \\cite{bal99}. The expected PAMELA performance, for a pure secondary component (full boxes) and of an additional primary component (full circles), are indicated. Only statistical errors are included in the expected PAMELA data.} \\label{posratio} \\end{center} \\end{figure}  Another prominent goal of PAMELA is to measure the antihelium/helium ratio with a sensitivity of the order of 10$^{-7}$. This would represent a factor of 50~improvement on contemporary limits, as shown in figure~\\ref{antihe} as a function of rigidity (momentum /charge)~\\cite{hebar},\\cite{antihe_flux}. The contribution to the antihelium flux from cosmic ray interactions is expected to be less than 10$^{-12}$~\\cite{antihe_flux} and so an observation of antihelium would be a significant discovery as it could indicate the presence of antimatter domains in a baryon symmetric Universe.  \\begin{figure} \\begin{center} \\epsfig{file=he_hebar.ps, width=9cm} \\caption{The ratio of anti-helium to helium in the cosmic radiation shown as a function of rigidity~\\cite{hebar}. No observation of antihelium has been made to date and so upper limits are shown. The expectation for PAMELA after a 3~year long mission is shown.} \\label{antihe} \\end{center} \\end{figure}  The quasi-polar orbit and low geomagnetic cut-off experienced by the PAMELA apparatus combined with its intrinsic ability to measure low momenta will allow phenomena connected with solar and earth physics to be investigated~\\cite{cas04}.  The ability to measure the combined electron and positron energy spectrum up to 2~TeV will allow the contribution of local sources to the cosmic radiation to be investigated (e.g. see~\\cite{ato95}).  This article is organized as follows. The subdetector components of the PAMELA instrument are discussed in section~\\ref{sec:pamela} along with results from performance studies with particle beams and cosmic rays. The data acquisition and trigger systems are described in section~\\ref{sec:daq}. The Resurs DK1 satellite which houses PAMELA is presented in section~\\ref{sec:dk1}. Tests on various qualification models of the PAMELA apparatus are summarised in section~\\ref{sec:qual}. The physics performance of the flight instrument is presented in section~\\ref{sec:ground}. ",
        "conclusions": "PAMELA is a multi-purpose satellite-borne apparatus designed to study charged particles in the cosmic radiation with a particular focus on antiparticles (antiprotons and positrons). The energy range over which observations are made and the foreseen statistics mean that PAMELA stands to deliver results of great scientific relevance in several fields of cosmic ray research. To ensure reliable operation in space extensive space qualification tests of the PAMELA detector systems, electronics and mechanical structures were performed prior to launch. The performance of the individual detector components and the PAMELA system as a whole were also verified at particle beam facilities. The PAMELA instrument was launched into orbit from the Baikonur Cosmodrome on-board a Resurs DK1 satellite on June 15$^{th}$ 2006. All systems have been observed to operate as expected and scientific data analysis is now on-going."
    },
    "0608/astro-ph0608433_arXiv.txt": {
        "abstract": "{} {The goal of this work is to constrain the strength and structure of the magnetic field in the nearby cluster of galaxies A2255. At radio wavelengths A2255 is characterized by the presence of a polarized radio halo at the cluster center, a relic source at the cluster periphery, and several embedded radio galaxies. The polarized radio emission from all these sources is modified by Faraday rotation as it traverses the magnetized intra-cluster medium.  The distribution of Faraday rotation can be used to probe the magnetic field strength and topology in the cluster. } {For this purpose, we performed Very Large Array observations at 3.6 and 6\\,cm of four polarized radio galaxies embedded in A2255, obtaining detailed rotation measure images for three of them. We analyzed these data together with the very deep radio halo image recently obtained by us. We simulated random 3-dimensional magnetic field models characterized by different power spectra and produced synthetic rotation measure and radio halo images. By comparing the simulations with the data we are able to determine the strength and the power spectrum of the intra-cluster magnetic field fluctuations which best reproduce the observations.} {The data require a steepening of the power spectrum spectral index from $n=2$, at the cluster center, up to $n=4$,  at the cluster periphery and the presence of filamentary structures on large scales. The average magnetic field strength at the cluster center is 2.5\\,$\\mu G$. The field strength declines from the cluster center outward with an average magnetic field strength calculated over 1 Mpc$^3$ of $\\sim$1.2\\,$\\mu$G. } {} ",
        "introduction": "The existence of magnetic fields associated with the intracluster medium in clusters of galaxies is now well established through different methods of analysis (see e.g. the review by Govoni \\& Feretti 2004, and references therein). The strongest evidence for the presence of cluster magnetic fields comes from radio observations. Magnetic fields are revealed through the synchrotron emission of cluster-wide diffuse sources, and from studies of the rotation measure (RM) of polarized radio galaxies. Other techniques, not discussed in this paper, include the study of inverse Compton hard X-ray emission (e.g. Fusco-Femiano et al. 2004, Rephaeli et al. 2006), cold fronts  (e.g. Vikhlinin \\& Markevitch 2002) and magneto hydrodynamic simulations (e.g. Dolag et al. 2002).  Direct evidence for the presence of relativistic electrons and magnetic fields in clusters of galaxies comes from the detection, in an ever increasing number of galaxy clusters, of large-scale, diffuse, steep-spectrum synchrotron sources known as 'radio halos' or 'relics' depending on their morphology and location (e.g. Giovannini \\& Feretti 2002). Under the assumption that radio sources are in a minimum energy condition, it is possible to derive a zero--order estimate of the magnetic field strength averaged over the entire source volume. Typical equipartition magnetic fields, estimated in clusters with wide diffuse synchrotron emission, are 0.1-1 $\\mu$G. However, the equipartition estimate is critically dependent on the low energy cut-off of the relativistic electrons for steep spectrum sources such as radio halos. Since this quantity is not known the equipartition estimates of the magnetic field strength in these radio sources should be used with caution.  The presence of a magnetized plasma between an observer and a radio source changes the properties of the polarized emission from the radio source. Therefore a complementary set of information on clusters magnetic fields along the line-of-sight can be determined, in conjunction with X-ray observations of the hot gas, through the analysis of the RM of radio sources. Many high quality RM images of extended radio galaxies are now available in the literature (see e.g. the review by Carilli \\& Taylor 2002, and references therein). These data are consistent with magnetic fields of a few $\\mu$G throughout the clusters. In addition, stronger fields exist in the inner regions of strong cooling core clusters. In a few cases it has been possible to study the cluster magnetic field in more detail by sampling several extended radio galaxies located in the same cluster of galaxies (e.g. Feretti et al. 1999, Taylor et al. 2001, Govoni et al. 2001a).  It is worth noting that the RM observed toward radio galaxies may not be entirely representative of the cluster magnetic field if the RM is locally enhanced by compression of the intracluster medium due to relative motions. There are, however, several statistical arguments against this interpretation. In particular the statistical RM investigation of point sources (Clarke et al. 2001, Clarke 2004) shows a clear broadening of the RM distribution toward small projected distances from the cluster center indicating that most of the RM contribution comes from the intracluster medium.  This study included background sources, which showed similar enhancements as the embedded sources.  Recent work (e.g. En{\\ss}lin \\& Vogt 2003, Murgia et al. 2004) showed that detailed RM images of radio galaxies can be used to infer not only the cluster magnetic field strength, but also the cluster magnetic field power spectrum. Moreover, Murgia et al. (2004) pointed out that morphology and polarization information of radio halos may provide important constraints on the power spectrum of the magnetic field fluctuations on large scales. In particular, their simulations showed that if the intracluster magnetic field fluctuates up to scales of some hundred kpc, then steep magnetic field power spectra may give rise to detectable polarized emission in radio halos.  A2255 is a nearby (z=0.0806, Struble \\& Rood 1999), rich cluster with signs of undergoing a merger event (e.g. Yuan et al. 2005, Sakelliou \\& Ponman 2006). It is a suitable target to study the intracluster magnetic field because it is characterized by the presence of a diffuse radio halo source at the cluster center, a relic source at the cluster periphery, and several embedded radio galaxies (Jaffe \\& Rudnick 1979, Harris et al. 1980, Burns et al. 1995, Feretti et al. 1997). Govoni et al. (2005) found, for the first time in a cluster, that the radio halo of A2255 shows filaments of strong polarized emission ($\\simeq$ 20-40\\%). The distribution of the polarization angles in the filaments indicates that the cluster magnetic field fluctuates up to scales of $\\sim$400 kpc in size. In the rest of the cluster they did not detect significant diffuse polarized emission except in the brighter regions of the relic ($\\simeq $15-30\\%).  Here we present multi frequency (3.6 and 6 cm) Very Large Array (VLA\\footnote{The Very Large Array is a facility of the National Science Foundation, operated under cooperative agreement by Associated Universities, Inc.}) observations of four polarized radio galaxies in the A2255 cluster. We apply the numerical approach proposed by Murgia et al. (2004), of analyzing the polarization properties of both radio galaxies and halo, to investigate the cluster magnetic field strength and power spectrum.  The paper is organized as follows. In Sect. 2 we discuss details about the radio observations and the data reductions. In Sect. 3 we present the total intensity and polarization properties of the four radio galaxies at 3.6 and 6 cm. In Sect. 4 we present the RM images, discuss the results, and consider the presence of a cluster magnetic field. In Sect. 5, by following the same approach presented in Murgia et al. (2004), we introduce the 3D multi-scale magnetic field modeling used to determine the intra-cluster magnetic field strength and structure. in Sects. 6 and 7 we show the results obtained with a constant and variable magnetic field power spectrum slope, respectively. In Sect. 8, we draw conclusions from this study.  Throughout this paper we assume a $\\Lambda$CDM cosmology with $H_0$ = 71 km s$^{-1}$Mpc$^{-1}$, $\\Omega_m$ = 0.3, and $\\Omega_{\\Lambda}$ = 0.7. At the distance of A2255, 1\\arcsec corresponds to 1.5 kpc. ",
        "conclusions": "We present new VLA observations at 3.6 and 6\\,cm for four polarized radio galaxies embedded in A2255, obtaining detailed RM images for three of them. The dispersion and the mean of the RM decrease with increasing distance from the cluster center. We analyze these data, together with the very deep radio halo image recently obtained by Govoni et al. (2005). Using the numerical approach described in Murgia et al. (2004), we simulate random 3-dimensional magnetic field models characterized by different power spectra and produce synthetic RM and radio halo images. By comparing the simulations with the data we investigate the strength and the power spectrum of the intra-cluster magnetic field fluctuations.  We find that a magnetic field power spectrum with a single spectral index is not suitable to reproduce both the rotation measure distribution of radio galaxies, and the radio halo polarization. The data require a steepening of the power spectrum spectral index from $n=2$, at the cluster center, up to $n=4$ at the cluster periphery and the presence of filamentary structures on large scale. This result seems to indicate that the magnetic field at the cluster center would be more dominated by small-scale structures, while at the periphery most of the magnetic field energy density is concentrated on the large-scale structures. A magnetic field with power spectrum steepening with radius is likely to reflect a complex behaviour of turbulence and motions in clusters. Subramanian et al. (2006) discuss the evolving turbulence due to dynamo action, and find that turbulence can coexist with ordered filamentary gas structure. Since turbulence is likely driven by a merger event, it may be expected that this driving is stronger in the cluster central region. In this case, turbulence in the peripheral regions can be in a more advanced stage of decay than at the cluster center.  Since smaller scales decay first, and the integral scale of decaying turbulence increases with time, larger-scale magnetic field structures may be more easily present in the outer cluster region.  The average magnetic field strength at the cluster center is 2.5\\,$\\mu G$. The field strength declines from the cluster center outward and the average magnetic field strength calculated over 1 Mpc$^3$ is about 1.2\\,$\\mu$G.  In conclusion, we find that the above magnetic field model provides a good description of the RM radial profiles and is able to reproduce both the luminosity and the fractional polarization of the radio halo, thus reconciling the magnetic field strength estimated using two complementary methods of analysis."
    },
    "0608/astro-ph0608119_arXiv.txt": {
        "abstract": "We compare a new $R=120,000$ spectrum of PG~$1634+706$ ($z_{QSO}=1.337,m_V=14.9$) obtained with the HDS instrument on Subaru to a $R=45,000$ spectrum obtained previously with HIRES/Keck. In the strong {\\MgII} system at $z=0.9902$ and the multiple cloud, weak {\\MgII} system at $z=1.0414$, we find that at the higher resolution, additional components are resolved in a blended profile.  We find that two single-cloud weak {\\MgII} absorbers were already resolved at $R=45,000$, to have $b=2$-$4$~{\\kms}.  The narrowest line that we measure in the $R=120,000$ spectrum is a component of the Galactic {\\NaI} absorption, with $b=0.90\\pm0.20$~{\\kms}.  We discuss expectations of similarly narrow lines in various applications, including studies of DLAs, the {\\MgI} phases of strong {\\MgII} absorbers, and high velocity clouds. By applying Voigt profile fitting to synthetic lines, we compare the consistency with which line profile parameters can be accurately recovered at $R=45,000$ and $R=120,000$. We estimate the improvement gained from superhigh resolution in resolving narrowly separated velocity components in absorption profiles. We also explore the influence of isotope line shifts and hyperfine splitting in measurements of line profile parameters, and the spectral resolution needed to identify these effects. Super high resolution spectra of quasars, which will be routinely possible with 20-meter class telescopes, will lead to greater sensitivity for absorption line surveys, and to determination of more accurate physical conditions for cold phases of gas in various environments. ",
        "introduction": "\\label{sec:1}  With the advent of high resolution spectroscopy on 8-meter class telescopes in the early 1990's came a revelation in the study of quasar absorption line systems.  Previous to this time, several large surveys were conducted with a typical resolution of $R = \\lambda/\\Delta \\lambda = 1000$--$3000$, corresponding to $100$--$300$~{\\kms}.  These surveys tabulated the redshift path densities of {\\MgII}, {\\CIV}, and {\\Lya} lines over the redshift range accessible from optical telescopes \\citep*{lanzetta87, sbs88, caulet}. However, most individual systems were not resolved, so that the kinematics of the different metal lines could not be compared system by system.  It is not surprising that an increase of resolution to $R=45,000$ ($\\sim7$~{\\kms}) provided detailed views of the multiphase structure of gas inside and around galaxies, since parcels of gas within the different components of galaxies are often moving at several to tens of {\\kms} relative to one another.  The increase of resolution from thousands to tens of thousands also brought a dramatic increase in the sensitivity of surveys.  This allowed measurements of the metallicity in the most diffuse gas sampled by weak {\\Lya} forest lines \\citep[e.g., ][]{cowie95}, and for deeper surveys of metal lines as well \\citep[e.g., ][]{weak1,toru02}.  Presently, hundreds of quasar spectra have been obtained with $R\\sim45,000$ using the High Resolution Echelle Spectrograph on Keck I \\citep[HIRES, ][]{hires}, the High Dispersion Spectrograph on Subaru \\citep[HDS, ][]{hds} and the Ultraviolet-Visual Echelle Spectrograph on the VLT \\citep[UVES, ][]{uves}.  There is undoubtedly much to be gained from detailed studies of the absorption lines detected in these databases.  However, in this paper we look toward the future and ask what advances we should expect in studies of quasar absorption lines with an increase in resolution from the presently attainable value of $7$~{\\kms} to $0.5$-$2$~{\\kms}, which is within the reach of 20-meter class telescopes for hundreds of quasars.  This follows a study by \\citet{tappeblack} of the quasar PKS~$2145+067$ at $R\\sim100,000$ with UVES on the VLT UT2 telescope.  That study focused on the $z=0.79089$ strong {\\MgII} absorber, and measured $b$ parameters in the range $1$--$7$~{\\kms}, showing considerably more structure than previous spectra \\citep{cwcvogt01} with $R\\sim45,000$.  More recently, \\citet{chand06} considered the effect of observing at $R>100,000$ for studies of variation in the fine-structure constant.  In addition to the increased sensitivity to weak lines, gained at higher resolution, there is reason to expect we should see substructure on {\\kms} scales.  Distinct small-scale structures are seen in the interstellar medium of the Milky Way on these velocity scales \\citep{points}.  Very narrow components, indicating cold clouds, are seen through 21-cm absorption in some damped {\\Lya} systems \\citep{lane00}.  Also, detailed models of transitions detected in DLA systems indicate the presence of a cold phase \\citep{wolfe03}.  The large {\\Mgn} to {\\Mgi} ratios in many components of strong {\\MgII} absorbers, most of which are produced by lines of sight through $L^*$ galaxies, could be explained by a cold phase that produces the bulk of the {\\MgI} absorption \\citep{cwcvogtchar03}. In the case of the $z=0.9902$ system toward PG$1634+706$, \\citet{z99} proposed cold clouds with Doppler parameters of $b(Mg) \\sim 0.75$~{\\kms} to explain the strong {\\MgI} absorption, but these could not be resolved at FWHM$\\sim7$~{\\kms}.  Finally, VLA observations of Milky Way high velocity clouds indicate that they have $T<900$~K, and most Milky Way lines of sight show narrow {\\CaII} lines from HVCs \\citep{richter05}.  However, at FWHM$\\sim7$~{\\kms} these {\\CaII} lines were unresolved. \\citet{richter05} suggested that if the Milky Way is typical of spiral galaxy halos these narrow absorption features should be commonly seen in the profiles of strong {\\MgII} absorbers and Lyman limit systems.  For the brightest quasars it is already feasible to use the highest resolution settings on 8-meter class telescopes ($\\sim$~40 with $m_V < 16$ for non-exorbitant exposure times) to study quasar absorption line systems.  In this paper, we present an exploration of the power of $R=120,000$ in studying four {\\MgII} absorbers along the line of sight toward the magnitude $m_V=14.9$ quasar PG$1634+706$.  Our observations were carried out in the service mode with the HDS on the Subaru Telescope.  The paper begins, in \\S~\\ref{sec:2}, with a general consideration of how accurately line parameters can be determined at super-high resolution. In \\S~\\ref{sec:3}, we present our Subaru/HDS observations of several absorption line systems in the spectrum of one of the brightest quasars, PG~$1634+706$, and discuss the implications of our $R=120,000$ observations of these systems.  We also show the Milky Way absorption along this line of sight.  In \\S~\\ref{sec:4}, we speculate about the progress that will be made in the study of quasar absorption lines when many quasars are observed at $R>100,000$. ",
        "conclusions": "\\label{sec:4}  The ability to routinely obtain quasar spectra with resolutions of $R>100,000$ is not likely to provide as large a change in the field of quasar absorption lines as the previous advance to $R=45,000$, made possible with $8$--meter class telescopes.  However, the ability to probe structure on scales of $\\sim1$~{\\kms} will certainly be valuable in select applications and in detailed studies of particular absorbers.  Examples of such applications include probing the physical conditions of cold regions in the ISM of galaxies, particularly in DLAs and other strong {\\MgII} absorbers.  Also, it appears that high velocity clouds, both around the Milky Way and at high redshift, may have a cold phase. Finally, surveys extending to small equivalent widths will be aided by the enhanced sensitivity of higher resolution observations.  Our comparison of $R=120,000$ Subaru/HDS observations to $R=45,000$ Keck/HIRES observations of various systems toward PG~$1634+706$ show that at the higher resolution, more accurate Doppler parameter measurements can be made, and that additional Voigt profile components are resolved.  However, for the particular {\\MgII} absorbers along the PG~$1634+706$ sightline no extremely narrow individual components were measured.  It was only for the Galactic {\\NaI} line in this line of sight that $b<1$~{\\kms} was measured.  Because such narrow components are expected in a variety of environments, it would be valuable to push the limits of 8-meter class telescopes in order to detect them in a variety of types of systems toward the brightest quasars.  We are grateful to Chris Churchill for giving us permission to use the Keck/HIRES spectrum for this work.  We also wish to thank the referee, John Black, for several insightful comments that have increased the scope of this paper."
    },
    "0608/astro-ph0608269_arXiv.txt": {
        "abstract": "We investigate the value of future dark energy experiments by modeling their ability to constrain the dark energy equation of state. Similar work was recently reported by the Dark Energy Task Force (DETF) using a two dimensional parameterization of the equation of state evolution.  We examine constraints in a nine dimensional dark-energy parameterization, and find that the best experiments constrain significantly more than two dimensions in our 9D space. Consequently the impact of these experiments is substantially beyond that revealed in the DETF analysis, and the estimated cost per ``impact'' drops by about a factor of ten as one moves to the very best experiments. The DETF conclusions about the relative value of different techniques and of the importance of combining techniques are unchanged by our analysis. ",
        "introduction": "The observed cosmic acceleration requires ``dark energy'' to achieve consistency with the current cosmological paradigm.  The dark energy must be the dominant component of the universe today (about 70\\% of the energy density), yet an understanding of its fundamental nature has proved elusive. Many believe that resolving the mystery of the dark energy will force a radical change in our understand of fundamental physics. This expectation has generated great interest in the dark energy and widespread enthusiasm for an aggressive observational program to help resolve this mystery.  Recently, the Dark Energy Task Force (DETF)\\cite{Albrecht:2006um} released a report to guide the planning of future dark energy observations. The DETF used a dark energy ``figure of merit'' (FoM) based on a two-parameter description of the dark energy evolution in order to produce quantitative findings. An interesting question is whether using the DETF FoM might lead to poor choices in shaping an observational program because of its simplicity.  In particular, could the relative value of two possible experiments be distorted by the DETF FoM?  We consider this question by examining an alternative FoM.  We model the dark energy evolution with a multi-parameter model and formulate a FoM (the ``D9 FoM'') which gives an experiment credit for {\\em any} constraint it places on the dark energy evolution. Because theory currently offers little guidance on the functional form of this evolution, we should seek from experiment any available constraint on its behavior. We use this alternative FoM to assess many of the same simulated data sets (or ``data models'') considered by the DETF.  The DETF dark energy parameters are almost completely unconstrained by current data which is typically analyzed in smaller parameters spaces in order to manifest a noticeable impact (see for example \\cite{Spergel:2006hy}). In contrast, the DETF approach of evaluating experiments making no prior assumptions on the cosmic curvature and using a two-parameter model for the equation of state looks very ambitious ({\\em too} ambitious to some \\cite{Linder:2005nh}). Our work shows that the best data models constrain significantly more than two equation-of-state parameters and thus their impact was underestimated even using the ambitious DETF parameterization. As a consequence, we find that the highest quality large-scale projects also have a much lower estimated cost per FoM improvement than is achieved by the medium scale projects. Our work represents only one of several interesting advances that can reveal impact of dark energy experiments that is greater than seen by the DETF (see for example \\cite{Zhan:2006gi,Schneider:2006br}). In order to give a more focused discussion our new FoM is the only significant technical difference between our methods and those of the DETF. ",
        "conclusions": "We have analyzed a FoM for dark energy probes which is defined in a nine-dimensional parameter space, up from the (already ambitious) two-parameter space used by the DETF. Our 9D FoM gives a more complete account of the impact of a given data model. We find the DETF data models constrain significantly more parameters than the two used by the DETF, leading to vastly improved FoM's in the D9 space. Our rescaling law gives an intuitive account of the impact of measuring more parameters.  The 9D FoM follows the same logic as the DETF FoM, namely it attaches equal weight to any constraint on $w(a)$.  The measurement of a non-zero value for any of the independently measured combinations $\\vec{f}_i$ would  be equally significant in that it would rule out a pure cosmological constant. Thus the higher values of the 9D FoM (vs the DETF) reflect genuinely greater discovery power.  The effective number of parameters measured ($D_e$) increases with the quality of the project, and consequently so does the gap in the impact registered by the 9D vs DETF FoM's.  Good Stage 4 projects measure more than four parameters and produce a three or four order of magnitude improvement in the 9D FoM vs a one order of magnitude improvement in the DETF FoM. As a result, our 9D work indicates that good Stage 4 projects achieve a ten times lower estimated cost per FoM increase than Stage 3 projects, the reverse of the DETF conclusion.  Aside from this major difference, the detailed DETF conclusions about the relative value of different techniques and the importance of combining techniques are unchanged by the 9D analysis."
    },
    "0608/astro-ph0608096_arXiv.txt": {
        "abstract": "Space-based instruments provide new and, in some cases, unique opportunities to search for dark matter.  In particular, if dark matter comprises sterile neutrinos, the x ray detection of their decay line is the most promising strategy for discovery.  Sterile neutrinos with masses in the keV range could solve several long-standing astrophysical puzzles, from supernova asymmetries and the pulsar kicks to star formation, reionization, and baryogenesis. The best current limits on sterile neutrinos come from Chandra and XMM-Newton.  Future advances can be achieved with a high-resolution x-ray spectrometry in space. ",
        "introduction": "There is an overwhelming amount of evidence that most of the matter in the universe is not made of ordinary atoms, but, rather, of new, yet undiscovered particles\\cite{dark_review}.  The evidence for dark matter is based on several independent observations, including cosmic microwave background radiation, gravitational lensing, the galactic rotation curves, and the x-ray observations of clusters. None of the Standard Model particles can be dark matter.  Hence, the identification of dark matter will be a discovery of new physics beyond the Standard Model.  To detect dark matter one must guess its properties, which ultimately determine one's strategy for detection.   One can base one's guesses on compelling theoretical ideas or on some observational clues.  One of the most popular theories for physics beyond the Standard Model is supersymmetry.  A class of supersymmetric extensions of the Standard Model predict dark matter in the form of either the lightest supersymmetric particles\\cite{SUSY_dark_review}, or SUSY Q-balls\\cite{Q_balls}. Another theoretically appealing possibility is dark matter in the form of axions\\cite{axion}.  Axion is a very weakly interacting field which accompanies the Peccei--Quinn solution of the strong CP problem. There are several other dark-matter candidates that are well motivated by theoretical reasoning.  A comprehensive review of possibilities is not our purpose; rather, we will focus on the forms of dark matter which are well-motivated and for which there are new opportunities in space research.  Right-handed or sterile neutrinos can be the dark matter\\cite{dw}\\cdash\\cite{nuMSM}. The existence of such right-handed states is implied by the discovery of the active neutrino masses. Although it is not impossible to explain the neutrino masses otherwise, most models introduce gauge singlet fermions that give the neutrinos their masses via mixing.  If one of these right-handed states has mass in the $\\sim 1-50$~keV range, it can be the dark matter.  Several indirect astrophysical clues support this hypothesis. Indeed, if the sterile neutrinos exist, they can explain the long-standing puzzle of pulsar velocities\\cite{Kusenko:review}.  In addition, the x rays produced in decays of the relic neutrinos could increase the ionization of the primordial gas and can catalyze the formation of molecular hydrogen at redshift as high as 100. Since the molecular hydrogen is an important cooling agent, its increased abundance could cause the early and prompt start formation\\cite{reion,reion1}. Sterile neutrinos can also help the formation of supermassive black holes in the early universe\\cite{mun}. For smaller masses, the sterile neutrinos have a long enough free-streaming length to rectify several reported inconsistencies between the predictions of cold dark matter on small scales and the observations.  The consensus of these indirect observational hints helps make a stronger case for the sterile dark matter. ",
        "conclusions": "Several independent observational hints point to sterile neutrinos with masses in the keV range.  Pulsar velocities can be explained by the emission of such sterile neutrinos from a supernova, because the sterile neutrino emission is anisotropic in the presence of the magnetic field. The x-ray photons from the decays of the sterile neutrinos can ionize the primordial gas and can cause an increase in the fraction of molecular hydrogen, which makes a prompt star formation possible at a relatively high redshift.   The sterile neutrinos can be the dark matter.  In addition, they could have played a role in generating the matter-antimatter asymmetry of the universe. Future observations of x-ray telescopes may be able to discover the relic sterile neutrinos by detecting keV photons from their decays."
    },
    "0608/astro-ph0608575_arXiv.txt": {
        "abstract": "{% The photometric calibration of the Sloan Digital Sky Survey (SDSS) is a multi-step process which involves data from three different telescopes: the 1.0-m telescope at the US Naval Observatory (USNO), Flagstaff Station, Arizona (which was used to establish the SDSS standard star network); the SDSS 0.5-m Photometric Telescope (PT) at the Apache Point Observatory (APO), New Mexico (which calculates nightly extinctions and calibrates secondary patch transfer fields); and the SDSS 2.5-m telescope at APO (which obtains the imaging data for the SDSS proper).\\\\ In this paper, we describe the Monitor Telescope Pipeline, {\\tt MTPIPE}, the software pipeline used in processing the data from the single-CCD telescopes used in the photometric calibration of the SDSS (i.e., the USNO 1.0-m and the PT).  We also describe transformation equations that convert photometry on the USNO-1.0m $u'g'r'i'z'$ system to photometry the SDSS 2.5m $ugriz$ system and the results of various validation tests of the {\\tt MTPIPE} software.  Further, we discuss the semi-automated PT factory, which runs {\\tt MTPIPE} in the day-to-day standard SDSS operations at Fermilab.  Finally, we discuss the use of {\\tt MTPIPE} in current SDSS-related projects, including the Southern $u'g'r'i'z'$ Standard Star project, the $u'g'r'i'z'$ Open Star Clusters project, and the SDSS extension (SDSS-II).} ",
        "introduction": "\\label{sec:intro}  \\begin{figure*} \\includegraphics[angle=-90,scale=0.7]{fg01.ps} \\caption{Locations of the 158 Smith et al.\\ (2002) primary standards (asterisks) and the 64 fields of the southern extension to the primary standard star network (unfilled squares; see \\S~\\ref{sec:otherprojects}).  Circled symbols indicate $u'g'r'i'z'$ standard stars or fields for which there is currently Landolt $UBVR_cI_c$ photometry (Landolt 1973, 1983, 1992) or for which Landolt is currently obtaining $UBVR_cI_c$ photometry (Landolt, in prep.). \\label{allprifieldlocationsRaDec}} \\end{figure*}  \\begin{figure*} \\includegraphics[angle=-90,scale=0.7]{fg02.ps} \\caption{ Locations of the SDSS secondary patches in equatorial coordinates (Aitoff projection): {\\em (black)} Northern SDSS patches, {\\em (red)} Southern SDSS patches, and {\\em (green)} SEGUE patches.  (SEGUE is part of the Sloan extension, or SDSS-II, and is discussed in \\S~\\ref{sec:otherprojects}.) \\label{secpatchlocationsRaDec}} \\end{figure*}     The Sloan Digital Sky Survey (SDSS; York et al.\\ 2000; Stoughton et al.\\ 2002a; Abazajian et al.\\ 2003, 2004, 2005; Adelman-McCarthy et al.\\ 2006) is a modern, CCD-based optical imaging and spectroscopic survey of the Northern Galactic Cap.  Both imaging and spectroscopy are performed using a 2.5m f/5 Ritchey-Chr\\'etien telescope (Gunn et al.\\ 2006).  The imaging camera (Gunn et al. 1998) contains an imaging array of thirty 2048$\\times$2048 SITe CCDs and scans the sky in drift scan mode along great circles in five different filters ($ugriz$).  A complete scan from the western to the eastern borders of the survey area is called a {\\em strip}.  Since there are gaps within the CCD mosaic, it requires two strips, offset by 93\\% of a CCD width, to fill in a complete rectangular area (which is called {\\em stripe}).  The imaging portion of the SDSS has a depth of $r=22.2$ (95\\% detection repeatablity for point sources; Ivezi\\'{c} et al.\\ 2004).  Followup spectrosopy is performed for a variety of targetted galaxies (Eisenstein et al.\\ 2001; Strauss et al. 2002), quasars (Richards et al. 2002), and stars via a pair of 320-fiber multi-object spectrographs.  The SDSS imaging covers several thousand square degrees of sky, and over this region the SDSS photometric calibrations achieve an accuracy of $\\approx$0.02~mag (2\\%) rms in $g$, $r$, and $i$ and $\\approx$0.03~mag (3\\%) in $u$ and $z$ (Ivezi\\'{c} et al.\\ 2004; Adelman-McCarthy et al.\\ 2006). The photometric calibration of the SDSS 2.5m imaging data is a multi-step process which involves the data from three different telescopes and three different data processing pipelines.  The telescopes in question are: \\begin{enumerate}  \\item the US Naval Observatory (USNO) 1.0-m telescope at Flagstaff Station, Arizona, which was used to set up a network of 158 primary standard stars for the $u'g'r'i'z'$ photometric system in the magnitude range $9 \\la r'\\la 14$ (Fig.~\\ref{allprifieldlocationsRaDec}; Fukugita et al.\\ 1996; Smith et al.\\ 2002);  \\item the 0.5-m Photometric Telescope (PT) at Apache Point Observatory (APO), New Mexico, which observes a set of $u'g'r'i'z'$ primaries over a range of airmasses to determine the photometric solution for the night (zeropoints \\& extinctions), and calibrates stars down to $r'\\approx 18$ in transfer fields --- or secondary patches --- that are placed throughout the SDSS survey area (Fig.~\\ref{secpatchlocationsRaDec}); and  \\item the SDSS 2.5-m telescope itself (Gunn et al.\\ 2006), also located at APO, whose imaging camera (Gunn et al.\\ 1998) scans over the secondary patches (Fig.~\\ref{secpatchlocationsRaDec}) during normal course of operations.  Note that the imaging camera saturates under normal operating conditions at $r \\approx 14$, necessitating the use of faint stars in the secondary patches (rather than observing the primary standards directly).  The secondary patches, which are $40\\arcmin \\times 40\\arcmin$ in size, are grouped in sets of four, so that each group spans the full $2.5^{\\circ}$ width of a survey stripe, and each group is spaced at roughly $15^{\\circ}$ (1~hour) intervals along a survey stripe.  [The positions of patches along a stripe have been optimized to avoid bright ($V \\la 8$) stars; hence the slight variability in the spacings between patch groups and between patches within patch groups.]  To calibrate the 2.5m imaging camera photometry, atmospheric extinction measurements are taken from PT observations for the night, and photometric zeropoints are obtained by matching stars in the imaging camera data with those from the secondary patches. This method of calibration permits the 2.5m telescope to spend its time more efficiently, in that the onus of obtaining nightly extinction measurements and of calibrating primary and secondary standard stars is shifted to other, smaller telescopes.  \\end{enumerate}  The three software pipelines are: \\begin{enumerate}  \\item the Monitor Telescope Pipeline ({\\tt MTPIPE})\\footnote{The Monitor Telescope Pipeline is named after the original SDSS 0.6m calibration telescope at APO, called the Monitor Telescope (or MT), which was de-commissioned in Summer 1998 and replaced by the PT, which was commissioned in Spring 1999.  Overlapping this time frame, the USNO-1.0m telescope was used to set up the $u'g'r'i'z'$ standard star network, a task originally planned for the old MT.}, which was used to process the USNO-1.0m data for setting up the $u'g'r'i'z'$ primary standard star network (Smith et al. 2002), and which is currently used to process the PT data for determining the atmospheric extinctions at APO each night and for calibrating the secondary patch fields to the SDSS $ugriz$ system;  \\item the Photometric Pipeline ({\\tt PHOTO}; Lupton et al.\\ 2001; Lupton 2006), which processes the 2.5m imaging camera data, yielding among its many outputs accurate astrometry (Pier et al.\\ 2003) and $ugriz$ instrumental photometry (counts); and  \\item the (New) Final Calibrations Pipeline ({\\tt NFCALIB}; \\S~4.5.3. of Stoughton et al.\\ 2002a), which matches the stars in the PT secondary transfer fields with the 2.5m imaging camera photometry, calculates the resulting photometric zeropoints, and applies them to the 2.5m instrumental magnitudes.  (To deal with 2.5m imaging scans that are observed over a range of airmasses, {\\tt NFCALIB} also makes use of the first-order atmospheric extinction coefficients measured by the PT on the same night as a given imaging scan.  Since most 2.5m imaging scans only cover a small range in airmass, this is typically a small effect.)   \\end{enumerate}  \\begin{figure} \\includegraphics[angle=0,scale=0.35]{fg03.eps} \\caption{A flowchart of {\\tt MTPIPE}. \\label{mtpipeFlowChart1}} \\end{figure}  These pipelines were written with the expressed purpose of meeting the image processing and calibration needs of the SDSS, and are in no way meant to supplant the many excellent general-purpose astronomical image processing and analysis packages like IRAF\\footnote{{\\tt iraf.noao.edu}} (Tody 1986, 1993), MIDAS\\footnote{{\\tt http://www.eso.org/projects/esomidas/}} (Warmels 1985; Grosb{\\o}l 1989), Gypsy\\footnote{{\\tt http://www.astro.rug.nl/$\\sim$gipsy/}} (Shostek \\& Allen 1980; Allen et al. 1992; van der Hulst et al. 1992), SExtractor (Bertin \\& Arnouts 1996), DAOPHOT (Stetson 1987), and DoPHOT (Schechter et al.\\ 1993), to name but a few.  Thus, the SDSS imaging and photometric calibration pipelines have more in common with such special-purpose software packages as the ESO Imaging Survey pipeline (Nonino et al. 1999), the INT Wide Field Camera pipeline (Irwin \\& Lewis 2001), the VLT Survey Telescope pipeline (Grado et al. 2004), and the Liverpool Telescope Gamma Ray Burst pipeline (Guidorzi et al. 2006).  The topic of the current paper is {\\tt MTPIPE}.  In the following sections, we will discuss in turn {\\tt MTPIPE}'s component packages (\\S~\\ref{sec:packages}), validation tests of the {\\tt MTPIPE} software (\\S~\\ref{sec:tests}), the {\\tt MTPIPE}-based semi-automated PT data processing factory (\\S~\\ref{sec:ptfactory}), and other SDSS-related projects using {\\tt MTPIPE} (\\S~\\ref{sec:otherprojects}); in \\S~\\ref{sec:conclusions} we comment on future plans for {\\tt MTPIPE}.  In closing, we note that, although the SDSS magnitudes and colors are shown without primes ($ugriz$), the magnitudes of the primary standard stars are shown with primes ($u'g'r'i'z'$).  This is no accident.  The primed system is defined in the natural system of the USNO-1.0m, its CCD, and its set of SDSS filters.  The SDSS magnitudes, however, are defined in the natural system of the SDSS 2.5m imaging camera and its SDSS filters.  These two systems are very similar and the coefficients of the transformation equations are quite small.  We will describe the differences in these two systems and provide transformation equations in \\S~\\ref{sec:packages}. ",
        "conclusions": ""
    },
    "0608/astro-ph0608213_arXiv.txt": {
        "abstract": "Observations of H$_{2}$CO lines and continuum at 1.3 mm towards Sgr B2(N) and Sgr B2(M) cores were carried out with the SMA. We imaged  H$_{2}$CO line absorption against the continuum cores and the surrounding line emission clumps. The results show that the majority of the dense gas is falling into the major cores where massive stars have been formed. The filaments and clumps of the continuum and gas are detected outside of Sgr B2(N) and Sgr B2(M) cores. Both the spectra and moment analysis show the presence of outflows from Sgr B2(M) cores. The H$_{2}$CO gas in the red-shifted outflow of Sgr B2(M) appears to be excited by a non-LTE process which might be related to the shocks in the outflow. ",
        "introduction": "The giant molecular cloud Sgr B2 is located in the Galactic center ($\\sim$44$^{\\prime}$ from SgrA*). High resolution radio continuum and recombination line observations at centimeter wavelenghts showed  the two most active star forming cores, Sgr B2(N) and (M), hosting clusters of massive stars associated with UCHII regions [1-6]. Previous observations showed that massive star formation is taking place in both Sgr B2(N) and (M), suggesting that the two hot cores are at different evolutionary stages [7-12]. H$_{2}$CO molecule has a simple chemical formation path and has been proven to be a useful probe of physical conditions of star formation regions [13,14]. We observed Sgr B2 at 1.3 mm with the SMA including the lines of H$_{2}$CO. The new results regarding the ongoing star formation in the Galactic center region are reported in this paper. ",
        "conclusions": "H$_{2}$CO lines and continuum at 1.3 mm were observed with SMA. The infalling gas was detected through the red-shifted absorbing gas against the continuum in Sgr B2 compact cores and their associated dusty clumps and filament Z10.24. The gas clumps in emission were also observed in both H$_{2}$CO lines. Based on an LVG model, our analysis suggests that the gas in Sgr B2 complex consists of  cold dense cores and a warm but less dense envelope. The red-shifted outflow in Sgr B2(M) is observed in H$_{2}$CO (3$_{21}$-2$_{20}$) line and appears to be excited by C-shock. \\ack We thank all the SMA staff of SAO and ASIAA for making SMA observations possible."
    },
    "0608/gr-qc0608062_arXiv.txt": {
        "abstract": "The coalescence of massive black holes generates gravitational waves (GWs) that will be measurable by space-based detectors such as LISA to large redshifts.  The spins of a binary's black holes have an important impact on its waveform.  Specifically, geodetic and gravitomagnetic effects cause the spins to precess; this precession then modulates the waveform, adding periodic structure which encodes useful information about the binary's members.  Following pioneering work by Vecchio, we examine the impact upon GW measurements of including these precession-induced modulations in the waveform model. We find that the additional periodicity due to spin precession breaks degeneracies among certain parameters, greatly improving the accuracy with which they may be measured.  In particular, mass measurements are improved tremendously, by one to several orders of magnitude. Localization of the source on the sky is also improved, though not as much --- low redshift systems can be localized to an ellipse which is roughly $\\mbox{a few} \\times 10$ arcminutes in the long direction and a factor of $2-4$ smaller in the short direction.  Though not a drastic improvement relative to analyses which neglect spin precession, even modest gains in source localization will greatly facilitate searches for electromagnetic counterparts to GW events. Determination of distance to the source is likewise improved: We find that relative error in measured luminosity distance is commonly $\\sim 0.2\\%-0.7\\%$ at $z \\sim 1$.  Finally, with the inclusion of precession, we find that the magnitude of the spins themselves can typically be determined for low redshift systems with an accuracy of about $0.1\\%-10 \\%$, depending on the spin value, allowing accurate surveys of mass and spin evolution over cosmic time. ",
        "introduction": "\\label{sec:intro}  \\subsection{Background to this analysis}  Observations have now demonstrated that massive black holes are ubiquitous in the local universe. It appears that all galaxies with central bulges contain black holes whose masses are strongly correlated with the properties of the bulge \\cite{fm00, g00}. Hierarchical structure formation teaches us that these galaxies assembled over cosmic history through the repeated coalescence of the dark matter halos in which they reside {\\cite{hierarchy}}.  Taken together, these suggest that coalescences of massive black holes should be relatively frequent events, especially at high redshift when halo coalescences were common \\cite{vmh03}.  Massive black hole coalescences are extremely strong gravitational wave (GW) sources.  In the relevant mass band --- thousands to millions of solar masses --- the GWs these binaries generate are at low frequency ($f \\sim 10^{-4.5} - 10^{-1}$ Hz) where ground-based GW antennae have poor sensitivity due to geophysical and other terrestrial noise sources.  Measuring GWs from massive black holes requires going into the quiet environment of space.  LISA, the Laser Interferometer Space Antenna, is being designed as a joint NASA-ESA mission to measure GWs in this frequency band; cosmological massive black hole coalescences are among its highest priority targets.  By measuring these GWs, one can infer the properties of the source that generated the waves.  Some particularly important and interesting properties are the masses of the binary's members, their spins, the binary's location on the sky, and its distance from the solar system barycenter.  Measuring a population of coalescence events could thus provide a wealth of data on the cosmological distribution and evolution of black hole masses and spins.  Most analyses of how well binary black hole parameters can be determined by LISA measurements have ignored the impact of spin-induced precession \\cite{c98,h02,bbw05}.  Under such an assumption, subsets of parameters can be highly correlated with each other, increasing the errors in parameter estimation.  One such subset comprises the binary's ``chirp mass'' $\\mathcal{M}$, its reduced mass $\\mu$, and the spin parameters $\\beta$ and $\\sigma$ (which are written out explicitly in Sec. \\ref{sec:intrinsic}).  These four parameters influence the GW phase $\\Phi$.  As discussed in \\cite{cf94,pw95}, the correlation coefficient between $\\mu$ and $\\beta$ is nearly 1.  It is thus difficult to ``detangle'' these parameters from one another in a measurement.  Another such subset consists of a binary's sky position, orientation, and distance.  To see why these parameters are strongly correlated, consider the form of the two polarizations of the strongest quadrupole harmonic of the gravitational waveform:  \\begin{align} h_+(t) &= 2 \\frac{{\\mathcal M}^{5/3}(\\pi f)^{2/3}}{D_L}(1+\\cos^2\\iota) \\cos\\Phi(t) \\, , \\label{eq:hplusintro}\\\\ h_{\\times}(t) &= -4 \\frac{{\\mathcal M}^{5/3}(\\pi f)^{2/3}}{D_L}\\cos \\iota \\sin\\Phi(t) \\, . \\label{eq:hcrossintro} \\end{align} (We work in units with $G = 1 = c$; a convenient conversion factor in this system is $10^6\\,M_\\odot = 4.92\\,{\\rm seconds}$.)  The quantity $\\iota$ is the binary's inclination relative to the line of sight: $\\cos\\iota \\equiv \\mathbf{\\hat{L}} \\cdot \\mathbf{\\hat{n}}$, where $\\mathbf{\\hat{L}}$, the direction of the binary's orbital angular momentum, defines its orientation and $\\mathbf{\\hat{n}}$ is the direction from observer to source.  The quantity $D_L$ is the luminosity distance to the source, and $f(t) \\equiv (1/2\\pi) d\\Phi/dt$.  One does not measure the polarizations $h_+$ and $h_\\times$ directly; rather, one measures a sum $h_M(t)$ in which the two polarizations are weighted by antenna response functions as follows: \\begin{equation} h_M(t) = F_+(\\theta_N,\\phi_N,\\psi_N) h_+(t) + F_\\times(\\theta_N,\\phi_N,\\psi_N) h_\\times(t) \\;. \\label{eq:hM_schem} \\end{equation} (This equation should be taken as schematic; see \\ref{sec:extrinsic} for a more detailed and definitive description.)  The angles $\\theta_N$ and $\\phi_N$ denote the location of the source on the sky in some appropriate coordinate system.  The angle $\\psi_N$, known as the ``polarization angle,'' fixes the orientation of the component of $\\mathbf{\\hat{L}}$ perpendicular to the line of sight.  (In other words, $\\mathbf{\\hat{L}}$ is fixed by $\\iota$ and $\\psi_N$.)  Measuring the phase determines chirp mass with high accuracy; the fractional error in ${\\cal M}$ is often $\\sim 10^{-3} - 10^{-4}$.  As far as amplitude is concerned, the chirp mass can be regarded as measured exactly.  What remains is to determine, from the measured amplitude and the known ${\\cal M}$, the angles $\\theta_N$, $\\phi_N$, $\\psi_N$, $\\iota$, and the distance $D_L$.  As Eqs.\\ \\eqref{eq:hplusintro}, \\eqref{eq:hcrossintro}, and \\eqref{eq:hM_schem} illustrate, these five parameters are strongly correlated.  The motion of LISA around the sun\\footnote{LISA is being designed as a constellation of three spacecraft whose centroid orbits the sun with a period of one year; see {\\tt http://lisa.nasa.gov} for further details.} breaks these degeneracies to some extent --- the angles $\\theta_N$ and $\\phi_N$ appearing in Eq.\\ (\\ref{eq:hM_schem}) can be regarded as best defined in a coordinate system tied to LISA. As the antenna orbits the sun, these angles become effectively time dependent.  The one-year periodicity imposed by this motion makes it possible to detangle these parameters.  Analyses typically find that the position of a merger event at $z \\sim 1$ can be determined, on average, to an ellipse which is $1.5-2$ degrees across in the long direction and $1.5-2$ times smaller in the short direction\\footnote{It is worth bearing in mind that the full moon subtends an angle of about 30 arcminutes.} {\\cite{c98,bbw05,hh05}}.  The distance to such a binary can be determined to $1\\%-2\\%$ accuracy on average (less in some exceptional cases) \\cite{h02,bbw05,hh05}.  \\subsection{Black hole spin and spin precession}  The preceding discussion ignores an important piece of relativistic physics: the precession of each binary member's spin vector due to its interaction with the spacetime in which it moves.  In general relativity, the spacetime of an isolated object can be regarded as having an ``electric piece,'' arising from the object's mass and mass distribution, and a ``magnetic piece,'' arising from the object's mass currents and their distribution\\footnote{This analogy is most apt in the weak field.  In that limit, one can recast the Einstein field equations of general relativity into a form quite similar to Maxwell's equations; see {\\cite{membrane}} for detailed discussion.  Though the analogy does not fit quite so well in strong field regions, it remains accurate enough to be useful.}.  Spin precession consists of a geodetic term, arising from the parallel transport of the spin vector in the gravitoelectric field of the other hole, and Lense-Thirring terms, caused by the gravitomagnetic field of the other hole.  The basic physics of gravitomagnetic precession can be simply understood by analogy with a similar (and closely related) electromagnetic phenomenon --- the precession of a magnetic dipole {\\mbox{\\boldmath$\\mu$}} immersed in an external magnetic field ${\\bf B}$.  An object's spin angular momentum ${\\bf S}$ can be regarded as a gravitational ``magnetic dipole.''  When immersed in a ``gravitomagnetic field,'' one finds that ${\\bf S}$ feels a torque, just as a magnetic dipole {\\mbox{\\boldmath$\\mu$}} experiences a torque when immersed in magnetic field ${\\bf B}$.  In a binary black hole system, the gravitomagnetic field arises from the binary's orbital motion and the spins of its members.  Precession thus includes both spin-orbit (geodetic and orbital gravitomagnetic) and spin-spin effects \\cite{th85}.  (The major goal of the ``Gravity Probe B'' experiment is to measure the effects of geodetic and spin-spin Lense-Thirring precession upon a gyroscope in low Earth orbit {\\cite{gpb}}.)  As the spins precess, they do so in such a way that the {\\it total} angular momentum ${\\mathbf{J}} = {\\mathbf{L}} + {\\mathbf{S}}_1 + {\\mathbf{S}}_2$ is held constant; the orbital angular momentum ${\\mathbf{L}}$ precesses to compensate for changes in ${\\mathbf{S}}_1$ and ${\\mathbf{S}}_2$.  As a consequence, the inclination angle $\\iota$ and polarization angle $\\psi_N$ become time varying (as do certain other quantities appearing in the GW phase function $\\Phi$).  Figure \\ref{fig:Apcomparison} shows the so-called ``polarization amplitude,'' defined in Sec.\\ \\ref{sec:extrinsic}, of the waveform measured by a particular detector.  Without precession, this quantity is modulated by the orbital motion of LISA, helping to provide some information about the binary's sky position.  The polarization amplitude also depends on the angles $\\iota$ and $\\psi_N$, so it undergoes additional modulation when precession is included.  Such precession-imposed time variations quite thoroughly break many of the degeneracies which have been found to limit parameter measurement accuracy in earlier analyses.  It is without a doubt that black holes in nature spin.  Observations are not yet precise enough to indicate the value of typical black hole spins; the evidence to date does, however, seem to indicate that fairly rapid rotation is common.  For example, the existence of jets from active systems seems to require non-negligible black hole spin --- jets appear to be ``launched'' by the shearing of magnetic field lines (supported by the highly conductive, ionized material accreting onto the black hole) by the differential rotation of spacetime around a rotating black hole {\\cite{bz77,m03}}.  Also, observations of highly distorted iron K-$\\alpha$ lines --- a very sharp flourescence feature in the rest frame of the emitting iron ions --- indicate that this emission is coming from very deep within a gravitational potential (at radii less than the Schwarzschild innermost radius $6M$) and is smeared by near luminal relativistic speeds to boot \\cite{rn03}. Though perhaps influenced somewhat by selection effects\\footnote{The systems for which we have constraints on spin are systems which are actively accreting, and, {\\it ipso facto}, those which are most likely to be rapidly spinning \\cite{hb03}.}, these pieces of evidence are strong hints that the black holes which will form the binaries we hope to measure will be strongly influenced by spin.  The only limit in which spin precession can be neglected is that in which the spins of the binary's members are exactly parallel (or antiparallel) to one another and to the orbital angular momentum $\\mathbf{L}$.  Since the target binaries of this analysis are created by galactic merger processes, their members will almost certainly have no preferred alignment --- random spin and orbit orientation is expected to be the rule.  (This expectation is borne out by work {\\cite{spck02}} showing that jets in active galaxies are oriented randomly with respect to the disks of their host galaxies.)  Taking into account spin precession is thus of paramount importance for GW observations of merging black hole systems.  A great deal of work has gone into developing families of model waveforms (``templates'') sufficiently robust to {\\it detect} GWs from spinning and precessing binaries, at least in the context of measurements by ground-based detectors \\cite{bcv03,bcv041,bcv042,bcv05,gkv03,gk03,gikb04}.  The key issue in this case is that the various modulations on the waveform imposed by the binary's precession smear its power over a wider spectral range, making it much more difficult to detect at the (relatively) low signal-to-noise ratios (SNRs) expected for ground-based observations. Not as much work has gone into the complementary problem of {\\it measuring} these waves --- examining the impact precession has upon the precision with which binary properties may be inferred from the waves.  To date, the most complete and important analysis of this type is that of Vecchio {\\cite{v04}}.  Vecchio focuses (for simplicity) on equal mass binaries and only includes the leading ``spin-orbit'' precession term.  This limit is particularly nice as a first analysis of this problem, since it can be treated (largely) analytically (cf.\\ discussion in Sec.\\ III B of Ref.\\ {\\cite{v04}}).  Vecchio's work largely confirms the intuitive expectation discussed above --- the precision with which masses are measured is substantially improved; in particular, the reduced mass of the system can be measured with several orders of magnitude more accuracy. Parameters such as the sky location of the binary and the luminosity distance are also measured more accurately, but only by a factor of 2 -- 10.  \\begin{figure}[htb] \\psfragscanon \\psfrag{xlabel}{$t$ (s)} \\psfrag{ylabel}{$A_{\\mathrm{pol, I}}(t)$} \\includegraphics[scale=0.54]{ApcomparisonNAL.eps} \\includegraphics[scale=0.54]{Apcomparison2NAL.eps} \\caption{These figures depict the ``polarization amplitude'' $A_{\\mathrm{pol}}(t)$ of the signal measured in detector I as a function of time.  The curves are as follows: solid line, $\\chi_1 = \\chi_2 = 0$; dashed line (nearly overlapping the solid line), $\\chi_1 = \\chi_2 = 0.1$; dotted line, $\\chi_1 = \\chi_2 = 0.5$; and dot-dashed line, $\\chi_1 = \\chi_2 = 0.9$.  ($\\chi = S/m^2$ is the dimensionless spin parameter.)  The top figure covers the last two years of inspiral.  The spinless curve has periodicity of one year, corresponding to the motion of LISA around the sun.  Notice that as spin is introduced, the curves become more strongly modulated, with the number of additional oscillations growing as the spin is increased.  By tracking these spin-precession-induced modulations, it becomes possible to better measure parameters like mass and sky position and measure spin for the first time.  The bottom figure shows a close-up of the final months of inspiral.} \\label{fig:Apcomparison} \\end{figure}  \\subsection{This analysis}  Our goal here is to update Vecchio's pioneering analysis by taking the precession equations and the wave phase to the next higher order and by performing a broader parameter survey (including the impact of mass ratio).  By taking the precession equations to higher order, we include ``spin-spin'' effects --- precessional effects due to one black hole's spin interacting with gravitomagnetic fields from the other hole's spin.  By taking the wave phase to higher order, we include, among other terms, a time-dependent spin-spin interaction. Finally, when the mass ratio differs from 1, the geodetic spin-orbit term causes the two spins to precess at different rates, even without the spin-spin corrections.  Including these effects means that we {\\it cannot} model the precession with a simple, analytic rule --- we are forced to integrate the equations of precession numerically as inspiral proceeds, incurring a significant performance cost.  Fortunately, the basic ``engine'' on which this code is based {\\cite{h02}} runs extremely fast, thanks largely to the use of spectral integrators (which, in turn, is thanks to a suggestion by E.\\ Berti {\\cite{bbw05}}), so total run time remains reasonable.  The cost in efficiency due to the inclusion of higher-order effects is offset by the more complete description of the signal they provide. An important consequence is that it now becomes possible from GW measurements to determine the spin of each member of the binary.  With Vecchio's approximations, only three components of the black holes' vector spins can be determined --- enough to constrain, but not determine, their spin magnitudes.  Our more general approach allows us to measure all six vector spin components.  To our knowledge, this is the first analysis indicating how well spin can be measured from merging comparable mass binary systems.  (As Barack and Cutler have shown {\\cite{bc04}}, spin is very well determined by measurements of GWs from {\\it extreme mass ratio} binaries --- those in which the system's mass ratio $m_2/m_1 \\lesssim 10^{-4}$ or so.)  Our error estimates are computed using the maximum likelihood formalism first introduced in the context of GW measurements by Finn {\\cite{f92}}.  A potential worry is that we are using a Gaussian approximation to the likelihood function.  This approximation is very convenient since it allows us to directly compute a Fisher information matrix.  Its inverse is the covariance matrix, which directly encodes the estimated 1-$\\sigma$ errors in measured parameters, as well as correlations among different parameters.  The Gaussian approximation is known to be accurate when the SNR is ``high enough'' \\cite{f92,cf94}.  Unfortunately, it is not particularly obvious what ``high enough'' really means.  In our case, we are estimating measurement errors on 15 parameters\\footnote{Two masses; 2 angles specifying the initial orientation of the binary's orbit; 4 angles specifying the initial orientation of the spins; 2 spin magnitudes; the time at which coalescence occurs; the phase at coalescence; 2 angles specifying the binary's position on the sky; and the distance to the binary.}  --- a rather fearsome number to fit.  The Gaussian approximation almost certainly {\\it underestimates} measurement error, since it assumes the likelihood function is completely determined by its curvature in the vicinity of a maximum, missing the possibility of a long tail to large error.  We thus fear that our estimates are likely to be optimistic, especially for events with relatively small SNR.  It would be quite salubrious to ``spot check'' a few cases by directly computing the likelihood function in a few important corners of parameter space and comparing to the Gaussian predictions.  This would both quantify the degree to which our calculations are too optimistic and help to determine how large SNR must be for this approximation to be reliable.  In addition to concerns about the Gaussian approximation, it must be noted that the waveform family we use for our analysis is somewhat limited.  We use a post-Newtonian description of the GWs from these binaries.  Since our analysis requires us to follow these binaries deep into the strong field where the usual post-Newtonian expansion is likely to be somewhat unreliable, it is likely that we are introducing some systematic error.  In particular, the equations of spin precession that we are using are only given to the leading order needed to see spin-orbit and spin-spin precession effects \\cite{acst94}.  Higher spin-orbit corrections to the equations of motion and precession have recently been derived {\\cite{fbb06}}, as have their impact on the the waves' phasing {\\cite{bbf06}}.  Another analysis {\\cite{pr06}} has worked out higher-order spin-spin corrections to the post-Newtonian metric, from which it would not be too difficult to work out equations of motion and precession and then the modification to the waves' phase.  It would be interesting to see what effect the higher-order corrections have on these results.  Finally, it should be noted that the frequency domain expression of the signal which we are using is derived formally using a ``stationary phase'' approximation.  This approximation is based on the idea that the binary's orbital frequency is changing ``slowly.''  The orbital frequency is thus well-defined over ``short'' time scales. Quantitatively, this amounts to a requirement that the time scale on which radiation reaction changes the orbital frequency, $T_{\\rm insp}$, be much longer than an orbital period, $T_{\\rm orb}$. Precession introduces a new time scale, $T_{\\rm prec}$, the time it takes for the angular momentum vectors to significantly change their orientations.  For the stationary phase approximation to be accurate, we must in addition require $T_{\\rm prec} \\gg T_{\\rm orb}$, a somewhat more stringent requirement than $T_{\\rm insp} \\gg T_{\\rm orb}$.  No doubt, a certain amount of error is introduced due to the breakdown of this condition late in the inspiral.  Thus, the results which we present here should be taken as {\\it indicative} of how well LISA is likely to be able to measure the parameters of massive black hole binaries, but cannot be considered definitive.  We are confident however that the {\\it improvement} in measurement accuracy obtained by taking spin precession into account is robust.  Specifically, we see that errors in masses are reduced dramatically, from one to several orders of magnitude.  Errors in sky position and distance are also reduced, but by a smaller factor.  Such improvement may nonetheless critically improve the ability of LISA to interface with electromagnetic observatories {\\cite{an02,mp05,kfhm06}}.  Finally, the added information in the precession signal allows us to measure the spins of the holes.  These improvements due to precession will certainly survive and play an important role even in an analysis which addresses the caveats we list above.  \\subsection{Organization of this paper}  The remainder of the paper is organized as follows.  In Sec.\\ \\ref{sec:inspiral}, we discuss the gravitational waveform generated by binary black hole coalescence, focusing upon the slow, adiabatic inspiral.  Section \\ref{sec:intrinsic} describes the ``intrinsic'' waveform produced by the motion of the orbiting black holes as given in the ``restricted post-Newtonian expansion'' of general relativity. Section {\\ref{sec:precession}} then describes the post-Newtonian precession equations which we use to model the evolution of the spins of a binary's members, as well as how those precessions influence the waveform.  Finally, in Sec.\\ {\\ref{sec:extrinsic}} we describe ``extrinsic'' effects which enter the measured waveform through its measurement by the LISA constellation.  In Sec.\\ {\\ref{sec:meas}} we summarize our parameter estimation formalism; this section will be largely review to readers familiar with the literature on GW measurements.  Section {\\ref{sec:meastheory}} first summarizes the maximum likelihood formalism we use to estimate measurement errors.  In Sec.\\ {\\ref{sec:lisanoise}}, we then describe the up-to-date model for the noise which we expect to accompany LISA measurements.  Section {\\ref{sec:results}} presents our results.  After describing some critical procedural issues in the setup of our calculations in Sec.\\ {\\ref{sec:procedure}}, we summarize our results for parameters intrinsic to the binary (particularly masses and spins) in Sec.\\ {\\ref{sec:resI}} and for extrinsic parameters (particularly sky position and luminosity distance) in Sec.\\ {\\ref{sec:resII}}.  In both cases, we compare, when appropriate, to results from a code which does not incorporate spin-precession physics.  (This code was originally developed for the analysis presented in Ref.\\ {\\cite{h02}}.)  The general rule of thumb we find is that the accuracy with which masses can be determined is improved by about one to several orders of magnitude when precession physics is taken into account.  In addition, we find that for low redshift ($z \\sim 1$) binaries LISA should be able to determine the spins of the constituent black holes with a relative precision of $0.1\\% - 10\\%$, depending (rather strongly) on the spin value.  Likewise, we find improvement in the measurement accuracy of extrinsic parameters, though not quite as striking --- half an order of magnitude improvement in source localization and distance determination is a good, rough rule of thumb.  An important consequence of these improvements is that LISA should be able to localize low redshift binaries --- using GW measurements alone --- to an elliptical ``pixel'' that is perhaps $\\mbox{a few} \\times 10$ arcminutes across in its widest direction and a factor of $2-4$ smaller along its minor axis.  For higher redshift binaries ($z \\sim 3 - 5$), this pixel is several times larger, perhaps a few degrees in the long direction and tens of arcminutes to a degree or two in the narrow one.  These results suggest that it should not be too arduous a task to search for electromagnetic counterparts to a merging binary black hole's GW signal {\\cite{an02,mp05,kfhm06}} --- particularly at low redshift, these pixel sizes are comparable to the field of view of planned large scale surveys.  A concluding and summarizing discussion is given in Sec.\\ {\\ref{sec:disc}}.  Along with summarizing our major results and findings, we discuss future work which could allow us to quantitatively assess the consequences of some of the simplifying assumptions we have made.  At several points in this analysis, we need to convert between a source's redshift $z$ and luminosity distance $D_L$.  To make this conversion, we assume a flat cosmology ($\\Omega_{\\rm total} = 1$) with contributions from matter ($\\Omega_M = 0.25$) and from a cosmological constant (equation of state parameter $w = -1$, $\\Omega_\\Lambda = 0.75$).  We also choose a Hubble constant $H_0 = 75 \\, \\mathrm{km\\ s^{-1}\\ Mpc^{-1}}$.  These choices are in concordance with the latest fits presented by the WMAP team in their three-year analysis of the cosmic microwave background {\\cite{setal06}}.  The luminosity distance as a function of redshift is then given by  \\begin{equation} D_L(z) = \\frac{(1+z)c}{H_0}\\int_0^z \\frac{dz^\\prime}{\\sqrt{\\Omega_M(1+z')^3 + \\Omega_\\Lambda}} \\, . \\label{eq:Dzrelation} \\end{equation} ",
        "conclusions": "\\label{sec:disc}  The general relativistic precession of black holes in binary systems can have a strong influence on the binary's dynamics {\\cite{acst94,s04}} and thus upon the GWs that it generates.  It has been known for some time that it will be necessary to take these dynamics into account in order to detect these black holes in noisy detector data {\\cite{bcv03,bcv041,bcv042,bcv05,gkv03,gk03,gikb04}}; clearly, taking these dynamics into account will be just as (if not more) important for the complementary problem of determining the parameters which characterize a detected system.  Vecchio {\\cite{v04}} first demonstrated that, by taking into account precession physics, quite a few near degeneracies among binary source parameters can be broken, making our estimates for how accurately they can be determined more optimistic.  This analysis largely confirms and extends Vecchio's pioneering work.  By taking the equations of motion to higher order to include spin-spin couplings, and by surveying measurement accuracy as a function of mass ratio, we have found that the improvement noted by Vecchio holds rather broadly.  The degeneracy breaking due to precession physics is a rather robust phenomenon.  Two conclusions from this work are particularly important with regard to the astrophysical reach of future LISA measurements.  The first is that modeling spin-precession physics makes it possible to determine the magnitudes of the spins of the black holes which constitute the binary.  If the spins are rapid, they can be measured quite accurately (as good as $0.1 \\%$ accuracy for high spin, low redshift systems) due to the strong modulation imposed on the signal by their interaction. Coupled with the fact that the black hole masses can likewise be measured with good precision, this suggests that LISA will be a valuable tool for tracking the evolution of both mass and spin over cosmic time.  Such observations could provide a direct window into the growth of cosmological structures.  Measuring spin may also make it possible to indirectly test the black hole area theorem {\\cite{h71}}. The requirement that black hole area can only grow implies a consistency relation between the initial and final masses and spins. By measuring the initial masses and spins through the inspiral, and the mass and spin of the merged remnant hole through the ringdown waves {\\cite{dkkfgl04,bcw06}}, we can check this consistency relation in a manner analogous to the mass loss test proposed in {\\cite{hm05}}. We intend to investigate whether this test is feasible in future work.  Second, we confirm Vecchio's result that precession breaks degeneracies between the angles which determine a binary's orientation and its position on the sky, improving the accuracy with which sky position can be fixed using GWs alone.  At low redshift ($z \\sim 1$), we find that sources can be localized to within an ellipse whose major axis is typically $\\mbox{a few}\\times 10$ arcminutes across and whose minor axis is typically a factor $\\sim 2-4$ smaller.  This is small enough that searching the GW pixel for an electromagnetic counterpart to the merger event should not be too arduous a task \\cite{kfhm06}. For mergers at higher redshift, the waves weaken and the source is not so well localized.  The field which would need to be searched for sources at $z \\sim 3 - 5$ is typically a few degrees across in the long axis and tens of arcminutes to a degree or two in the short direction --- a rather more difficult challenge, but not hopeless.  We intend to more thoroughly investigate the nature of localization with spin precession, including how the pixel evolves with observation time up to final merger, in future work.  As mentioned in the Introduction, our analysis makes many assumptions and approximations which are likely to affect our results; a goal of future work will be to lift these approximations.  One major concern is the Gaussian approximation we have taken to the likelihood function.  As already discussed, this approximation is known to be good when the SNR is ``large'' {\\cite{f92,cf94}}; however, it is not apparent what large really means, particularly given that we are fitting for 15 parameters.  Lifting this simplifying approximation can be done by simply computing the likelihood function (\\ref{eq:signalprob}) directly and examining how well parameters are thereby determined.  In the context of GW measurements, Markov Chain-Monte Carlo techniques have been investigated and found to be very useful {\\cite{cm01,cdwm04,umdvwc04}}; the application of these techniques to LISA measurement problems is now being rather actively investigated {\\cite{cc05,cp06_1,cp06_2}}.  Because we have used the Gaussian approximation (among other simplifications taken in this analysis), we cannot claim that this analysis gives a definitive statement about the accuracy with which LISA could measure binary black hole source parameters.  However, it is certainly {\\it indicative} of the accuracy which we expect LISA to achieve.  In particular, we are confident that the trends we have seen as parameters are varied (e.g., masses, redshift, spin magnitude) are robust.  Most importantly, it is very clear that the influence of spin-induced precession upon the measured waveform allows parameters to be measured to greater accuracy than before.\\\\"
    },
    "0608/astro-ph0608355_arXiv.txt": {
        "abstract": "We investigate the interaction between a giant planet and a viscous circumstellar disk by means of high-resolution, two-dimensional hydrodynamical simulations. We consider planet masses that range from $1$ to $3$ Jupiter masses ($\\MJup$) and initial orbital eccentricities that range from $0$ to $0.4$. We find that a planet can cause eccentricity growth in a disk region adjacent to the planet's orbit, even if the planet's orbit is circular. Disk-planet interactions lead to growth in a planet's orbital eccentricity. The orbital eccentricities of a $2\\,\\MJup$ and a $3\\,\\MJup$ planet increase from $0$ to $0.11$ within about $3000$ orbits. Over a similar time period, the orbital eccentricity of a $1\\,\\MJup$ planet grows from $0$ to $0.02$. For a case of a $1\\,\\MJup$ planet with an initial eccentricity of $0.01$, the orbital eccentricity grows to $0.09$ over $4000$ orbits. Radial migration is directed inwards, but slows considerably as a planet's orbit becomes eccentric. If  a planet's orbital eccentricity becomes sufficiently large, $e \\ga 0.2$, migration can reverse and so be directed outwards. The accretion rate towards a planet depends on both the disk and the planet orbital eccentricity and is pulsed over the orbital period. Planet mass growth rates increase with planet orbital eccentricity. For $e \\sim 0.2$ the mass growth rate of a planet increases by $\\sim 30\\%$ above the value for $e=0$.  For $e \\ga 0.1$, most of the accretion within the planet's Roche lobe  occurs when the planet is near the apocenter. Similar accretion modulation occurs for flow at the inner disk boundary which represents accretion toward the star. ",
        "introduction": "\\symbolfootnote[0]{\\hspace*{-0.5em}$^\\dag$% To appear in % \\textsc{The Astrophysical Journal} % (v652 n2 December 1, 2006 issue). Also available as ApJ preprint doi: % \\texttt{10.1086/508451}. } \\label{sec:introduction} \\footnotetext[1]{NASA Postdoctoral Fellow.} \\setcounter{footnote}{1} \\thispagestyle{empty} A striking property of the extrasolar planets known to date is the range of their orbital eccentricities, which is far wider than that of planets in the solar system.\\footnote{See the \\textit{California \\& Carnegie Planet Search} at \\url{http://exoplanets.org} and the \\textit{Extrasolar Planets Encyclopaedia} at \\url{http://exoplanet.eu}. } Eccentricities are typically $\\sim 0.2$--$0.3$, but very low and high values are also found \\citep[e.g.,][]{marcy2005}. A variety of explanations have been proposed to explain the eccentricities. The high eccentricity cases may owe their explanation to the presence of a distant binary companion \\citep{wu2003,takeda2005}. In fact, the large eccentricity ($e=0.93$) of HD~80606b could be the result of a three-body interaction process, known as ``Kozai cycle'' \\citep{wu2003}. Moreover, numerical experiments show that the observed distribution of orbital eccentricities for $e\\gtrsim 0.6$ can be reproduced by assuming the action of a Kozai-type perturbation \\citep{takeda2005}.  For more typical eccentricities, other processes are likely to be at work. Planet-planet interactions involving scattering on dynamical timescales is a possibility \\citep{rasio1996}. However, numerical experiments indicate that interactions between equal mass planets would produce more isolated planets on low-eccentricity orbits than those observed \\citep{ford2001}. Secular interactions between planets is another possible means of eccentricity excitation \\citep{juric2005}. This mechanism assumes that the planets have appropriate initial separations for evolution  to occur on secular timescales ($\\sim 10^{10}$ years).  Disk-planet interactions can also give rise to planetary eccentricities \\citep{gt1980,pawel1992,goldreich2003}. The evolution of orbital eccentricity depends on a delicate balance between Lindblad and corotation resonances. If the planet is massive enough to clear a gap, Lindblad resonances promote eccentricity growth, while corotation resonances damp eccentricity. If the corotation resonances operate at maximal efficiency, they dominate over Lindblad resonances by a slight margin, and the result is eccentricity damping \\citep[][]{gt1980}. Two mechanisms have been proposed to weaken the effects of corotation resonances and thereby provide eccentricity growth. The first mechanism relies on a large enough gap about a planet's orbit to exclude the corotation resonances from the disk, while leaving a remaining Lindblad resonance, the 1:3 outer resonance. This mechanism requires a massive enough companion and/or low enough viscosity to open a wide enough gap. Such a mechanism has been demonstrated for binary stars \\citep{pawel1991}. The second mechanism relies on the delicate nature of corotational resonances in their ability to weaken (saturate) when the disk viscosity is sufficiently small and the planet eccentricity sufficiently large \\citep[][]{ogilvie2003,goldreich2003}. A small finite initial eccentricity, typically $e\\sim 0.01$, is required for the eccentricity to grow by this mechanism. Moreover, a companion object on a circular orbit can drive eccentricity in a circumstellar disk. This process is believed to occur in disks involving $10$--$20$ Jupiter-mass ($\\MJup$) companions \\citep{papa2001} and binary stars \\citep{lubow1991a,lubow1991b}. The disk transfers eccentricity to the planet's orbit by exchanging energy and angular momentum. The disk-planet system undergoes a coupled eccentricity evolution.  One goal of this paper is to determine whether the eccentricity growth by disk-planet interactions occurs. Simulations by \\citet{papa2001} suggested that orbital eccentricity is excited when the mass of the embedded body is larger than about $10$--$20$ Jupiter-masses, while lower mass companions experienced eccentricity damping. By applying higher resolution and simulating over longer timescales, we aim to see whether eccentricity growth can occur for lower mass, planetary mass, companions. In addition, we are interested in the effects of eccentricity on planet migration.  Accretion of gas onto a planet is likely affected by the planet's orbit eccentricity. Circumbinary disks surrounding  eccentric orbit binaries undergo pulsed accretion on  orbital timescales \\citep{pawel1996}. A similar process could occur with eccentric orbit planets. The mean accretion rate could also be modified by the orbital eccentricity. Another goal of this paper is to investigate this accretion process.  \\pagestyle{myheadings} \\markboth{\\hfill Evolution of Giant Planets in Eccentric Disks \\hfill}% {\\hfill \\textsc{G. D'Angelo, S. Lubow, \\& M. Bate} \\hfill} We apply high-resolution hydrodynamical simulations to investigate disk-planet interactions over several thousand orbital periods. In \\refSect{sec:model_description} we describe the physical model. The numerical aspects of the calculations are detailed in \\refSect{sec:numerical_issues}. Results on the growth of the disk eccentricity for a fixed planet orbit are presented in \\refSect{sec:disk_eccentricity}. Similar results on the growth of disk eccentricity were recently reported by \\citet{kley2006}. We describe results on the planet's orbital evolution in \\refSect{sec:orbital_evolution}. Results on the pulsed accretion are described in \\refSect{sec:planet_accretion}. Finally, our conclusions are given in \\refSect{sec:conclusions} ",
        "conclusions": "\\label{sec:conclusions} We simulated the orbital evolution of circular and eccentric orbit giant planets embedded in circumstellar disks.  The disks were  analyzed using a two-dimensional hydrodynamics code that utilizes nested grids to achieve high resolution in a large region ($2\\,a\\times 2\\pi/3$) around the planet. The disks were modeled as an $\\alpha$-disk and a few values of $\\alpha$ were considered. We investigated planet masses of $1\\,\\MJup$, $2\\,\\MJup$, and $3\\,\\MJup$ and initial orbital eccentricities that ranged from $0$ to $0.4$.  Disk gaps become broader and shallower as the planet eccentricity increases (see \\refFgp{fig:globalsigma} and \\ref{fig:avsigma}). The density near the orbit of the planet is very small compared with the density in the disk for all eccentricities considered. A planet on a fixed circular orbit can cause an initially circular disk to become eccentric (see \\refFgp{fig:rcmvstime}). The disk eccentricity is suppressed at lower planet masses ($\\Mp\\lesssim 1\\,\\MJup$) and higher disk viscosities ($\\alpha\\gtrsim 0.01$), as also found by \\citet{kley2006} and by \\citet{papa2001} at higher planet masses. We attribute the eccentricity growth to a tidal instability associated with a series of eccentric outer Lindblad resonances in the inner parts of the outer disk (\\refFgp{fig:lam}). The same type of instability, involving an inner disk, is thought to be responsible for the superhump phenomena in binary star systems \\citep{lubow1991a,osaki2003}.  The simulations indicate that planet eccentricity can grow, as a consequence of disk-planet interactions (\\refFgp{fig:et_mp23} and \\ref{fig:et_mp1vse}). The growth is stronger in the $2\\,\\MJup$ and $3 \\,\\MJup$ cases than for $1\\,\\MJup$, and for lower disk viscosity ($\\alpha\\lesssim 4\\times10^{-3}$). Planet eccentricities of $\\sim 0.1$ were found in the simulations over the course of a few thousand orbits for $2 \\,\\MJup$ and $3 \\,\\MJup$  planets. A similar eccentricity growth is obtained for a $1 \\,\\MJup$ planet in a disk with viscosity $\\alpha\\approx 10^{-3}$. The planet and disk both acquire eccentricity as they interact, which may lead to complicated time-dependent behavior of their eccentricities. The planet eccentricity growth is likely aided by the disk eccentricity growth. The results suggest that the eccentric growth found for $\\sim 10\\,\\MJup$ planets by \\citet{papa2001} also occurs for lower planet masses. The higher resolution achieved by our calculations may be playing a role in obtaining this growth.  For circular orbit planets, migration occurs on roughly the local viscous timescale, as expected for Type~II migration. However, it is slowed for eccentric orbit planets. This result appears for several configurations with either dynamically determined (\\refFgp{fig:at_mp123} and \\ref{fig:eat_mp3vsnu}) or imposed planet eccentricities  (\\refFgp{fig:at_mpvse}). For a $2\\,\\MJup$ case, even migration reversal (outward migration) is found for a dynamically determined eccentricity (\\refFgp{fig:at_mp123}). Migration slowing or reversal would have important consequences for the planet formation process. The cause is not yet clear. It may involve torques from outer disk \\citep{papa2002} or instead from the coorbital region. Some preliminary evidence suggests the latter.  Mass accretion both within a planet's Roche lobe and through a gap can be strongly modulated with orbital phase for eccentric orbit planets or eccentric disks (\\refFgp{fig:avmpdot}, \\ref{fig:density_zoom}, and \\ref{fig:avmsdot}). The modulation was largest for planet eccentricity $e \\simeq 0.2$. This pulsating accretion is similar to what is found for eccentric orbit binary stars embedded in a circumbinary disk \\citep{pawel1996}, although the phasing is different. Both disk and planet eccentricity also lead to enhanced accretion onto the planet. This enhancement likely helps planets achieve higher masses.  The simulations lend support to the idea that disk-planet interactions cause planet eccentricity growth, along the lines of \\citet{goldreich2003}. The simulations suggest that planet eccentricities are easier to achieve for higher mass planets ($\\Mp\\gtrsim 2\\,\\MJup$). Our results are subject to the usual limitations in approximate initial conditions, simulation time, radial range for coverage of the disk (likely resulting in the lack of an inner disk), the $\\alpha$-disk model, and the use of various numerical devices. We also neglected disk self-gravity, which may affect migration especially for higher mass disks \\citep{anelson2003a}.  However, it is not clear that typical extra-solar planet eccentricities of $0.2$--$0.3$ can be achieved through disk-planet interactions. The eccentricity growth at later times shows indications of slowing and possibly stalling for $e \\la 0.15$ (see Fig~\\ref{fig:eat_mp3vsnu}.) Perhaps higher eccentricities can be achieved for disks with different properties (e.g., lower viscosity and smaller disk's aspect ratio). Eccentricity may be limited by damping due to high order eccentric inner Lindblad resonances that lie outside a planet's orbit. Simulations of eccentric orbit binary star systems suggest that little eccentricity growth occurs for $e \\ga 0.5$ \\citep{lubow1992b}. Although the simulated planets do not achieve orbital eccentricities in excess of $0.15$ over the duration of the simulated evolution (for configurations that start from circular orbits), the simulation times correspond to less than $10^{5}$ years. Migration slowing and reversal may permit the planets to achieve higher eccentricities on longer timescales while avoiding orbit decay into the disk center/host star.  We have not yet investigated eccentricity evolution of sub-Jupiter mass planets. They may also provide important constraints. Other simulations suggest that disks with standard viscosity have only mild gaps for smaller planet masses of $\\Mp \\la 0.1\\,\\MJup$ \\citep[e.g.,][]{gennaro2003b,bate2003}. Under those conditions, disk-planet interactions likely lead to eccentricity damping, due to the dominance of the coorbital Lindblad resonance \\citep{ward1986,pawel1993}. The observational determination of eccentricities for small mass planets would help constrain these models. The planet around HD~49674 is close to this regime. It has a minimum mass of $0.11\\,\\MJup$ and a best-fit eccentricity of $0.29$ (P.~Butler, private communication). Since it is close to the central star (the period is $4.9$ days), it is possible that the eccentricity evolution could be more complicated, especially if it became trapped in a central disk hole. Examples of isolated planets like this, but at longer periods would provide useful constraints."
    },
    "0608/astro-ph0608449_arXiv.txt": {
        "abstract": "We have carried out a detailed study of symbiotic stars with white dwarf accretors by means of a population synthesis code. We estimate the total number of symbiotic stars with white dwarf accretors in the Galaxy  as 1,200 -- 15,000. This range is compatible with observational estimates. Two crucial physical parameters that define the birthrate and number of symbiotic stars are the efficiency of accretion by white dwarfs (which greatly depends on the separation of components after common envelope stage and stellar wind velocity) and the mass of the hydrogen layer which the white dwarf can accumulate prior to the hydrogen ignition. The theoretical estimate of the Galactic occurrence rate of symbiotic novae ranges from about 1.3 to about 13.5 yr$^{-1}$, out of which weak symbiotic novae comprise about 0.5 to 6.0 yr$^{-1}$, depending on the model assumptions. We simulate the distributions of symbiotic stars over orbital periods, masses of components, mass-loss rates of cool components, mass-accretion rates of hot components and  luminosity of components. Agreement with observations is reasonable. ",
        "introduction": "Symbiotic stars (SySs) are an inhomogeneous group of variable stars with composite spectra. Their spectacular spectral and photometric variability is a very important and interesting phenomenon. The spectra of SySs suggest that a three-component system consists of a binary system in which an evolved giant transfers matter to a much hotter compact companion by means of stellar wind and an HII region \\citep{l1,l2,l3}.  The cool component is a red giant (RG) which is a first giant branch (FGB) or an asymptotic giant branch (AGB) star. In the majority of SySs the hot component is, most probably, a white dwarf (WD), a subdwarf or an accreting low-mass main-sequence (MS) star \\citep{b50,b18,b32,b66,b16}. The variability of SySs may be due to the thermonuclear runaways on the surface of an accreting WD \\citep{b50,b60,b16} or to the variations in the accretion rate onto the hot component \\citep{d86,l18,l13}. Recent reviews of the properties of SySs can be found in \\citet{b34} and \\citet{b29}. The nova-like eruptions of SySs were reviewed by \\citet{b28}.   In astrophysics today, the interaction between components in the binaries is of special interest, for which SySs offer an exciting laboratory \\citep{b19}. For instance, the nature of components' interaction and of the activity of WD accretors are still controversial \\citep{b34}; mass loss and stellar wind (including the terminal velocity and acceleration mechanism) from a cool giant are not clear; SySs have been frequently discussed as Type Ia supernovae progenitors \\citep{b50,b31,b20,b60,b16,b5}, but their contribution to the rate of these events remains uncertain.  Theoretical studies on the formation and evolution of the SySs have been published, e.g., by YLTK, \\citet{b6,b16,b12}. Their investigations reproduced successfully many observed properties of these objects. However, in the recent years, a number of papers have appeared that provide new insights in the different aspects of stellar evolution relevant to the SySs. \\citet{b44} and \\citet{b57} did a more detailed study of Novae models; \\citet{b36} and \\citet{b37} put forward an alternative algorithm for angular momentum loss during the common-envelope stages of evolution; Winters and his collaborators \\citep{b53,b54,b55} studied the hydrodynamical structure of the stellar wind around AGB stars with low mass-loss rate and low wind outflow velocity. New observational catalogue and an analysis of SySs have been given in \\citet{b61} and \\citet{b29}. It is time for a new study of the SySs.  In the present work we model the subpopulation of SySs with WDs as the hot components and FGB or AGB stars as the cool components. We study evolutionary channels leading to the formation of symbiotic systems in which He, CO or ONe WDs accrete hydrogen-rich matter from the stellar wind of FGB or AGB stars. We obtain SySs' birthrate, lifetime and number in the Galaxy, and some potentially observable parameters of SySs, such as orbital periods, masses of the components, their luminosities and  mass-loss rates. Special attention is paid to the dependence of population model on the parameters entering population synthesis. This paper is the first one of a series of papers on SySs. It offers a basis for further studies of symbiotic stars, such as D-SySs which are SySs with thick dust shells (paper in preparation).   In $\\S$ 2 we present  our assumptions and describe some details of the modeling algorithm. In $\\S$ 3 we discuss the main results and the effects of different parameters. In $\\S$ 4 the main conclusions are given. ",
        "conclusions": "We performed a detailed study of the formation of symbiotic systems with WDs as hot components, employing the population synthesis approach to the evolution of binaries. Several important conclusions can be drawn:  1. The number of nuclear symbiotic stars in the Galaxy may range from  about 1,200 to 15,000. These numbers are compatible with observational estimates. The model birthrate of SySs in the Galaxy is from 0.035 to 0.131 yr$^{-1}$.   2. The estimated occurrence rate of symbiotic Novae is between 1.3 yr$^{-1}$ and 13.5 yr$^{-1}$; weak SyNe contribute to this rate from about 0.5 to 6.0 events per yr. This estimate greatly depends on the critical ignition mass of the hydrogen shell $\\Delta M^{{\\rm crit}}_{{\\rm WD}}$ and accretion efficiency. It may be compared to the estimate of the Galactic occurrence rate of classical Novae $41\\pm20$\\, yr$^{-1}$ \\citep{hbf97}. It looks likely that the models with high rates of symbiotic novae strongly overestimate this rate. This can mean that (i) the relative fraction of massive accretors is overestimated because of too steep initial-final mass relation; (ii) the values of hydrogen-ignition mass $\\Delta M^{{\\rm crit}}_{{\\rm WD}}$ implied by us  may be underestimated; (iii) mass loss by stellar winds from accreting dwarfs is underestimated or (iv) accretion efficiency is overestimated. The latter issue has to be resolved by gas-dynamical calculations of mass-flows in symbiotic systems. Note also, that tripling of $\\Delta M^{{\\rm crit}}_{{\\rm WD}}$ in case 11 brings occurence rate of SyNe to about 1 per yr; having in mind the paucity of observed SyNe this low number may be preferred. Ignition masses given by Eq.~(\\ref{eq:yunm}) are typically higher than the masses computed by \\citet{b44,b57}. This may mean that  the latter studies also strongly underestimate $\\Delta M^{{\\rm crit}}_{{\\rm WD}}$.  There  may exist two varieties of symbiotic stars --- those with short outbursts and those with long ones.  3. The evolution of symbiotic stars is unlikely to lead to SNe Ia via accumulation of Chandrasekhar-mass or to the edge-lit detonations. Helium novae may occur in symbiotic systems if the energy is efficiently dissipated at the base of accumulated He-layer.  4. The results of the modeling depend on the assumed common-envelopes formalism. Within $\\alpha$-formalism, variations of combined parameter $\\alpha_{\\rm ce}\\lambda$ do not strongly affect model population, since it is dominated by wide systems and systems that experience stable RLOF. Increase of $\\alpha_{\\rm ce}\\lambda$ shifts the systems produced via channel I into the ranges of parameters of descendants of channels II and III.  The main effect of applying $\\gamma$-formalism is the dominance of systems formed through common envelopes. Comparison with observations, whenever possible (Table 2, Figs. 4 -- 11), suggests that using $\\gamma$-formalism it is easier to explain variety of parameters of observed SySs. This especially concerns existence of short-period systems and the tendency of SySs for harbouring relatively low-mass accretors. There is apparent excess of Novae in $\\gamma$-formalism models. However, this may be also an effect of the underestimate of critical ignition masses or the overestimate of the number of systems with massive accretors.    5. The wind velocity law affects the occurrence rate of SyNe and the number of SySs within a factor of 2. Observational data is still not sufficient to provide real restrictions on the models.   6. To summarize, the models for the population of SySs mostly depend on the assumed formalism for the common envelope evolution and critical mass for hydrogen ignition $\\Delta M_{\\rm {WD}}^{{\\rm crit}}$. These factors have the strongest effect upon the occurrence rate of SyNe and the total number of SySs, introducing an uncertainty up to a factor of about 4."
    },
    "0608/astro-ph0608480_arXiv.txt": {
        "abstract": "{We present the first measurement of the mass function of galaxy clusters based directly on cluster masses derived from observations of weak gravitational lensing. To investigate the degree of sample incompleteness resulting from the X-ray based selection of the target clusters, we use a sample of 50 clusters with weak lensing mass measurements to empirically determine the relation between lensing mass and X-ray luminosity and the scatter about this relation. We use a complete, volume-limited sub-sample of 35 X-ray luminous clusters of galaxies at $0.15 < z < 0.3$ to constrain the abundance of very massive ($M \\gtrsim 10^{15} h^{-1} M_{\\sun}$) clusters. From this, we constrain $\\sigma_8 (\\Omega_m /0.3)^{0.37} = 0.67^{+0.04}_{-0.05}$ ($68\\%$ confidence limits), agreeing well with constraints from the 3-year WMAP CMB measurements and estimates of cluster abundances based on X-ray observations, but somewhat lower than constraints from ``cosmic shear'' weak lensing measurements in random fields.} ",
        "introduction": "\\label{sec:intro}  It has long been recognized that the abundance of massive clusters of galaxies is a sensitive probe of density fluctuations in the universe. The cluster mass function can be used to derive interesting constraints on parameters that govern structure growth, in particular the normalization of the primordial spectrum of density fluctuations and the average mass density of the universe $\\Omega_{m}$ (e.g., Evrard 1989; Frenk et al. 1990; Henry \\& Arnaud 1991; Lilje 1992; White, Efstathiou, \\& Frenk 1993; Cen \\& Ostriker 1994; Eke, Cole, \\& Frenk 1996; Henry 1997; Oukbir \\& Blanchard 1997; Bahcall \\& Fan 1998; Pen 1998; Viana \\& Liddle 1999; Henry 2000).  If the matter-energy density in the universe is dominated by a dark energy component, structure growth is hindered when this component becomes dynamically important. The resulting effect on the evolution of the cluster mass function makes it an interesting probe of dark energy (e.g., Wang \\& Steinhardt 1998; Haiman, Mohr, \\& Holder 2001; Levine, Schulz, \\& White 2002; Munshi, Porciani, \\& Wang 2004; Henry 2004; Majumdar \\& Mohr 2004). Being a sensitive probe of the growth factor, cluster abundances might be combined with probes of the expansion history such as the magnitude-redshift relation for type Ia supernovae to constrain photon-axion oscillations (e.g., Cs{\\'a}ki, Kaloper, \\& Terning 2005) or even modifications to gravity arising in brane-world cosmologies (Lue, Scoccimarro, \\& Starkman 2004). Finally, cluster abundances can provide limits on the level of non-Gaussianity in the primordial density perturbations (Scoccimarro, Sefusatti, \\& Zaldarriaga 2004; Mathis, Diego, \\& Silk 2004).  Many current studies use the X-ray temperature function (XTF) or the X-ray luminosity function (XLF) of galaxy clusters as proxies for the mass function (e.g., Pierpaoli et al.\\ 2001; Ikebe et al.\\ 2002; Reiprich \\& B{\\\"o}hringer 2002; Seljak 2002; Allen et al.\\ 2003; Pierpaoli et al.\\ 2003; Viana et al.\\ 2003; Henry 2004; Stanek et al.\\ 2006). Assuming that the X-ray emitting intra-cluster medium (ICM) is in a state of hydrostatic equilibrium, the observed X-ray temperatures $T_{\\rm X}$ and luminosities $L_{\\rm X}$ can be simply related to cluster masses. However, this assumption is not always valid, as some clusters have complex temperature structures due to ongoing merging events (e.g., Henry \\& Briel 1995; Henriksen \\& Markevitch 1996; Pratt, B{\\\"o}hringer, \\& Finoguenov 2005). A significant spatial offset between the peaks of the X-ray emission from the ICM and the corresponding concentrations of (collision-less) dark matter is observed in some of these systems (Clowe, Gonzalez, \\& Markevitch 2004; Jee et al.\\ 2005). Simulations indicate that major mergers will temporarily boost ICM temperatures and luminosities significantly above their equilibrium values and thereby significantly alter the observed XTF and XLF (e.g., Randall, Sarazin, \\& Ricker 2002). This could potentially bias constraints on cosmological parameters, if the resulting scatter around the XTF and XLF is not precisely determined for clusters in the high-mass tail of the distribution (i.e., if the box sizes of the simulations are too small to adequately sample the most massive clusters).  These problems can be remedied by obtaining independent values for the cluster masses through measurements of gravitational lensing, which do not rely on assumptions about the dynamical state of the clusters. A set of lensing-derived cluster masses can then be used to empirically calibrate the normalization and scatter in the $M-L_{\\rm X}$ and $M-T_{\\rm X}$ relations (e.g., Allen et al.\\ 2003; Smith et al.\\ 2005; Pedersen \\& Dahle 2006). Still, the complicated gas processes in merging clusters may plausibly lead to a significantly asymmetric dispersion about the mean relations which can not be adequately probed by the limited current data sets of clusters with both good X-ray and lensing data.  A different strategy is to derive the cluster mass function directly from lensing masses. The main advantage of this approach is that it should in principle avoid the problem of relating any observables of the baryonic cluster component to actual cluster mass (in practice, however, such issues are still relevant whenever the cluster sample is selected based on baryonic tracers of mass, but at a less problematic level). This approach requires lensing masses for a large sample of clusters with well-understood selection criteria. Samples from cluster surveys that efficiently pick out the extremely rare, most massive clusters in a large volume are most suitable, since the high-mass end of the mass function is most sensitive to cosmological parameters. Also, given the current level of statistical and systematic uncertainties in lensing measurements, very massive systems are needed to produce a useful lensing signal from individual clusters. However, the lensing measurements are subject to projection effects that provide an additional source of error, and need to be calibrated to reasonable precision, e.g. using N-body simulations (e.g., Clowe, de Lucia, \\& King 2004).  Here, we present the first attempt to constrain the cluster mass function by using weak gravitational lens measurements directly. Our results are only dependent on observables of the baryonic component of the clusters in the sense that the clusters were selected based on their X-ray luminosity. We describe our sample selection in \\S~\\ref{sec:sample}, and in \\S~\\ref{sec:massfun} we describe in detail how we use our mass measurements to derive the mass function. Particular emphasis is put on the treatment of statistical and systematic errors and how these affect our results.  Throughout this paper we will assume a spatially-flat cosmology with $\\Omega_m + \\Omega_{\\Lambda} = 1$. Numbers are generally given for the ``concordance model'' case ($\\Omega_m$, $\\Omega_{\\Lambda}$) = ($0.3$, $0.7$). The Hubble parameter is generally given by $H_0 = 100 h\\, {\\rm km}\\, {\\rm s}^{-1} {\\rm Mpc}^{-1}$ (however, in the discussion of the $M-L_{\\rm X}$ relation we use $H_0 = 50 h_{50}\\, {\\rm km}\\, {\\rm s}^{-1} {\\rm Mpc}^{-1}$ for easy comparison to results in the literature). ",
        "conclusions": "\\label{sec:disc}  The results presented in this paper demonstrate that weak gravitational lensing based measurements of cluster abundances provide competitive constraints in the $\\sigma_8$-$\\Omega_m$ plane. In a recent work (Pedersen \\& Dahle 2006), $\\sigma_8$ was estimated by measuring the normalization of the mass-X ray temperature relation of galaxy clusters. If the intrinsic scatter about the mass-temperature relation is low (10\\% or less), as this normalization implies a higher value of $\\sigma_8 = 0.88 \\pm 0.09$ for $\\Omega_m = 0.3$. While most theoretical studies predict a rather tight mass-temperature relation, Pedersen \\& Dahle (2006) showed that the scatter about the relation is most likely larger than 10\\% , producing a $\\sigma_8$ estimate which is systematically biased towards high values. Hence, as remarked in their discussion, the $\\sigma_8$ estimate of Pedersen \\& Dahle (2006) should most accurately be regarded as an upper limit, and the true value is likely to be lower.  Being based directly on weak gravitational lensing mass estimates, the results presented here do not depend on knowing the mass-temperature relation and its scatter, and we obtain a lower value of $\\sigma_8 = 0.67^{+0.04}_{-0.05}$ for $\\Omega_m = 0.3$. As shown in Figure~\\ref{fig:bananaSpergel} our joint constraints on $\\sigma_8$ and $\\Omega_m$ agree very well with the 3-year WMAP data (Spergel et al.\\ 2006). The 3-year WMAP data give a preference for $\\sigma_8$ and $\\Omega_m$ values lower than indicated by the first year WMAP data, combined with large scale structure data from SDSS (Tegmark et al.\\ 2004). Our result also agrees well with many estimates of cluster abundances based on X-ray data (e.g., Reiprich \\& B{\\\"o}hringer 2002; Seljak 2002; Viana, Nichol, \\& Liddle 2002; Allen et al.\\ 2003; Henry 2004), which have tended to produce $\\sigma_8$ values in the lower end of the range of published estimates. Reiprich (2006) discusses cluster abundances in light of the 3-year WMAP results, and gives best-fit results for $\\sigma_8$ and $\\Omega_m$ based on the HIFLUGCS cluster sample that agree very well with both WMAP data and the results presented here. These results imply a lower value of $\\sigma_8$ than implied by ``cosmic shear'' measurements of weak gravitational lensing in random fields (Hoekstra et al.\\ 2005; Semboloni et al.\\ 2006; see figure~\\ref{fig:bananaSpergel}).  One possible source of systematic error in the results presented in this paper is the use of X-ray luminosity selected clusters to estimate the scatter about the mass-X-ray luminosity relation. Some weak lensing studies (e.g., Erben et al.\\ 2001; Dahle et al.\\ 2003) may hint toward the existence of a class of X-ray under-luminous clusters, although detections of shear-selected clusters are significantly affected by projection effects (White, van Waerbeke, \\& Mackey 2002; Hennawi \\& Spergel 2005), and may include ``proto-clusters'' (Weinberg \\& Kamionkowski 2002) that are still undergoing collapse and have not yet become fully virialised systems. The latter type of objects may pose less of a problem for the present study, since such systems should not be included when making the fit to a theoretically predicted mass function for fully collapsed systems. A sub-population of X-ray under-luminous clusters is also found by Popesso et al.\\ (2006), based on a correlation of Abell clusters in the SDSS database and RASS sources. Since these systems show signs of being in formation, still undergoing significant mass accretion, their relation to theoretically predicted abundances of virialised clusters is somewhat unclear. Weak lensing surveys with the large sky coverage required to find statistically useful numbers of $M_{180c} \\gtrsim 10^{15} M_{\\sun}$ clusters are still some years into the future, but ongoing surveys (Miyazaki et al.\\ 2003; Hetterscheidt et al.\\ 2005; Gavazzi \\& Soucail 2006; Wittman et al.\\ 2006; Schirmer et al.\\ 2006) are beginning to detect a significant number of lower-mass clusters based on weak gravitational shear. If there is indeed a large population of massive, X-ray under-luminous galaxy clusters, our estimates of sample incompleteness as a function of mass would be too low, implying a somewhat higher true value of $\\sigma_8$.  The constraints presented here may be improved in several ways in the near future: Firstly, obtaining new wide-field imaging data for all the clusters in this sample will remove the need for extrapolation of an adopted mass model out to $r_{180c}$, removing one potential source of systematic error while significantly reducing the statistical errors on the mass measurements. Furthermore, by obtaining imaging in a larger number of pass-bands, approximate photometric redshifts may be calculated, providing sufficient 3D information to distinguish the gravitationally lensed background galaxies from cluster and foreground galaxies. Although we have made an attempt to correct for this effect by applying an average, radially dependent, contamination correction to all clusters (and estimate the extra scatter caused by cluster richness variations), the true correction factor will of course depend on cluster richness and cluster redshift, and cluster substructure will cause it not to be a simple function of cluster radius. Hence, the appropriate correction for individual clusters can differ significantly from the average correction calculated here (see e.g., Broadhurst et al.\\ 2005). Future photometric redshift estimates of background galaxies, and the uncertainties in these estimates could also be included in a optimized weighting scheme (e.g., Kaiser 2000; Benitez 2000), significantly improving the signal-to-noise ratio of the lensing mass measurements.  The sample used in this study is restricted to galaxy clusters in the Northern celestial hemisphere, and the sample size could easily be doubled by including clusters with similar X-ray selection criteria in in the Southern hemisphere, e.g., from the REFLEX sample of B{\\\"o}hringer et al.\\ (2004), thus improving the statistical errors. The sample could also be extended by lowering the X-ray luminosity cutoff $L_{\\rm min}$, sampling the mass function at lower masses. For the (e)BCS sample, such a low-luminosity extension of the cluster sample would be restricted to the near end of the redshift range probed here, since $L_{\\rm min}$ corresponds to the lower eBCS X-ray flux limit at the far end of the survey volume. However, more recent cluster surveys based on RASS data (e.g., Ebeling, Edge, \\& Henry 2001) go down to lower X-ray fluxes and would be useful for probing further down the mass function without restricting the survey volume. Finally, by combining the data presented in this paper with a weak lensing survey of a similar cluster sample at higher redshifts, interesting constraints can also be put on the dark energy equation of state parameter $w$, based on the evolution of the cluster mass function."
    },
    "0608/astro-ph0608163_arXiv.txt": {
        "abstract": "A deep 75 ks {\\it Chandra} ACIS--I data of NGC 2024 was analyzed, aiming at a search for diffuse X-ray emission in this one of the most nearby (415 pc) massive star-forming regions. After removing point sources, an extended emission was detected in the central circular region with a radius of 0.5 pc. It is spatially associated with the young massive stellar cluster. Its X-ray spectrum exhibits a very hard continuum ($kT>8$ keV) and a sign of He-like Fe K$_\\alpha$ line with the 0.5--7 keV absorption corrected luminosity of 2$\\times10^{31}$ \\ergs. Undetected faint point sources, estimated from the luminosity function of the detected sources, contribute less than 10\\% to this emission. Hence the emission is truly diffuse in nature. Because of the proximity of NGC 2024 and the long exposure, this discovery is one of the most strong supports for the existence of the diffuse X-ray emission in massive star-forming regions. ",
        "introduction": "\\label{sec:intro}  Increasing evidences of diffuse X-ray emission has been found in massive star-forming regions (MSFRs) with the {\\it Chandra X-ray Observatory}; the Rosette Nebula (at a distance of $D=1.4$ kpc; \\citealt{Townsley2003}), M17 ($D=1.6$ kpc; \\citealt{Townsley2003}), RCW 38 ($D=1.6$ kpc; \\citealt{Wolk2002}), NGC 6334 ($D=1.7$ kpc; \\citealt{Ezoe2006}), RCW 49 ($D=$1.9-7.9 kpc; \\citealt{Townsley2004}), W51A ($D=$5.5-7.5 kpc; \\citealt{Townsley2004}), NGC 3603 ($D=7$ kpc; \\citealt{Moffat2002}), the Arches Cluster ($D=8.5$ kpc; \\citealt{Yusef-Zadeh2002}), and the Quintuplet Cluster ($D=8.5$ kpc; ; \\citealt{Wang2002}). Townsley et al. (2003) explained the diffuse soft X-ray emission found in M17 in the context of strong shocks by fast stellar winds from young massive stars \\citep{Dyson1972,Castor1975,Weaver1977}. Recent results on NGC 6334 by \\cite{Ezoe2006} indicate that the spectra of the diffuse emission varies from place to place; those in tenuous molecular cloud regions are soft and thermal with temperatures of several keV, while those in dense cores exhibit harder continua with a photon index of $\\Gamma\\sim$1. They also have shown that these thermal and non-thermal spectra of the diffuse X-ray emission in MSFRs, found as a mixture in the NGC 6334 case, may be generally explained by the stellar-wind shock model. In spite of these observational progresses, even with {\\it Chandra}, there remains an uncertainty how much undetected faint point sources contribute to the emission, because these MSFRs are relatively distant ($>1.4$kpc).  In order to unambiguously examine the diffuse X-ray emission in MSFRs as possible as we can, we analyzed archival {\\it Chandra} data of NGC 2024. This region, known as the Flame nebula, is one of the nearest MSFRs ($D=415$ pc; \\citealt{Anthony1982}), located in the Orion Nebula. It is an HII region considered to be illuminated by an O8V-B2V star IRS 2b \\citep{Bik2003}. In the vicinity of IRS 2b, there are one early B star candidate associated with an ultra compact HII region G206.543-16.347 and an infrared source IRS 2 \\citep{Lenorzer2004}, and seven compact dust condensations named FIR 1-7, possibly massive protostars \\citep{Mezger1988}. In addition, $\\sim300$ low-mass ($\\simlt2$\\msolar) young stars have been found by near infrared observations \\citep{Lada1991}. The estimated age of NGC 2024 ranges from 0.3 \\citep{Meyer1996} to several Myr \\citep{Comeron1996}. A previous analysis of the same {\\it Chandra} data on point sources has been published by \\cite{Skinner2003}. In this paper, we focus on a search for the diffuse X-ray emission in this representative MSFR. ",
        "conclusions": "\\label{sec:discuss}  We have found the diffuse X-ray emission in NGC 2024. Because of the proximity of NGC 2024 and the long exposure time, this discovery itself is one of the strongest supports for the existence of the diffuse emission in MSFRs, among the previous {\\it Chandra} results. Also, this result provides us with a new observational evidence that, even in a MSFR where only late O to early B stars exist, the diffuse X-ray emission can be observed if the sensitivity is enough high. The spectral analysis suggests that the diffuse emission is dominated by the thermal emission. At the same time, since the continuum ($kT>8$ keV) of the diffuse emission is harder than typical spectra of young stars, a part of the emission may come from non-thermal origin. In NGC 6334, the soft-thermal and hard possibly non-thermal regions can be spatially distinguished \\citep{Ezoe2006}. In NGC 2024, the soft and hard regions may be co-spatial and hence be observed as a mixture of both thermal and non-thermal emission.  We then discuss whether the thermal or non-thermal interpretation is feasible in terms of energetics. First, if we assume that the whole diffuse emission is thermal, the total plasma energy $U$ will be $\\sim10^{47}\\eta^{0.5}$\\ergs where $\\eta$ is a filling factor, from the equation (3) in \\cite{Ezoe2006}. Using the X-ray luminosity $L_{\\rm X}$ and $U$, the plasma cooling time is estimated as $U/L_{\\rm X}\\sim10^8\\eta^{0.5}$ yr, which is far longer than the age of NGC 2024, from 0.3 \\citep{Meyer1996} to several Myr \\citep{Comeron1996}. Although the sound crossing time in a 10 keV plasma across the region of 0.5 pc in size is $\\sim10^3$ yr and hence short, the plasma with the estimated pressure of $\\sim5\\times10^7 \\eta^{-0.5}$ K cm$^{-3}$, calculated from the equation (5) in \\cite{Ezoe2006}, may be confined by the dense HII region known to exist around the molecular cloud dark lane (figure \\ref{fig:image} c)  \\citep{Subrahmanyan1997} and also the strong magnetic field within the molecular cloud \\citep{Crutcher1999}. Then, by dividing the total energy by the assumed age of NGC 2024 of 0.3 Myr, the average energy input is estimated as $10^{34}$ \\ergs  Second, we consider the non-thermal possibility. The flat continuum ($\\Gamma\\sim1$ or $kT>8$ keV) of the emission strongly suggests the bremsstrahlung emission by 10 keV to several MeV electrons, rather than the synchrotron or inverse Compton emission \\citep{Ezoe2006}. Since the Coulomb loss overwhelms the bremsstrahlung emission, if we assume that the diffuse emission is totally non-thermal, the necessary kinematic energy to be supplied is at least $\\sim$10$^5$ times larger than the observed X-ray luminosity, $\\sim$2$\\times$10$^{36}$ \\ergs.  One of the most plausible energy sources for the diffuse emission is the shocks generated by fast stellar winds ($\\sim2000$ km s$^{-1}$) from young massive stars. Its huge kinematic energy of the stellar winds can be easily converted via the strong shocks, among dense molecular clouds and HII regions, into the thermal (and non-thermal) energy of particles in the surrounding gases. This explanation has been proposed in M17 \\citep{Townsley2003} and NGC 6334 \\citep{Ezoe2006}. In NGC 2024, there are at least one late O to early B star IRS 2b (O8V-B2V), one early B star candidate IRS 2, and 7 possibly massive protostars. As shown in figure \\ref{fig:image}, these massive stars are spatially associated with the diffuse emission. A typical kinematic luminosity of the stellar wind is $\\sim10^{33-35}$ \\ergs per one B2V-O8V star \\citep{Howarth1989,Prinja1990,Panagia1973}. Therefore, the necessary energy supply of the thermal interpretation is explained if there is at least one massive star with a strong wind, comparable or stronger than those of a typical B0.5V star. The non-thermal interpretation is also possible if we consider all the 9 sources are massive stars and have strong winds. Wind-wind collisions may effectively increase the energy of a shock. Hence, the stellar wind scenario is possible from the viewpoint of the energetics.   YE is financially supported by the Japan Society for the Promotion of Science."
    },
    "0608/astro-ph0608177_arXiv.txt": {
        "abstract": "{We present here a study of chemical abundances in the giant elliptical galaxy M~87 using high-resolution spectra obtained with the Reflection Grating Spectrometers during two deep XMM-Newton observations. While we confirm the two-temperature structure of the inter-stellar medium (ISM) in M~87, we also show that a continuous temperature distribution describes the data equally well. The high statistics allows us, for the first time, to determine relatively accurate abundance values also for carbon and nitrogen. The comparison of the abundance ratios of C, N, O and Fe in the ISM of M~87 with those in the stellar population of our Galaxy shows that the relative contribution of core-collapse supernovae to the enrichment of the ISM in M~87 is significantly less than in the Milky Way and indicates that the enrichment of the ISM by iron through Type~Ia supernovae and by carbon and nitrogen is occurring in parallel. This suggests that the main source of carbon and nitrogen in M~87 are the low- and intermediate-mass asymptotic giant branch stars. From the oxygen to iron abundance ratio in the hot gas we estimate that the relative number of core collapse and type Ia supernovae contributing to the enrichment of the ISM in the core of M~87 is $\\sim60$\\% and $\\sim40$\\% respectively. The spatial distributions of iron and oxygen are different. While the oxygen abundance distribution is flat the iron abundance peaks in the core and has a gradient throughout the $4\\arcmin$ wide field of view of the instrument, suggesting an early enrichment by core-collapse supernovae and a continuous contribution of type Ia supernovae. ",
        "introduction": "\\label{intro}  The giant elliptical galaxy M~87 resides at the center of the nearby Virgo Cluster. Its relatively small distance of $\\sim16$~Mpc makes it a perfect target to study processes that are important in all clusters of galaxies, but to a much greater detail than possible elsewhere. Among such studies, the investigation of the structure and physics of cooling cores (previously thought to harbour cooling flows) and of the heavy element enrichment of the intra-cluster medium (ICM)  gain a central importance. This is due to the fact that M~87 is on one hand a famous cooling core cluster and on the other hand provides with an ICM temperature range of about $1-3$~keV an X-ray spectrum very rich in metal emission lines that allows us to perform detailed diagnostics of the chemical ICM abundances.  Consequently, M~87 is among the best studied objects in the sky.  Studies of M~87 at all wavelengths offer a detailed insight into the balance of the ICM cooling and the heating by the central active galactic nucleus (AGN).  The radio map of M~87 shows two lobes of synchrotron emission located about $20-30$~kpc East and South-West of the nucleus \\citep[]{owen2000}. The AGN and its inner jet probably supply the lobes with energy and relativistic electrons. Based on the previous XMM-Newton observation \\citet{belsole2001} and \\citet{matsushita2002} found that the ICM in M~87 has locally a single phase structure, except for the regions associated with radio lobes, where the X-ray emission can be modelled with two thermal components with temperatures of about 2 and 1~keV. The cooler gas fills a small volume compared to the hotter component and is confined to the radio arms, rather than being associated with the potential well of the galaxy \\citep{molendi2002}. Based on radio and X-ray observations \\citet{churazov2001} developed a model, according to which the bubbles seen in radio buoyantly rise through the hot plasma in the cluster and during their rise they uplift cooler X-ray emitting gas from the central region out to larger distances. This scenario was also studied in numerical hydrodynamical simulations by \\citet{bruggen2002} who illustrated the turbulent mixing of hot, cooler and relativistic plasma initiated by the bubble-ICM interaction. This effect probably gives rise to the observed two-phase or multi-phase temperature structure in the interaction region. Plasma cooling below a temperature of 0.8~keV is not observed in the XMM Reflection Grating Spectrometers (RGS) X-ray spectra \\citep{sakelliou2002} and neither in the XMM European Photon Imaging Camera (EPIC) spectra of M~87 \\citep{bohringer2002}.  \\citet{bohringer2001} and \\citet{finoguenov2002} showed that oxygen has a relatively flat radial abundance distribution compared to the steep gradients of silicon, sulfur, argon, calcium and iron. The strong abundance increase of iron in the core of M~87 compared to that of oxygen indicates an enhanced contribution of type Ia supernovae (SN~Ia) in the central regions of the galaxy. The flatter oxygen gradient was also confirmed by \\citet{matsushita2003}.  On the other hand, while \\citet{gastaldello2002} confirm the strong heavy-element gradient they also find a significant increase of abundance for light elements, in particular oxygen, toward the center of the galaxy which contradicts the results of \\citet{bohringer2001}, \\citet{finoguenov2002} and \\citet{matsushita2003}.  In this paper we report the results obtained by combining two deep observation with the Reflection Grating Spectrometer \\citep[RGS,][]{herder2001} aboard XMM-Newton. The strongly peaked high X-ray luminosity makes M~87 an ideal target for the RGS.  The high statistics allows us for the first time to reliably determine the abundance of carbon and nitrogen.  While it is well known that elements from oxygen up to the iron group are primarily produced in supernovae, the main sources of carbon and nitrogen are still being debated. Both elements are believed to originate from a wide range of sources including winds of short-lived massive metal rich stars, longer-lived low- and intermediate-mass stars and also an early generation of massive stars \\citep[e.g.][]{gustafsson1999,chiappini2003,meynet2002}. \\citet{shi2002} find that carbon is enriched by winds of metal-rich massive stars at the beginning of the Galactic disk evolution, while at the later stage it is produced mainly by low-mass stars. They also conclude that nitrogen is produced mainly by intermediate-mass stars. \\citet{bensby2006} find that the carbon enrichment in the Galaxy is happening on a time scale very similar to that of iron. They conclude that while at low metallicities the main carbon contributors are massive stars, carbon is later produced mainly by asymptotic giant branch (AGB) stars.  The deep observation also allows us to extract spatially resolved high-resolution spectra and to determine radial temperature profiles and abundance profiles of several elements out to $3.5\\arcmin$ from the core of M~87.  For the distance to M~87 we adopt a value of 16.1~Mpc \\citep{tonry2001} which implies a linear scale of 4.7~kpc~arcmin$^{-1}$. Unless specified otherwise, all errors are at the 68\\% confidence level for one interesting parameter ($\\Delta \\chi^{2}=1$). ",
        "conclusions": "We have analysed the temperature structure and the chemical abundances in the giant elliptical galaxy M~87 using high-resolution spectra obtained during two deep XMM-Newton observations  with the Reflection Grating Spectrometers. We found that:  \\begin{itemize} \\item While we confirm the two-temperature structure of the ISM we show that a multi-temperature {\\it{wdem}} fit also describes the data well and gives the same relative abundances as the two-temperature fit. \\item The O/Fe and O/N ratios in the ISM of M~87 are lower than in the stellar population of our Galaxy, which shows that the relative contribution of core collapse supernovae in the old stellar population of M~87 to the enrichment of the ISM was significantly less than in our Galaxy. \\item The comparison of the C/Fe and N/Fe ratios in the ISM of M~87 with those in the stellar population of our Galaxy suggests that the enrichment of the ISM by iron through Type~Ia supernovae and by carbon and nitrogen is occurring in parallel and thus the dominant source of carbon and nitrogen in M~87 are the intermediate- and low-mass AGB stars. \\item From the oxygen to iron abundance ratio we estimate that the relative number of core collapse and type Ia supernovae contributing to the enrichment of the ISM in the core of M~87 is $\\sim60$\\% and $\\sim40$\\% respectively. \\item The spatial distributions of oxygen and iron in M~87 are different. While the oxygen abundance distribution is flat the iron abundance peaks in the core and has a gradient throughout the field of view of the instrument, suggesting an early enrichment by core-collapse supernovae and a continuous contribution of SN~Ia. \\end{itemize}"
    },
    "0608/astro-ph0608494_arXiv.txt": {
        "abstract": "Weak lensing surveys have become a powerful tool for mapping mass distributions and constraining the expansion history of our Universe, but continuum surveys must average over a large number of galaxies to average down the ellipticity noise due to the unknown ellipticity and orientation of the lensed galaxies. This Letter presents a technique for measuring weak lensing with velocity maps that avoids the ellipticity noise. By studying the inherent noise characteristics, we argue this new technique could rival or exceed the sensitivity of traditional continuum observations for upcoming \\HI radio surveys. ",
        "introduction": "Weak lensing has become one of the most useful techniques for measuring foreground mass distributions, and more recently has emerged as one of the most promising ways of constraining the expansion history of the universe \\citep{Miralda-Escude1991, KaiserWL,DETF}.  Weak lensing measurements traditionally work by measuring the increase in ellipticity and tangential alignment of background galaxies \\citep{Miralda-Escude1991,BartelmannSchneiderWL}. Because the intrinsic ellipticity and alignment of each of the background galaxies is unknown, a large number of these galaxies must be averaged over. This effect is commonly called ``ellipticity noise,'' and requires weak lensing measurements to have a very high density of background galaxies.  This Letter presents a technique for measuring weak lensing with galactic velocity maps. The use of the projected velocity distribution allows weak lensing measurements based on a single galaxy, and avoids the ellipticity noise inherent in broad band lensing surveys. The same underlying effect has been previously identified by \\citet{Blain2002}. In this Letter we clarify the underlying physical process, how it can be observed, and identify the error characteristics. The attributes of the observational errors are then used to argue that this technique could be significant for upcoming \\HI observations.  Section \\ref{VMsec} introduces how lensing can be identified in velocity maps, and the inherent error characteristics. Section \\ref{Impsec} then analyzes the sensitivities, and argues that the velocity map technique may be of comparable sensitivity to broad band observations for upcoming \\HI radio surveys.  Whether broad band or velocity map techniques are more sensitive is likely to depend on the details of the observation and instrument, and the currently unknown systematics. ",
        "conclusions": "In this paper we have introduced a technique for measuring weak lensing with galaxy velocity maps. The primary advantage of the technique is that it avoids the ellipticity noise inherent in continuum weak lensing surveys. We have also argued that due to the increased utility of bright galaxies, the velocity lensing technique with planned \\HI surveys may be competitive with proposed continuum lensing surveys.  Much additional work is needed to evaluate the real-world utility of the velocity mapping technique. In the future we hope to use VLA and EVLA observations to study the systematic limits in real galaxy observations. Combined with careful studies of the \\HI luminosity function and instrumental effects of the xNTD and SKA designs, we hope determine the applicability of this technique to future weak lensing surveys."
    },
    "0608/astro-ph0608388_arXiv.txt": {
        "abstract": "Consider a perfect AO system with a very fine wavefront sampling interval and a very small actuator interval. If this AO system senses wavefront at a wavelength, $\\lambda_{WFS}$, and does science imaging at another wavelength, $\\lambda_{SCI}$, the light paths through the turbulent atmosphere at these two wavelengths are slightly different for a finite zenith distance, $z$. The error in wavefront reconstruction of the science channel associated with this non-common path effect, or so-called chromatic shear, is uncorrectable and sets an upper bound of the system performance. We evaluate the wavefront variance, $\\sigma^2(\\lambda_{WFS},\\lambda_{SCI},z)$ for a typical seeing condition at Mauna Kea and find that this effect is not negligible at a large $z$. If we require that the Strehl ratio be greater than 99 or 95\\%, $z$ must be less than about 50 or 60$^\\circ$ respectively, for the combination of visible wavefront sensing and infrared science imaging. ",
        "introduction": "It is well known that a high-performance adaptive optics system with a Strehl ratio well over 90\\% is required for the direct detection of an exoplanet from the ground especially in reflected light. Such an adaptive optics system is often called an extreme AO system (ExAO), because of its high level of sophistication in instrumentation. The ExAO has strict requirements on guide stars and wavefront sensors. Fine wavefront sampling is achieved by many pixels in the wavefront sensor, each of which corresponds to the length scale less than cm. Fine temporal sampling is also required along with the fine spatial sampling. Therefore, a large number of photons are required for a tiny area on the wavefront for a short period of time. So an AO guide star must be very bright. Another requirement on the guide star is that it must be regarded as a point source, even after the AO correction, and a presently available laser guide star, located at a finite altitude with a finite angular extent, is not fitted for this purpose, even if the problem of brightness or laser power is overcome. If the Strehl reduction due to the finite extent of the laser guide star is inevitable, there is no point of sampling the wavefront so finely, in other words, there is no use of an ExAO. For the same reason, a stringent requirement on isoplanicity will result in a small field of view. For these reasons, the major application of an ExAO requires a very bright natural guide star and the high Strehl ratio is achieved only for a small field of view. All things considered, an exoplanet search around bright nearby stars, will remain the primary scientific goal for an ExAO project.   There are many sources of errors that can cause the reduction of the Strehl ratio, e.g. see \\citep{guyon}, but most of them are in principle controllable by sophisticated instrumentation. However, there are also some errors which are in principle uncorrectable. Here we focus on one of the major uncorrectable errors: atmospheric chromatic non-common-path error, or ``chromatic shear (CS)\", associated with the different paths through the turbulent atmosphere of the light beams incident on the visible wavefront sensor and the infrared science imager, when the target star is at a finite zenith distance.   In actual astronomical observations, total observing time for an object is finite and the object never stays at the zenith. Moreover, it often happens that an object needs to be observed at a large zenith distance, due a large difference between the declination of the object and the latitude of the observatory. In order to secure a significant amount of observing time for an exoplanet search, it is inevitable to observe a target star with a finite zenith distance.  Most of the current AO systems sense wavefront at visible wavelengths and image the target at near-infrared wavelengths. This is partly due to the higher availability of a visible detector such as a CCD for a wavefront sensing camera. However, if the CS effect is a significant limiting factor for the next-generation extreme AO systems whose scientific justification is direct exoplanet detection, the choice of the combination of $\\lambda_{WFS}$ and $\\lambda_{SCI}$, may need to be reconsidered.  In this paper, we quantify the magnitude of this CS effect, in terms of wavefront variance associated with it. ",
        "conclusions": "We have obtained the wavefront variance due to atmospheric non-common path error, or chromatic shear, as a function of wavelengths for wavefront sensing and science imaging, and zenith distance, assuming an idealized adaptive optics system with a very fine actuator interval for wavefront correction. For visible wavefront sensing and infrared science imaging, the wavefront variance exceeds 0.01 at $z \\sim 50^\\circ$. For an extreme AO system currently designed, the residual error due to the finite actuator spacing, appears to dominate the wavefront error at $z < 50^\\circ$ and the non-common path effect may become dominant only at $z > 50^\\circ$."
    },
    "0608/astro-ph0608341_arXiv.txt": {
        "abstract": "The astrophysical {\\it p} process, which is responsible for the origin of the proton rich stable nuclei heavier than iron, was investigated using a full nuclear reaction network for a type II supernova explosion when the shock front passes through the O/Ne layer. Calculations were performed with a multi-layer model adopting the seed of a pre-explosion evolution of a 25 solar mass star. The reaction flux was calculated to determine the main reaction path and branching points responsible for synthesizing the proton rich nuclei. In order to investigate the impact of nuclear reaction rates on the predicted {\\it p}-process abundances, extensive simulations with different sets of collectively and individually modified neutron-, proton-, $\\alpha$-capture and photodisintegration rates have been performed. These results are not only relevant to explore the nuclear physics related uncertainties in {\\it p}-process calculations but are also important for identifying the strategy and planning of future experiments. ",
        "introduction": "The majority of the observed heavy nuclei above iron have been produced by neutron induced nucleosynthesis processes such as the slow neutron capture process ({\\it s} process) in low mass AGB stars and massive red giant stars and the rapid neutron capture process ({\\it r} process) in the supernova shock front \\citep{KBW89}.  In nature, 35 nuclei can be found on the neutron deficient side of the valley of stability ranging from $^{74}$Se to $^{196}$Hg, which are shielded against production by neutron capture processes. For the nucleosynthesis of these isotopes the {\\it p} process has been proposed as the most likely scenario. The nature of the {\\it p} process is still under debate and it is possible that it represents different independent nucleosynthesis scenarios. The contribution from charged particle induced nuclear reactions are most likely negligible because the high Coulomb barrier reduces the associated reaction rates significantly, though the possibility of some contribution to the lightest p-nuclei cannot be entirely excluded \\citep{Sch98,JoM04}. In its current interpretation the {\\it p} process is described as $\\gamma$ induced photodisintegration of stable nuclei in the shock front of type II supernovae \\citep{WoH78,RPA90} or alternatively in the deflagration flame of a type I supernova detonation \\citep{HMW91}. More recently pre-explosive sites for the $p$-process have been suggested in the O-Ne burning zone of massive stars. This feature emerges in one dimensional stellar evolution models \\citep{RauHH02} and also in two dimensional simulations of convective oxygen rich burning zones. However, in the two dimensional study the calculation of the associated nucleosynthesis pattern has been limited to a one dimensional model adaptation and does not yet include quantitative abundance predictions \\citep{BaA01}. Supercritical accretion disks associated with jets in supernovae have also been proposed as a possible scenario \\citep{FHK03}.  The bulk of {\\it p} isotopes represents a small fraction of the total abundance and is presently based on the analysis of meteorite data \\citep{AnG89}. This is illustrated in Fig. \\ref{rapp_fig1}, where the abundance distributions are compared for all those heavy isotopes, which can be entirely ascribed to a specific production process. In most cases the predicted {\\it p}-process abundances agree within a factor of three with the observed values \\citep{Lam92,Mey94,RAH95}. However, there still are significant discrepancies for the light {\\it p} nuclei $^{92,94}$Mo and $^{96,98}$Ru, which are largely overabundant compared to the model predictions, as well as for the Dy and Gd p-isotopes in mass region $A=150-160$. Such discrepancies are reviewed in more detail in \\citet{ArG03}, who give a comprehensive review of the present interpretation and questions related to {\\it p}-process nucleosynthesis. While \\citet{ArG03} discuss the global aspects of {\\it p}-process nucleosynthesis, the present work concentrates on the nuclear physics related uncertainties in {\\it p}-process model predictions and in particular on the identification of critical {\\it p}-process reaction rates, which are most important for further experimental studies.  The presently favored scenario for the {\\it p} process are type II supernova (SN) explosions \\citep{ArG03,HIS04}. The emerging shock front causes a rapid increase in temperature and density in the different layers of the pre-supernova star. As displayed in Fig. \\ref{rapp_fig2} peak temperatures between 1.7$<$T$_9<$3.3 are reached in the Ne/O layer of the presupernova envelope \\citep{RAH95}. The associated intense photon flux induces a range of photo-disintegration processes shifting the existing distribution of seed abundances to the proton rich side of the valley of stability by ($\\gamma$,n)-reactions. When this process becomes less efficient because the neutron binding energy increases with neutron deficiency, the reaction flux is maintained by the ($\\gamma$,p) and ($\\gamma$,$\\alpha$) channels. After the shock front has passed the Ne/O layers, temperature and density drop exponentially and the unstable proton rich nuclei decay back to the valley of stability. In chapter 3 the {\\it p}-process reaction flux will be discussed more quantitatively in terms of our present model predictions.  The description of the entire synthesis process for the {\\it p} nuclei requires a comprehensive reaction network involving far more than ten thousand reactions. With very few exceptions the astrophysical reaction rates have been calculated by means of the statistical Hauser-Feshbach (HF) model, which can only be applied for nuclei with high level densities. The statistical model entries depend on free parameters for the particle potentials and level densities, as well as for the $\\gamma$ widths. These parameters have to be determined, tested, and improved by laboratory measurements. The HF predictions do also depend critically on the masses of the associated nuclei, which are rather well known in the case of the {\\it p} process \\citep{AWT03}. A detailed description of HF techniques for determining nuclear reaction rates are given in the literature \\citep{RaT00,RaT01,RaT04}. The presently available HF models describe the astrophysical reaction rates typically only within a factor of two. The specific limitations with respect to {\\it p}-process analyses have been discussed by \\cite{ArG03}. The observed discrepancies between experimental data and theoretical HF predictions seem to be largest in the case of $\\alpha$ capture reactions which has been interpreted as the result of the use of insufficient $\\alpha$ potential models. We therefore will give special consideration to the impact of uncertainties in ($\\gamma,\\alpha$) reactions on $p$-process reaction flow and $p$-nuclei abundance predictions. ",
        "conclusions": "We have investigated the reaction flux patterns of the {\\it p} process within the framework of a multi-mass zone type II supernova shock front model. Particular attention was paid to the influence of {\\it p}-process reaction rates within the uncertainty limits of the theoretical Hauser-Feshbach model. The predicted {\\it p} abundances are similar to the ones observed by previous work in this field. It could be shown that the endemic underproduction of the light {\\it p} nuclei $^{92,94}$Mo and $^{96,98}$Ru are not due to the uncertainties in the thermonuclear reaction rates and must be traced back to other reasons. Possible neutron poison reactions in core helium burning and secondary neutron sources during core carbon burning could modify significantly the weak {\\it s}-process seed abundance distribution for the {\\it p}-process. This was not considered in the present paper. It should be addressed in future studies to investigate systematically the {\\it s}-process related uncertainties for {\\it p}-process simulations.  Since ($\\gamma$,n) reactions control the overall feeding of the {\\it p}-process flux, the corresponding reaction rates affect the final {\\it p}-process abundance distribution over the entire mass range. This influence is particularly visible in the range above $N=$82. On the other hand we found that the impact of ($\\gamma$,p) and ($\\gamma,\\alpha$) reactions is limited to specific mass regions. Changes in the ($\\gamma$,p) reactions translate to direct modifications of the {\\it p}-process abundances in the lower mass range with $Z\\le$50, but similar correlations were not observed in the higher mass range. In contrast, variations in the ($\\gamma,\\alpha$) reaction rates impact the abundance predictions in the higher mass range above $N=$82, while only small effects can be observed at lower masses. The overall impact of reaction rates seems not as dramatic as for other processes, the here discussed modifications of the rates translated in general to a change in the absolute abundances of a factor two to three. This will directly affect the overproduction factors displayed in Fig.~\\ref{rapp_fig4}. Improved data are necessary since a reduction of the associated uncertainty will help to identify more clearly other uncertainty factors in the present simulation. This includes the uncertainties resulting from the initial seed distribution or uncertainties yielding from insufficient model descriptions of the p-process scenario.  Past experiments have confirmed that theoretically predicted {\\it p}-process reaction rates agree within a factor of two with experimentally determined proton and neutron capture and the respective photodisintegration data. For the corresponding $\\alpha$ rates considerably larger differences of up to a factor of ten were reported. More experimental work is necessary to test the validity of the $\\alpha$ potentials used in the model calculations. It is also necessary to expand the experimental work towards {\\it p}-process reaction studies on neutron deficient nuclei to verify the reliability of the theoretical predictions for reactions involving radioactive nuclei."
    },
    "0608/astro-ph0608031_arXiv.txt": {
        "abstract": "In general relativity, the energy conditions are invoked to restrict general energy-momentum tensors $T_{\\mu\\nu}$ on physical grounds. We show that in the standard Friedmann--Lema\\^{\\i}tre--Robertson--Walker (FLRW) approach to cosmological modelling where the equation of state of the cosmological fluid is unknown, the energy conditions provide model-independent bounds on the behavior of the distance modulus of cosmic sources as a function of the redshift. We use both the  \\emph{gold} and the \\emph{legacy} samples of current type Ia supenovae to carry out a model-independent analysis of the energy conditions violation in the context of standard cosmology. ",
        "introduction": "The standard approach to cosmological modelling commences with the assumption that our $3$--dimensional space is homogeneous and isotropic at large scales. The most general spacetime metric consistent with the existence of a cosmic time $t$ and the principle of spatial homogeneity and isotropy is the FLRW metric \\begin{equation} \\label{RWmetric} ds^2 = - dt^2 + a^2 (t) \\left[\\, d \\chi^2 + S_k^2(\\chi)\\, (d\\theta^2 + \\sin^2 \\theta  d\\phi^2) \\,\\right], \\end{equation} where the function $S_k(\\chi)=(\\chi\\,$, $\\sin\\chi$, $\\sinh\\chi)$ depends on the sign of the constant spatial curvature ($k=0,1,-1$), $a(t)$ is the cosmological scale factor, and we have set the speed of light $c = 1$. A third assumption in this approach is that  the large scale structure of the universe is essentially determined by gravitational interactions, and hence can be described by a  metrical theory of gravitation such as  General Relativity (GR), which we assume in this work.  Under these very general premises, the total density $\\rho$ and the total pressure $p$ of the cosmological fluid as a function of scale factor $a$ are given by \\begin{eqnarray} \\rho & = & \\frac{3}{8\\pi G}\\left[\\,\\frac{\\dot{a}^2}{a^2} +\\frac{k}{a^2} \\,\\right]\\;, \\label{rho-eq} \\\\ p & = & - \\frac{1}{8\\pi G}\\left[\\, 2\\,\\frac{\\ddot{a}}{a} + \\frac{\\dot{a}^2}{a^2} + \\frac{k}{a^2} \\,\\right] \\;, \\label{p-eq} \\end{eqnarray} where $G$ is the Newton constant.  Note that  if now one wishes to constrain the physical properties that hold for the matter fields in the Universe it is convenient to impose the so-called \\emph{energy conditions} \\cite{Hawking-Ellis,Visser,Carroll} that limit the arbitrariness of the energy-momentum tensor $T_{\\mu\\nu}$ of these fields or, equivalently, of the physical behavior of their associated energy density $\\rho$ and pressure $p$. These conditions can be stated in a coordinate-invariant way, in terms of $T_{\\mu\\nu}$ and vector fields of fixed character (timelike, null and spacelike). The most common energy conditions are\\footnote{Although physically well-motivated, over the past years views have changed as to how fundamental some of the specific energy conditions are, and by the end  the 1970's it became clear that not all forms of matter sources obey the energy conditions. The cosmological constant $\\Lambda$ is perhaps the best known example of matter fields that violates the strong energy condition, but fulfills the weak energy condition ($\\rho \\geq 0$ and $\\rho + p \\geq 0$). See, e.g. \\cite{Visser-Barcelo,Alametal} for a short historical review.}: \\begin{itemize} \\item[(i)] The null energy condition (NEC). NEC states that $T_{\\mu\\nu}n^{\\mu}n^{\\nu}\\geq 0$ for null vectors $n^{\\mu} \\in T_{s}(M)$, where $M$ is a real 4-dimensional space-time manifold and $T_{s}(M)$ denotes the tangent space to $M$ at a point $s \\in M$.  \\item[(ii)] The weak energy condition (WEC). WEC states that $T_{\\mu\\nu}t^{\\mu}t^{\\nu}\\geq 0$ for any timelike vector $t^{\\mu} \\in T_{s}(M)$. This will also imply, by continuity, the NEC.   \\item[(iii)] The strong energy condition (SEC). SEC is the assertion that for any timelike vector $(T_{\\mu\\nu} - T/2g_{\\mu\\nu})t^{\\mu}t^{\\nu}\\geq 0$, where $T$ is the trace of $T_{\\mu\\nu}$.   \\item[(iv)] The dominant energy condition (DEC). DEC requires that $T_{\\mu\\nu}t^{\\mu}t^{\\nu}\\geq 0$ for any timelike vector $t^{\\mu} \\in T_{s}(M)$ and the additional requirement that $T_{\\mu\\nu}t^{\\mu}$ be a non-spacelike vector. By continuity this will also hold for any null vector $n^{\\mu} \\in T_{s}(M)$. \\end{itemize}    In terms of the energy-momentum tensor for a perfect fluid, $T_{\\mu\\nu} = (\\rho+p)\\,u_\\mu u_\\nu + p \\,g_{\\mu \\nu}$, the  above conditions  (see, e.g., \\cite{Hawking-Ellis,Visser,Carroll,Visser-Barcelo,Alametal,carroll1}) reduce to \\[ \\begin{array}{llll} \\mbox{\\bf WEC} \\ & \\Longrightarrow & \\rho \\geq 0 & \\ \\mbox{and} \\quad\\, \\rho + p \\geq 0 \\;,  \\\\ \\\\ \\mbox{\\bf SEC}   &\\Longrightarrow & \\rho + 3p \\geq 0 & \\ \\mbox{and} \\quad\\, \\rho + p \\geq 0 \\;, \\\\ \\\\ \\mbox{\\bf DEC}   &\\Longrightarrow & \\rho \\geq 0  & \\ \\mbox{and} \\; -\\rho \\leq p \\leq\\rho \\;, \\end{array} \\] where the NEC restriction ($\\rho+p \\geq 0$) has been incorporated by the WEC.  Thus, from  equations (\\ref{rho-eq})--(\\ref{p-eq}) one can easily rewrite the energy conditions as a set of dynamical constraints involving the scale factor $a(t)$ and its derivatives for any spatial curvature, i.e., \\begin{eqnarray} \\mbox{\\bf WEC} & \\, \\Longrightarrow & \\quad\\, - \\frac{\\ddot{a}}{a} +  \\frac{\\dot{a}^2}{a^2}  + \\frac{k}{a^2} \\geq 0 \\;, \\label{wec-eq} \\\\ \\mbox{\\bf SEC} & \\, \\Longrightarrow & \\qquad \\frac{\\ddot{a}}{a} \\leq 0 \\;, \\label{sec-eq} \\\\ \\mbox{\\bf DEC} & \\, \\Longrightarrow & \\; - 2\\left[ \\frac{\\dot{a}^2}{a^2}+\\frac{k}{a^2} \\right] \\leq \\frac{\\ddot{a}}{a}  \\leq \\frac{\\dot{a}^2}{a^2}+\\frac{k}{a^2}  \\label{dec-eq}. \\end{eqnarray} Clearly, in a expanding FLRW universe  the SEC implies that the expansion of the universe is decelerating irrespective of the sign of the spatial curvature. Furthermore, since most of the ordinary forms of matter obey the DEC they inevitably fulfill the less restrictive WEC. However, violation of the DEC does not necessarily implies violation of WEC or even of the SEC.   In this \\emph{Brief Report}, to shed some light on the energy conditions interrelations from an observational viewpoint, we use the above dynamical formulation of the energy conditions to derive model-independent bounds on the luminosity distance $d_L$ of cosmic sources in the expanding FLRW flat universe. We then concretely confront  these bounds with current SNe Ia observations, as provided by the \\emph{High-z Supernovae Team} (HzST) \\cite{Riess2004} and the \\emph{Supernova Legacy Survey} (SNLS) collaboration \\cite{Legacy2005}. ",
        "conclusions": "The energy conditions  play an important role in the understanding of several properties of our Universe, including the current and past accelerating expansion phases in the context of FLRW models. They are also necessary in the formulation and proof of various singularity theorems in classical black hole thermodynamics (e.g., the proof of the second law of black hole thermodynamics requires the null energy conditions ($\\rho + p \\geq 0$), whereas the Hawking-Penrose singularity theorems invoke the strong energy condition, $\\rho + 3p \\geq 0$ and, by continuity, $\\rho + p \\geq 0$) \\cite{Hawking-Ellis,Visser}.   In this paper, by using the fact that the classical energy conditions can be translated into differential constraints involving the scale factor $a$ and its derivatives [see Eqs.(\\ref{wec-eq}-\\ref{dec-eq})], we have derived model-independent bounds on the luminosity distance of extragalactic sources. We have confronted these energy-condition-fulfillment bounds with current SNe Ia observations from \\emph{gold} and SNLS samples and shown that all the energy conditions seem to have been violated in a recent past of the cosmic evolution. The most surprising fact is that these violations may have happened even at high-$z$ [$\\sim {\\cal{O}}(1)$], when the Universe is expected to be dominated by usual matter fields.  Finally, we emphasize that in agreement with other recent studies \\cite{visserS}, the results reported here reinforce the idea that no possible combination of \\emph{normal} matter is capable of fitting the current observational data."
    },
    "0608/astro-ph0608207_arXiv.txt": {
        "abstract": "In the present study, considering the physical conditions that are relevant in interactions between supernova remnants (SNRs) and dense molecular clouds for triggering star formation we have  built a diagram of  SNR radius versus  cloud density in which the constraints above delineate a shaded zone where star formation is allowed. We have also performed fully 3-D radiatively cooling numerical simulations of the impact between SNRs and clouds under different initial conditions in order to follow the initial steps of these interactions. We determine the conditions that may lead either to cloud collapse and star formation or to complete cloud destruction and find that the numerical results are consistent with those of the SNR-cloud density diagram. Finally, we have applied the results above to the $\\beta-$Pictoris stellar association which is composed of low mass Post-T Tauri stars with an age of 11 Myr. It has been recently suggested  that its formation could have been triggered by the shock wave produced by a SN explosion localized at a distance of about 62 pc that may have occurred either in the Lower Centaurus Crux (LCC) or in the Upper Centaurus Lupus (UCL) which are both nearby older subgroups of that association (Ortega and co-workers). Using the results of the analysis above we have shown that the suggested origin for the young association at the proposed distance is plausible only for a very restricted range of initial conditions for the parent molecular cloud, i.e., a cloud with a radius of the order of 10 pc and density of the order of 20 cm$^{-3}$ and a temperature of the order of 50$-$100 K. ",
        "introduction": "The young Beta Pictoris Association is a nearby unbound moving group formed by low-mass post-T Tauri stars. Its mean distance to the Sun is $\\sim$ 35 pc (Zuckerman et al. 2001; Torres et al. 2006). Due to its age  ($\\sim$ 12  Myr, Zuckerman et al. 2001; or 11.2 $\\pm$ 0.3 Myr, Ortega et al. 2002, 2004), this association has received a lot of attention, since it can be the host of planetary formation among its star members. Also, this  group contains the largest number of debris disks and very probably the oldest classical T Tauri type disk (Torres et al. 2006). Distinctly from the mechanism usually proposed for star formation in open clusters, Ortega et al. (2002, 2004) have suggested that a type II supernova could have triggered the formation of this loose group of stars. That could also be the case for the younger association of TW Hya 8 Myr ago (de la Reza et al. 2006).  A star formation process starts when a pressure-bounded, self-gravitating cloud or clump becomes gravitationally unstable. In the classical theory developed by Jeans, an instability occurs when the gravitational attraction overcomes the combined action of all dispersive and resistive forces. The simplest case is a system in virial equilibrium where only  the presence of the potential energy is considered. If the potential energy is greater than twice the total kinetic energy, the system collapses, while in the opposite case, it expands. This consideration can be extended by including any relevant additional physical forces acting on the gas.  It has long been known that essentially all present day star formation takes place in giant molecular clouds (GMCs, e.g. Blitz 1993, Williams, Blitz, \\& McKee 2000), so that it is vital to understand the properties, dynamical evolution and fragmentation of these clouds in order to understand star formation.  There are several possible ways of driving a cloud or clump to collapse, but at the present time is not well understood how a self-gravitating cloud becomes gravitationally unstable. GMCs are observed to contain a wealth of structures on all length scales with highly supersonic motions (Larson 1981; Blitz \\& Williams 1999; Elmegreen \\& Scalo 2004). Based on numerical simulations, a number of authors have argued that it is these supersonic motions, maintained by internal or external driving mechanisms, that induce the observed density inhomogeneities in the gas (Mac Low \\& Klessen 2004; Elmegreen \\& Scalo 2004), and that it is therefore the supersonic motions that drive star formation. Suggested candidates for an internal driving mechanism include feedback from low-mass star formation  although GMCs with and without star formation have similar kinematic properties (Williams, Blitz \\& McKee 2000). External candidates include galactic spiral shocks (Roberts 1969, Bonnell et al. 2006) and supernova and superbubble shocks (Wada \\& Norman 2001; Elemegreen \\& Scalo 2004). All of these processes seem to have sufficient energy to explain the kinematics of the ISM and can generate the observed velocity dispersion-sizescale relation (Kornreich \\& Scalo 2000). Other mechanisms, such as protostellar winds and jets, magnetorotational instabilities, expansion of H II regions and fluctuations in the UV field apparently inject energy into the ambient medium at a rate which is about an order of magnitude lower than the energy that is required to explain the random motions of the ISM at several scales.  Here we focus on one of these driving mechanisms and investigate the possibility that a SN that exploded in the recent past  produced a shock front that compressed a molecular cloud and induced the formation of the $\\beta$ Pictoris association, as suggested by Ortega et al. (2002, 2004). \\footnote{We notice that a more general study of the role that SN explosions play on the process of structure and star formation and on the generation of turbulence in the ISM of normal and starburst galaxies will be presented elsewhere (Melioli et al. 2006).}  In \\S 2, we consider the equations that are relevant for a SNR expansion and the equations that describe the interaction between an expanding SNR and a molecular cloud. In \\S 3, we obtain analytically a set of constraints from these interactions that may lead to star formation and build a diagram where these constraints approximately delineate a zone in the parameter space which is appropriate for star formation. In \\S 4, we apply these analytical results to the physical conditions that could have led to the formation of the $\\beta$ Pictoris association induced by a SNR-cloud interaction and  also describe radiatively cooling 3-D hydrodynamical simulations of the interactions between a SNR and a molecular cloud considering initial conditions which are appropriate to the $\\beta$ Pictoris environment and find that the results confirm those obtained from the analytical study. In \\S 5 we draw our conclusions. ",
        "conclusions": "We have here presented a preliminary study of the role that interactions between SNRs and dense molecular clouds play in the process of star formation.  Considering the physical conditions that are relevant  in these interactions for triggering star formation (like the choice of the cloud and SNR initial conditions, the derivation of an appropriate Jeans mass, the determination of the conditions for complete destruction of the cloud after the impact, and the determination of the extent of penetration of the shock front into the cloud) we have built in Section 3, a diagram of the SNR radius versus the cloud density for a fixed cloud radius in which the constraints above delineate a shaded zone where star formation induced by SN shock front-cloud interactions is possible.  We have also performed fully 3-D radiatively cooling numerical simulations of the impact between a SNR and a cloud for different initial conditions (in Section 4) and, although self-gravity has not been included in the present study, we have been able to track the first steps of these interactions and detect the conditions that lead  either to cloud collapse and star formation or to complete cloud destruction. We have found that the  numerical results are consistent with those established by the SNR-cloud density diagram in spite of the fact that the later have been built from approximate conditions derived from a simplified analytic theory.  Finally, we have applied the results above to the nearby young stellar association $\\beta-$ Pictoris which is composed of low mass Post-T Tauri stars with an age of 11 Myr. Ortega et al. (2004) have recently suggested that its formation could have been triggered by the shock wave produced by a SN explosion  that may have occurred either in the Lower Centaurus Crux (LCC) or in Upper Centaurus Lupus (UCL) older subgroups of the OB Cen-Sco association. Taking from their study the initial conditions that would be appropriate for both the ISM and the cloud and assuming an interacting SN shock front still in the adiabatic phase, we have found that the suggested  origin for the young association by Ortega et al. is implausible, for the proposed distance of the SNR, of  $\\sim 60$ pc, unless the parent molecular cloud had a radius of the order of 10 pc and a density of the order of 20 $cm^{-3}$, as indicated by the shaded zone of the third panel of Figure 3. A larger cloud radius ($\\sim 20$ pc) would require a much smaller cloud density (see bottom panel of Figure 3) which would not be much frequent in GMCs with a temperature greater than 50 K. These analytical results have also been confirmed by the numerical simulations (see Figures 5 to 7). The results indicate that, unless the SN had an extremely (unusual) high energy, or had exploded at a much smaller distance than the proposed one (of $\\sim 60$ pc), then the suggested scenario for triggering the formation of the $\\beta-$ Pictoris association would be possible only under the restrict conditions above. In fact, our numerical results also revealed that using a similar parent cloud density to that proposed above, but assuming an interaction with a SNR at half the distance (30 pc) and a cloud with larger radius (20 pc) the interaction may lead to the development of structuring and formation of dense cold clumps behind the shock in the cloud, which may eventually aggregate and generate dense Jeans unstable cores, as observed in most GMCs.  We should emphasize that the present study is still preliminary and the results have been mainly focused on the particular conditions of the $\\beta$ Pictoris association and its neighbourhood. Besides, it has been performed without taking into account the presence of magnetic fields. These may also play an important role on star formation as they tend to attain values as large as $\\sim  100 \\mu$G to few mG within the densest cores of the molecular clouds and can dominate over the thermal and turbulent pressures, so that  at least the densest clouds may be magnetically supported. A more general investigation of star formation processes induced by SNR shock fronts including not only radiative cooling but also magnetic fields and self-gravity is in preparation and will be presented elsewhere (Melioli et al. 2006)"
    },
    "0608/physics0608183_arXiv.txt": {
        "abstract": "\\noindent Focusing specifically on physics periodicals, I show that the journal Impact Factor is not correlated with Hirsch's $h$-index.  This implies that the Impact Factor is not a good measure of research quality or influence because the $h$-index is a reflection of peer review, and thus a strong indicator of research quality. The impact gap between multidisciplinary journals and physics-only journals is significantly reduced when $h$ is used instead of the Impact Factor.  Additionally, the impact of journals specializing in review articles is inherently deflated using $h$ because of the limited number of annual publications in such periodicals. Finally, a reordering of the top ranking journals occurs with $h$ when only the physics articles of multidisciplinary journals are considered, falling more in line with the average physicist's interpretation of a journal's prestige. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/cs0608027_arXiv.txt": {
        "abstract": "%  The myADS-arXiv service provides the scientific community with a one stop shop for staying up-to-date with a researcher's field of interest. The service provides a powerful and unique filter on the enormous amount of bibliographic information added to the ADS on a daily basis. It also provides a complete view with the most relevant papers available in the subscriber's field of interest. With this service, the subscriber will get to know the lastest developments, popular trends and the most important papers. This makes the service not only unique from a technical point of view, but also from a content point of view. On this poster we will argue why myADS-arXiv is a tailor-made, open access, virtual journal and we will illustrate its unique character.  The ADS is funded by NASA Grant NNG06GG68G. ",
        "introduction": "Searching and staying up-to-date with scholarly literature is an essential part of scientific research. With the advent of the World-Wide Web (WWW) and the evolution of electronic publishing, a powerful environment was created to open the vast universe of scientific literature on a world-wide scale. In early 1994 the WWW had become sophisticated enough to allow the search of electronic resources via \"web forms\". Over the past decade, this environment matured into an unavoidable and indispensable fact of life and tools have emerged in it that have become a crucial ingredient in scientific research. Being able to search vast amounts of data electronically obviously facilitates the review process, but applying advanced technologies such as pattern recognition to the electronic data or by allowing nested searches, one is able to produce results that are unattainable in conventional, non-electronic ways. In that sense one can argue that the available electronic search tools on the WWW even further scientific research.  Using a straight-forward search engine (Google, Yahoo, MSN, AltaVista,...) results in thousands of documents, ranked by some sophisticated algorithm. Even with the advanced versions of these tools, we still find ourselves awash in information. To search the electronic, scholarly literature, scientists need to be able to zoom in on bibliographic data using additional descriptors and search logic. What scholarly tools are available for specialists in astronomy and (astro)physics? The principal bibliographic services are the NASA Astrophysics Data System (NASA ADS), Google Scholar, INSPEC and the Astronomy and Astrophysics Abstracts. Important additional resources are the Science Citation Index (ISI Web of Science), Scopus and ZETOC. Although these tools allow researchers to zoom in on the scholarly literature, they do not offer additional tools to determine the most popular or most cited papers in a given subject. Especially for staying up-to-date, it is essential to be notified of the most popular and most cited papers. Late 2003, the ADS introduced the {\\it myADS} service \\citep{kurtz03}, a fully customizable newspaper covering (journal) research for astronomy, physics and/or the arXiv e-prints. This service will give the user an overview of the most recent papers by his/her favorite authors, and the most recent, most cited and most popular papers in a particular subject area. Additionally, the user will see an overview of citations to his/her papers. Between a fifth and a quarter of all working astronomers already subscribe to myADS \\citep{kurtz05}. {\\it myADS-arXiv} is a fully customizable, open access, virtual journal, covering the most important papers of the past week in physics and astronomy. In other words, for the specialist, the myADS-arXiv service provides a one stop shop for staying up-to-date in his/her field of interest. ",
        "conclusions": "Why is myADS-arXiv so unique and powerful? It is because of the combination of the following two factors: the search capabilities of the myADS-arXiv service guarantees the proper selection of bibliographic records for your queries, and the quality of the e-prints guarantees the relevance and importance of the bibliographic records.  The unique machinery that powers the queries consists of {\\it reference resolving} (associating reference strings with existing bibliographic records, see \\citet{accomazzi99}), {\\it bibcode matching} (setting up the e-print/paper concordance), {\\it second order operators} (operations on results lists) and {\\it also-read statistics} (\"people who read paper X, also read paper Y\", see \\citet{kurtz05}). By associating reference lists in newly added papers with already existing records, we keep citation statistics up-to-date, which allows us to generate lists of \"Most Cited Papers\". With the reads statistics, we can construct a list of \"Most Popular Papers\", using the second order operators. The bibcode matching procedure is only relevant in cases where the preprint already appeared in a journal. Titles, abstracts and author lists are indexed, so that they are available for searching.  \\begin{figure}[!ht] \\plottwo{procfig1.ps}{procfig2.ps} \\caption{Left: percentage of e-printed papers for the top 100 most cited papers for the journals indicated. Right: citations and reads per paper for ApJ papers (1997-1999)} \\end{figure}  The quality of the e-prints is reflected in the observation: the most important papers in astronomy and physics appear as e-prints on the arXiv. This fact is illustrated by figure 1 (left), which shows for a number of important astronomy and physics journals, the fraction of e-printed papers for the top 100 most cited papers, during the period of 1992 through 2005. Over 90\\% appears as e-print first (for Month. Notices RAS and Nucl. Phys. B, it is even 100\\%). Just the effect of primacy through early access is not enough to explain the fraction of e-printed papers in the top 100 most cited papers (at a given moment). According to \\citet{kurtz05b} there is an effect called \"Self-Selection Bias\" that results in a further increase of citation rates for e-printed papers. \"because papers in the arXiv are not refereed ... this suggests that authors self-censor or self-promote, or that for some reason the most citable authors are also those who first use the new publication venue\" \\citep{kurtz05b}. The papers are not cited more because they are read more, they are cited more and read more because of their quality. Figure 1 (right) illustrates the fact that e-printed papers are cited and read more than papers that did not appear as e-print.  The rule \"better searches give better results\" most definitely applies to the myADS-arXiv service. Based on the sophisticated search capabilities of the ADS, myADS-arXiv will provide the most powerful results for those who are able to characterize a field in a couple of keywords and key phrases. Additionally, the use of \"simple logic\" allows a user to really zoom in on a specific research field.  The myADS-arXiv service is unique in the world of electronic libraries and publishing. No other electronic newsletter or alerting service produces a view on the scholarly physics and astronomy literaturei as comprehensive as this service. From the Harvard-Smithsonian Center for Astrophysics press release in SpaceRef.ca (article number 16658) on the myADS-arXiv service: \"It's the best thing since two pieces of sliced bread were assembled to make a sandwich,\" said Paul Ginsparg, Professor of Physics and Information Science at Cornell University (April 17, 2005)."
    },
    "0608/astro-ph0608561_arXiv.txt": {
        "abstract": "We present the initial results from the {\\em Spitzer} Survey of the Small Magellanic Cloud (S$^3$MC), which imaged the star-forming body of the Small Magellanic Cloud (SMC) in all seven MIPS and IRAC wavebands.  We find that the \\fei/\\ftf\\ ratio (an estimate of PAH abundance) has large spatial variations and takes a wide range of values that are unrelated to metallicity but anticorrelated with 24 \\um\\ brightness and $\\ftf/\\fst$ ratio.  This suggests that photodestruction is primarily responsible for the low abundance of PAHs observed in star-forming low-metallicity galaxies. We use the S$^3$MC images to compile a photometric catalog of $\\sim400,000$ mid- and far-infrared point sources in the SMC.  The sources detected at the longest wavelengths fall into four main categories: 1) bright $5.8$ \\um\\ sources with very faint optical counterparts and very red mid-infrared colors ($[5.8]-[8.0]>1.2$), which we identify as YSOs. 2) Bright mid-infrared sources with mildly red colors ($0.16\\lesssim[5.8]-[8.0]<0.6$), identified as carbon stars. 3) Bright mid-infrared sources with neutral colors and bright optical counterparts, corresponding to oxygen-rich evolved stars. And, 4) unreddened early B stars (B3 to O9) with a large 24 \\um\\ excess. This excess is reminiscent of debris disks, and is detected in only a small fraction of these stars ($\\lesssim5\\%$). The majority of the brightest infrared point sources in the SMC fall into groups one to three. We use this photometric information to produce a catalog of 282 bright YSOs in the SMC with a very low level of contamination ($\\sim7\\%$). ",
        "introduction": "The Small Magellanic Cloud (SMC), one of the closest and most prominent neighbors of the Milky Way, is a southern hemisphere dwarf galaxy of low mass \\citep[M$_{dyn}\\sim2.4\\times10^9$ M$_\\odot$;][]{stanimirovic04} and small size (R$_*\\sim3$ kpc).  The SMC is actively interacting with its companion the Large Magellanic Cloud \\citep[LMC, studied by the {\\em Spitzer} SAGE project;][]{meixner06}. Both galaxies are being tidally disrupted by the Milky Way, and also suffer from ram-pressure interactions with the halo of hot gas surrounding the Galaxy \\citep[e.g.,][]{gardiner96}. The results of this complex dynamical interplay are the dramatic features of the Magellanic Bridge and Magellanic Stream, which dominate the high-velocity \\hi\\ emission in the Southern sky.  Befitting its nature as a gas-rich late-type dwarf galaxy, the SMC contains a number of sites of active star formation that are apparent in \\ha\\ images, although none are as prominent as the 30 Doradus starburst in the LMC.  The brightest SMC star forming complex is the giant \\hii\\ region N~66 \\citep[also known as DEM~103, or NGC~346;][]{henize56,dem76}, which contains approximately $60$ O stars \\citep{mpg89}. The star formation activity of the SMC, as well as most of its gas, are concentrated in the feature known as the ``Bar'' (probably the body of the SMC proper). A second feature, the ``Wing'', is an extension roughly perpendicular to the Bar on its eastern side that continues into the \\hi\\ bridge that joins both Clouds (Figure \\ref{coverage}).  As one may expect in an interacting system, the three-dimensional structure of the galaxy remains quite unclear. Recent analyses show that the prominent irregular features, which are apparent in the light of young stars, disappear when looking at the old stellar populations \\citep{cioni00,zaritsky00,maragoudaki01}. These are regularly distributed, and the shape of the SMC appears spheroidal. Observations indicate that the distances to the Bar and the Wing are measurably different, and that the line-of-sight depth of the SMC is not negligible, amounting to 6 to 12 kpc \\citep[][and references therein]{crowl01}. The distance to the center of the SMC is approximately 61.1 kpc, corresponding to a distance modulus of $m-M=18.93$~mag and a spatial scale of 0.3 pc/\\arcsec \\citep{hilditch05,keller&wood06}. To put this in context, the spatial resolution of the $3.6-8.0$ \\um\\ data presented here is similar to the 60 \\um\\ resolution of the {\\em Infrared Astronomical Satellite} (IRAS) at 2 kpc.  The thickness of the SMC introduces a dispersion of $\\sim0.2$~mag in the distance modulus for objects located near either edge. The line-of-sight to the SMC is very clear, with a foreground extinction $A_V\\sim0.12$, corresponding to a reddening of $E(B-V)\\sim0.04$ \\citep{schlegel98}.  The proximity of the SMC makes it the lowest metallicity gas-rich galaxy that can be studied in great detail. For observational reasons metallicity, $Z$, is commonly quantified using the oxygen abundance relative to hydrogen. The gas-phase oxygen abundances measured in SMC \\hii\\ regions are $12+\\log{\\rm(O/H)}\\approx8.0$ \\citep{dufour75,dufour77,dufour84}, only a factor of 4 to 5 higher than those observed in the lowest metallicity galaxies known (I~Zw~18 and SBS~0335-052), and a factor of 5 lower than solar metallicity. The same metallicity with very little dispersion is observed in several stellar types in the SMC, including K, F, and A supergiants and Cepheid variables \\citep{russell89,hill97,luck98,venn99}. Its low heavy-element abundance suggests that the SMC may be one of the best local templates for studying the primitive interstellar medium and providing insight into the processes at work in primordial galaxies at high redshift. In particular, the interaction between massive stars, gas, and dust is unusually accessible in this nearby galaxy.  Compared to the Milky Way, the dust properties of the SMC are peculiar. For quite a while it has been recognized that the UV extinction curve in the SMC lacks the 2175 \\AA ``hump'' \\citep[e.g.,][]{gordon03} which is interpreted as signaling the lack of a population of small carbonaceous dust grains \\citep{li02}. Indeed, most current models of dust composition in the SMC rely heavily on silicates to reproduce its extinction curve, with typical silicate to carbon ratios by mass in the range $2-12$ \\citep[e.g.,][]{weingartner01}.  Whether this dearth of carbonaceous grains is maintained for the large grains remains unclear, as there is a well known anticorrelation between the specific frequency of carbon stars (the main producers of carbonaceous grains) and galactic metallicity \\citep[e.g.,][]{blanco80}, that results in a large abundance of these stars in the SMC.  Other lines of research also suggest a dearth of the smallest dust grains in the low-metallicity interstellar medium of the SMC. The first wide-field infrared maps of the SMC were acquired by IRAS and presented by \\citet{si89}.  \\citet{sauvage90} studied the IRAS colors of the SMC, comparing them to the emission from larger, higher metallicity galaxies, and concluded that there was evidence for a paucity of the small grains that are responsible for the bulk of the extended 12 $\\mu$m emission. More recently, \\cite{snez00} reprocessed the IRAS data to produce maps with improved angular resolution of 1\\arcmin\\ at 12~\\um, 25~\\um, and 60~\\um, and 1\\farcm7 at 100~\\um.  By comparing these data with \\hi\\ observations of the SMC at similar resolution, they were able to study the large-scale properties of the ISM.  \\citeauthor{snez00} used these far-infrared (FIR) observations to measure the dust mass, the dust temperature and its spatial variations, and the dust-to-gas ratio, DGR, across the SMC.  This study found a mean dust-to-gas ratio in the SMC that is a factor of $\\sim100$ times lower than in the Milky Way, providing a different environment for star formation to take place.  It also concluded that the dust in the SMC is dominated by large grains, with weaker contributions from polycyclic aromatic hydrocarbons \\citep[PAHs, typical sizes 4\\AA$<a<$12\\AA;][]{desert90} and very small grains (VSGs, 12\\AA$<a<$150\\AA) than are found in the Galaxy. One of the limitations of this study is the use of the 60 \\um\\ IRAS waveband to determine dust temperature. This waveband is expected to have a substantial contribution from stochastically heated small dust grains, out of equilibrium with the radiation field.  \\citet{bot04} expanded this work by adding {\\emph{Infrared Space Observatory}} (ISO) 170~\\um\\ imaging of the SMC to the IRAS data. Using a combination of 100 and 170 \\um\\ data they measured an appreciably higher dust-to-gas ratio, $\\sim30$ times lower than Galactic. They also modeled the spectral energy distribution to decompose the emission into the contributions by PAHs, VSGs, and large grains, finding a substantial 60 \\um\\ excess. Such excess could be caused by an enhanced interstellar radiation field in the SMC, or perhaps by a change in the grain size distribution with respect to the Galaxy, with more VSGs relative to large grains. \\citeauthor{bot04} suggest that grain destruction, perhaps shattering by supernovae shocks, is the driver behind the different properties of the dust in the SMC and the Milky Way.  These studies point to a dust-to-gas ratio in the SMC much lower than Galactic, scaling with its metallicity as DGR~$\\propto Z^{2}$, a result similar to that found for a collection of galaxies by \\citet{lisenfeld98}. Note that measurements of the mass-loss rate of the most evolved asymptotic giant branch (AGB) stars in the SMC, however, find a much larger dust-to-gas ratio, scaling as DGR~$\\propto Z^{1}$ \\citep{vanloon06}. Because these stars are believed to dominate the ISM dust production, this observation suggests that efficient destruction of dust takes place in the metal-poor star-forming ISM, along the lines of the evidence mentioned in the previous paragraphs.  Detailed modeling of SMC UV absorption and FIR emission data by \\citet{weingartner01}, \\citet{li02}, and \\citet{clayton03} indicate that a deficit of PAHs, as well as the lower abundance of VSGs suggested by IRAS, and perhaps a change in composition with more silicate and fewer graphite grains are required to reproduce the observations. \\citet{contursi00} studied N~66 using the cameras onboard ISO, and found that the emission from the Aromatic Infrared Band (AIB) in the 6.75 $\\mu$m ISO waveband, commonly attributed to PAHs, was faint compared to the 15 $\\mu$m dust continuum. This was confirmed spectroscopically using ISO's continuously variable filter (CVF) to obtain low-resolution spectra toward emission peaks, all of which show very weak or nonexistent AIB emission. As a counterpoint, \\citet{reach00} used the CVF to obtain spectroscopy toward a quiescent molecular cloud in the SMC. They measured AIB-to-continuum ratios similar to those observed in the Milky Way, pointing to similar relative abundances of PAHs.  These data suggest that PAHs do form and survive in the SMC and other star-forming low-metallicity galaxies in quantities very similar to Milky Way PAHs, where the environment allows it. Indeed, SMC observations show us that dust properties are environment dependent, and one of the main motivations for a new infrared survey of the SMC is to study and characterize the influence of the local environment on the properties and composition of interstellar dust.  The \\emph{Spitzer Space Telescope} allows dramatic improvements in our ability to study the ISM of the SMC, with previously unattainable images at 3.6 to 8.0 \\um\\ and a factor of $4-10$ higher angular resolution and typically an order of magnitude better sensitivity at FIR wavelengths than previous missions.  The {\\em Spitzer} data enable us to take a giant step forward in the study of the detailed, distribution of all types of dust grains, including PAHs which primarily emit at wavelengths shorter than 12~\\um.  The high angular resolution of \\emph{Spitzer} means that point sources can be separated from diffuse emission out to 70~\\um, allowing investigations of large samples of IR point sources (such as carbon stars and young stellar objects) and simultaneously providing a cleaner picture of the ISM. The unprecedented sensitivity and resolution of the {\\em Spitzer} data makes it possible to determine the abundance and distribution of dust grains and PAHs in relationship to their environment, to test and constrain dust models, and to obtain a complete census of star formation in the SMC, the nearest ``primitive'' galaxy.  In this paper we present an overview of the \\emph{Spitzer} Survey of the Small Magellanic Cloud (S$^3$MC), discussing the observations, methodology, and some of the initial results on the extended emission and the infrared point sources in the SMC. This is the second in a series of papers presenting the results of the survey \\citep[see also][]{snez05}. In \\S\\ref{obs} we describe the observations, in \\S\\ref{datared} we describe in detail the data reduction process employed to obtain the final S$^3$MC images, in \\S\\ref{photometry} we discuss the point source photometry for the global survey and the comparison with existing datasets, in \\S\\ref{discussion} we analyze the extended dust emission and discuss the different stellar populations detected by {\\em Spitzer}, and in \\S\\ref{conclusion} we summarize our findings. ",
        "conclusions": "\\label{conclusion}  We have used the {\\em Spitzer Space Telescope} to image the star-forming body of the SMC in the mid and far IR with unprecedented sensitivity and resolution.  IRAC images of the SMC show a wealth of diffuse emission across the SMC seen for the first time. Particularly striking is the diffuse emission at 8 $\\mu$m, often associated with star-forming regions.  We have used these images to take a first look at the \\fei/\\ftf\\ ratio (an estimator of PAH abundance) across the SW Bar region of the SMC. We find that this ratio has large spatial variations, and takes a wide range of values even in a small region.  Extended emission in the southernmost part of the SW Bar has $\\fei/\\ftf>0.7$, typical of solar-metallicity galaxies, while near the sites of massive star formation and bright 24 \\um\\ emission (suggesting large radiation fields) the ratio drops to the $\\fei/\\ftf<0.1$ observed in metal-poor galaxies \\citep{engelbracht05}. Because existing surveys of abundances in \\hii\\ regions find no evidence for a gradient (or even a large scatter) for metallicity in the SMC, we conclude that the wide range of values taken by the \\fei/\\ftf\\ ratio in this galaxy, and the wide range of PAH abundance suggested by it, is primarily driven by the local ISRF and not metallicity. We find that the relationship between the \\fei/\\ftf\\ ratio and the FIR dust color represented by the \\ftf/\\fst\\ ratio can be described by $\\log[\\fei/\\ftf]\\approx -0.7\\log[\\ftf/\\fst]-1.4$, showing that regions of blue FIR color are deficient in 8 \\um\\ PAH emission relative to their 24 \\um\\ intensity.  We combine these observations with optical and near-IR photometry in the literature to produce a catalog of point sources that spans from optical wavelengths to 70 $\\mu$m. We use this catalog to investigate the nature of the sources detected in the {\\em Spitzer} wavebands.  We find that the most prominent sources detected at the long wavelengths fall into four categories: 1) sources with very faint optical counterparts and very red ($[5.8]-[8.0]>1.2$) colors.  This population corresponds to YSOs, with some contamination from background galaxies and foreground dwarfs that we estimate using the SWIRE ELAIS-N1 catalog.  2) Bright MIR sources ($-12<M_{5.8}<-7.5$) with midly red colors ($0.16\\lesssim[5.8]-[8.0]<0.6$), many of which have bright optical counterparts. These are identified as carbon stars.  3) Bright mid infrared sources with neutral colors and bright optical counterparts, corresponding to oxygen-rich AGB and red supergiant stars. And 4) a class of dusty early B stars (B3 to O9, based on the optical photometry) that are not appreciably reddened optically, but that have large excesses at 24 $\\mu$m most of which cannot be explained by free-free emission. We use our multiwavelength data to compile a catalog of 282 bright YSOs in the SMC, where we choose the color-magnitude cuts to allow for only 7\\% contamination by background or foreground sources (Table 4).  Finally, we discuss the population of early B stars detected at 24 \\um, which constitutes a few percent ($\\sim3-5\\%$) of the B stars in a narrow region of the optical color-magnitude space, and shares many of the observational signatures of debris disks or cirrus hotspots in the Milky Way. Based on their small near IR excess and the fact that most of these objects are not found in existing objective prism catalogs, we conclude that the bulk of these systems are not emission line stars. They may constitute a population of young massive stellar systems caught in the act of photoevaporating their accretion disks, but the timescale derived by models for such process suggests that this is unlikely. Dust associated with debris disks or cirrus hotspots (i.e., the Pleiades effect) appear to be the most likely explanation for the origin of their 24 \\um\\ emission.  The S$^3$MC point source catalog constitutes a deep and uniform photometric database for a large sample of stars located at the same distance, with similar metallicities, and low levels of foreground and background confusion. We expect that one of the contributions of this work will be to supply the basis on which to test different schemes for identifying mid- and far-infrared sources. The S$^3$MC observations provide a unique resource for the study of stellar evolution, dust production, dust properties, and the interaction between the ISM and star formation in an environment that, because of its proximity and low metallicity, has closer resemblance to the conditions prevalent in primitive galaxies than any other galaxy that can be studied in similar detail."
    },
    "0608/astro-ph0608427_arXiv.txt": {
        "abstract": "The five enigmatic ``Cocoon stars\" after which the Quintuplet cluster was christened have puzzled astronomers since their discovery. Their extraordinary cool, featureless thermal spectra have been attributed to various stellar types from young to highly evolved, while their absolute luminosities places them among the supergiants. We present diffraction-limited images from the Keck 1 telescope which resolves this debate with the discovery of rotating spiral plumes characteristic of colliding-wind binary ``pinwheel\" nebulae. Such elegant spiral structures, found around high-luminosity Wolf-Rayet stars, have recently been implicated in the behavior of supernovae lightcurves in the radio and optical. \\clearpage ",
        "introduction": " ",
        "conclusions": "Observations of the Quintuplet cluster are listed in Table~S1 and were made with the near-infrared NIRC camera on the Keck~1 telescope. The nomenclature of Moneti et al.\\citep{moneti94} is adopted here, identifying the dusty red Quintuplet proper members, or cocoon stars, as Q\\,1, Q\\,2, Q\\,3, Q\\,4 and Q\\,9. Locations of each of these stars within the cluster as imaged by the Hubble Space Telescope is given in Fig.~S1, which also depicts our diffraction-limited Keck images of Q\\,2, Q\\,3 overlayed with graphical indication of the relative spatial scale.  The high angular resolution imaging utilized rapid-exposure speckle interferometry techniques from a number of long-standing experiments at the Keck~1 \\citep{keckmask}. Typical datasets entailed from one to a few hundred short-exposures ($\\sim$0.14\\,sec) with the NIRC camera operating in a high-magnification (0.0206 arcsec/pixel) mode. Interleaved observations of galactic-center object IRS~7 were used to calibrate the telescope-atmosphere point-spread function (the unresolved nature of this object was itself checked against stars HD~159255 and HD~163042).  All five Quintuplet cocoon stars were found to be spatially resolved to some degree. The full-width at half-maximum (FWHM) of best-fitting circular Gaussian functions, used here to give an estimate of apparent size, are given in Table~S1. Images of stars Q\\,2 and Q\\,3, recovered with Fourier techniques\\citep{keckmask} and presented earlier in Fig.~1, show that complex and asymmetric structures exist within the dust shells surrounding these stars. The complexity of these visibility functions resulted in poor fits for the Gaussian model, and we therefore limited the fit range to low spatial frequencies/short baselines where visibility excursions are not yet strongly manifest ($<$2$\\times$10$^6$ and 1$\\times$10$^6$\\,rad$^{-1}$ for 2.21 and 3.08\\,$\\mu$m respectively: equivalent to $\\simle$4.5 and 3.1\\,m).  By implication, if similar physics is assumed to pertain to the remaining three of the five Quintuplet cocoon stars, then the utility of fitting Gaussian shells may appear limited. Despite this, for objects which are only partially resolved, the fitting of such simple functional forms to interferometer data gives a good estimate of the overall size, and permits quantitative comparison with models. Furthermore, this allows sizes for morphologically complex objects Q\\,2 and Q\\,3 to be compared against measurements of the partially-resolved Q\\,1, Q\\,4 and Q\\,9, and against other dusty Wolf-Rayets observed with high resolution techniques. Note also that FWHM in the range 10--15\\,mas are only marginally resolved with this experiment; relative errors in these cases are correspondingly high.  Apparent sizes have been combined with flux measurements to yield estimates of surface brightness in our two filter bandpasses for the target stars. Correction for the $A_v$=29$\\pm$5 visible extinction was made using the optical dust constants of Mathis\\citep{mathis90}. Surface brightnesses were then derived following Monnier et al.\\citep{wr_size}, which also gives a comparison population of Galactic dusty Wolf-Rayets with similarly measured surface brightnesses. For both filter bandpasses, the Quintuplet cocoon stars showed surface brightnesses within the range spanned by the WR population studied in Monnier et al.\\citep{wr_size}. In particular, for Q4 and Q9 the color temperature between 2.21 and 3.08\\,$\\mu$m appeared to be in reasonable accord with similar measurements from the Galactic population. This finding was found to be generally robust against variations in the extinction correction over the expected range.  However, in some regards the Quintuplet WRs, and in particular Q\\,2 and Q\\,3, did appear to be distinct from the Galactic population. The observed increase in size between 2.21 and 3.08\\,$\\mu$m which approaches a factor of $\\sim$2 in the Quintuplet, is found to be a more modest $\\sim$1.4 elsewhere\\citep{wr104-idlm,wr_size}. Such changes of size (and surface brightness) with wavelength reflect the fact that dust with a range of temperatures contributes to the near-IR emission. The dramatic enlargement between 2.21 and 3.08\\,$\\mu$m argues for a flatter thermal profile in the Quintuplet dust shells: plausibly due to external heating from stars in the dense central region of the cluster. The comparison population of Galactic WRs were in far less crowded regions where the outer dust shell likely receives little or no energy except that originating with the central WR star.  Perhaps the simplest way to test the hypothesis that the remaining 3 cocoon stars are also pinwheels is by photometric monitoring. Variability of these sources has already been noted\\citep{glass02}, although not with sufficient coverage to reveal a cyclic change consummate with any rotational period\\citep{kato02}. The $\\sim$2.5\\,yr periods inferred from the imaging extend the confirmed operation of continuous dust formation by the pinwheel mechanism to significantly longer periods/larger binary separations than previously known, with implications for models of these processes.  This finding of late-type WC binaries in the Quintuplet means that WC (Carbon rich) outnumber WN (Nitrogen rich) Wolf-Rayets by 11:6, and furthermore all WC stars are dusty. This makes an interesting contrast to the massive young WR-rich cluster Westerlund~1 where the WC:WN ratio is reversed to 7:12, with none of the WC's exhibiting dust. Clearly, the close binaries at the heart of the Pinwheel systems can be responsible for significant modification of the stellar evolutionary path, in which mass-transfer or envelope-stripping events might precipitate the WC phase. This entanglement of binarity, mass-loss history, and evolutionary path has the potential to skew population distributions, although it is unclear exactly what conditions resulted in the abundance of Pinwheels in this cluster."
    },
    "0608/astro-ph0608537_arXiv.txt": {
        "abstract": "We use cosmological hydrodynamic simulations with enriched galactic outflows to compare predictions for the galaxy mass-metallicity ($M_*-Z$) with observations at $z\\approx 2$ from \\cite{erb06}.  With no outflows included galaxies are over-enriched, indicating that outflows are required not only to suppress star formation and enrich the IGM but also to lower galaxy metal content.  The observed $M_*-Z$ slope is matched both in our model without winds as well as our favored outflow model where the outflow velocity scales as the escape velocity, but is too steep in a model with constant outflow speeds.  If outflows are too widespread at early times, the IGM out of which smaller galaxies form can become pre-polluted, resulting in a low-mass flattening of the $M_*-Z$ relation that is inconsistent with data.  Remarkably, the same momentum-driven wind model that provides the best agreement with IGM enrichment data also yields the best agreement with the $z\\approx 2$ $M_*-Z$ relation, showing the proper outflow scaling and strength to match the observed slope and amplitude.  In this model, the $M_*-Z$ relation evolves slowly from $z=6\\rightarrow 2$; an (admittedly uncertain) extrapolation to $z=0$ broadly matches local $M_*-Z$ observations.  Overall, the $M_*-Z$ relation provides critical constraints on galactic outflow processes during the heydey of star formation in the Universe. ",
        "introduction": "Galactic outflows are observed to carry mass and metals into the intergalactic medium~(e.g. \\cite{vei05}).  At high redshift, outflows are ubiquitous from star-forming galaxies~(\\cite{pet01}), and the metals that they carry are seen in a wide range of environments (\\cite{pet06}). \\cite{pet99} originally noted the missing metals problem, namely that the amount of metals present in Lyman break galaxies at $z\\sim 2.5$ falls far short of the amount expected to be produced by their stars (see \\cite{dav06b} for an updated discussion).  Hence outflows appear to have a large impact on the evolution of galaxies during the heydey of star formation in the Universe, and metals provide a convenient avenue with which to trace and quantify ejected gas.  In \\cite{opp06} (hereafter OD06) we implemented a number of parameterized models for galactic outflows into \\gad\\ cosmological hydrodynamic simulations.  We tested these models against observations of \\ion{C}{iv} quasar absorption line systems at $z\\sim 2-6$, and found that only a relatively narrow range of outflow parameters were capable of enriching the diffuse intergalactic medium (IGM) early enough while not overheating it.  Intriguingly, the most successful models were ones that employed a momentum-driven wind scenario~(\\cite{mur05}) that is also favored by local observations of starburst outflows~(\\cite{mar05}).  Hence it is plausible that distant outflow-driving galaxies obey scalings similar to local ones.  Another critical test of such outflow models is measuring the metals that get left behind in galaxies rather than driven into the IGM. Observationally, this can be traced using the galaxy mass-metallicity relation.  \\cite{tre04}, using the Sloan Digital Sky Survey, found that low-redshift emission-line galaxies had a gas-phase metallicity that increased with stellar mass up to $\\sim 10^{10.5}M_\\odot$, and then flattened to higher masses.  At $z\\sim 1$, smaller survey areas make it difficult to track the high-mass flattening, but indications are that galaxies at a given mass are only mildly underenriched compared to present-day ones~(\\cite{sav05}).  This evolution was extended out to $z\\sim 2$ by \\cite{erb06}, who found that galaxies back then are roughly one-half as metal-rich as galaxies today at a given stellar mass. This is mildly surprising as only about one-quarter of all stars (and therefore, presumably, metals) have formed by then~(\\cite{rud03}). However, most models of cosmic chemical evolution predict earlier enrichment in the highly biased environments that form early stars (e.g. \\cite{cen99}, \\cite{dav06}).  In any case, if the solution to the missing metals problem is that the majority of metals have been ejected from galaxies~(e.g. \\cite{dav06b}), then the amount and distribution of metals left behind in galaxies likely offers key insights into the nature of outflows.  In these proceedings we test our outflow models in OD06 against the $z\\approx 2$ mass-metallicity data of \\cite{erb06}, and study the evolution of the mass-metallicity relation.  We show that the same outflow model that best matches the IGM data also best matches the mass-metallicity relation, and furthermore the models that do not match the IGM data provide a poorer fit.  This lends further support to the idea that momentum-driven galactic winds in a hierarchical galaxy formation setting can explain a wide range of observables at $z\\geq2$. ",
        "conclusions": "Using our cosmological hydrodynamic simulations incorporating galactic outflows, we investigate the nature of the mass-metallicity relation in galaxies at high redshift, and compare it to $z\\approx 2$ observations by \\cite{erb06}.  We find that: \\begin{itemize} \\item Outflows are required in order to reduce the metallicity of galaxies at $z=2$.  With no outflows, galaxies at a given stellar mass are twice as enriched as they should be by that epoch, comparable to galaxies today. \\item The slope and amplitude of the $M_*-Z$ relation is naturally reproduced in our momentum-driven wind model that is also favored from comparisons to IGM enrichment observations from $z\\approx 5\\rightarrow 2$. \\item The constant wind-speed outflow model produces an $M_*-Z$ relation that is flat at low masses, possibly owing to pre-pollution from early widespread winds, and rises too steeply at high masses, both of which are in disagreement with observations. Hence it does not appear that smaller galaxies preferentially lose more material to outflows at these epochs. \\item The mass-metallicity slope is established at early times ($z>6$) and remains constant at $Z\\propto M_*^{0.3}$, while the amplitude evolves slowly with redshift ($\\approx 0.06$~dex per Gyr).  Galaxies tend to evolve mostly along the $M_*-Z$ relation towards higher stellar masses. \\item A simple, albeit questionable, extrapolation to $z=0$ yields a relation that is in good agreement with data from~\\cite{tre04}. The unbroken slope to the lowest observed masses~(\\cite{lee06}) provides another non-trivial constraint that is met in our momentum-driven wind scenario. \\end{itemize}  In summary, the mass-metallicity relation provides another critical constraint on galactic outflows across cosmic time.  Understanding the complex processes that establish the shape and evolution of the $M_*-Z$ relation will likely require moving beyond simple closed box model variants, and may benefit from insights gained using numerical simulations that include metal production and distribution mechanisms.  The comparisons here strengthen our claim that momentum-driven wind scalings (whether or not they actually arise from momentum-driven winds) are able to properly distribute metals throughout the Universe at high redshfits."
    },
    "0608/astro-ph0608701_arXiv.txt": {
        "abstract": "We reinvestigate the generation and accumulation of magnetic flux in optically thin accretion flows around active gravitating objects. The source of the magnetic field is the azimuthal electric current associated with the Poynting-Robertson drag on the electrons of the accreting plasma. This current generates magnetic field loops which open up because of the differential rotation of the flow. We show through simple numerical simulations that what regulates the generation and accumulation of magnetic flux near the center is the value of the plasma conductivity. Although the conductivity is usually considered to be effectively infinite for the fully ionized plasmas expected near the inner edge of accretion disks, the turbulence of those plasmas may actually render them much less conducting due to the presence of anomalous resistivity. We have discovered that if the resistivity is sufficiently high throughout the turbulent disk while it is suppressed interior to its inner edge, an interesting steady-state process is established: accretion carries and accumulates magnetic flux {\\em of one polarity} inside the inner edge of the disk, whereas magnetic diffusion  releases magnetic flux {\\em of the opposite polarity} to large distances. In this scenario, magnetic flux of one polarity grows and accumulates at a steady rate in the region inside the inner edge and up to the point of equipartition when it becomes dynamically important. We argue that this inward growth and outward expulsion of oppositely-directed magnetic fields that we propose may account for the $\\sim 30$~min cyclic variability observed in the galactic microquasar GRS1915+105. ",
        "introduction": "The issue of the orgin of cosmic magnetic fields, despite much progress in the field, remains an open issue in  astrophysics. The problem is basically of a topological nature, because the magnetic field, being a solenoidal vector ${\\bf B}=\\nabla\\times {\\bf A}$ (${\\bf A}$ is the magnetic vector potential), is necessarily  absent in a homogeneous and isotropic universe. For the same reason, the magnetic field cannot be generated by the potential fluid motions that result from the growth of scalar inhomogeneities in such a universe, as current models contend. As such, the evolution of the magnetic field is akin to that of vorticity ${\\bf \\omega = \\nabla \\times {\\bf v}}$ (${\\bf v}$ is the flow velocity), a fact that has been noticed by at least one (to our knowledge) paper in the literature (see Eqs.~[4] and [7] in Kulsrud et al. [1997]). It is also of interest to note that the source of both ${\\bf\\omega}$ and ${\\bf B}$ is proportional to  $\\nabla p \\times  \\nabla \\rho/\\rho^2$ ($p$ is the pressure and $\\rho$ is the density), a quantity which is nonzero only in the presence of nonbarotropic fluids (Kulsrud et al. [1997] consider as such curved shocks and photoheating). This source term, while generally small, is nonetheless of great importance  because in its absence the initial values of  {\\bf $ \\omega$}, {\\bf B} = 0 are preserved by the evolution equations. It is generally referred to as the {\\em Biermann battery} and leads to magnetic fields that are usually quite weak initially ($\\sim 10^{-20}$ G) but thought to be subsequently amplified to the observed mean galactic values ($\\sim 10^{-6}$ G) by dynamo processes.   Some years ago, we proposed a source-term alternative to the term of the Biermann battery (Contopoulos \\& Kazanas [1998], hereafter CK). More specifically, we posited that the Poynting-Robertson radiation force acting predominantly on the electrons of an accretion flow around an active astrophysical source may generate toroidal electric currents sufficiently large to support poloidal magnetic fields that in certain situations could approach equipartition values. The Poynting-Robertson source term is proportional to $\\nabla\\times(\\ell \\, {\\bf\\omega})$, where $\\ell= L \\, \\sigma_{\\rm T}/4 \\pi r m_ec^3$ is the usual accreting source compactness ($L$ is the source luminosity, $\\sigma_T$ is the Thompson cross-section, $m_e$ is the electron mass, $c$ is the speed of light, and $r$ is the distance from the center); as such, besides its much larger magnitude compared to the term of the Biermann battery, this term is additionally important because it provides a hitherto unexplored, direct coupling  between the magnetic field and the vorticity equations of motion.   In our earlier work (CK) we concentrated our attention on the effects of the Poynting--Robertson source term on the evolution of the magnetic field within the volume of the disk alone, by solving the corresponding (induction) equation in that region.  Our study showed that, in a highly conducting plasma, currents and fields grow only for a restricted time of the order of the accretion time, and thus they quickly saturate to values well below equipartition (but still much higher than those induced by the Biermann battery). However, for plasma conductivity below a certain critical value (and a free inner boundary condition that allows the accumulation of magnetic flux interior to that point), we found a different regime in which the magnetic field accumulated around the gravitating object grows at a constant rate; we then showed that in several astrophysically interesting systems, magnetic fields may reach even equipartition values within astrophysically relevant time scales.   This `unrestricted' field growth has its origin on the fact that the magnetic field loops, generated by the toroidal Poynting-Robertson current inside the disk, open up to infinity once they reach the surface of the disk because of its differential rotation (and the ideal MHD conditions outside the disk). This opening up, coupled with sufficiently low conductivity, allows the separation of {\\em inward advected} field of one polarity from the {\\em outward diffusing} field of the opposite polarity. This effect of magnetic loop opening above and below the disk is central in our mechanism and it has been clearly described in our earlier paper. Our work was criticized by Bisnovatyi-Kogan, Lovelace \\& Belinski (2002), who, arguing along the lines of earlier work by Bisnovatyi-Kogan \\& Blinnikov (1977), claimed that the magnetic field grows only over an accretion time scale or so and saturates at values well below equipartition. In support of their view they provided an exact  solution of the induction equation with the Poynting-Robertson term included that indeed exhibited the claimed saturation in the ideal MHD limit, while in an Appendix to their paper they argued that their results would not change significantly for finite plasma conductivity.    In our opinion, the reason for the difference in the conclusions of the two treatments lies in the consideration of the inner boundary condition to the solution of essentially the same equation. In addition, the one-dimensional character of the model used, although it showed that unrestricted field growth can take place, did not help the reader to visualize the actual geometry of the magnetic field on the poloidal plane. We therefore decided to return to this problem and perform more sophisticated two-dimensional numerical simulations with a variety of boundary conditions which, we hope, will convince the reader that the mechanism described in our original paper can produce astrophysically interesting results. In \\S~2 we describe the setup of our simulations and in \\S~3 we present the results of our calculations. Finally, in \\S~4 our results are discussed and some conclusions are drawn. ",
        "conclusions": "We presented above  a two-dimensional analysis of the cosmic battery mechanism first proposed in CK, paying particular attention to the criticism of Bisnovatyi-Kogan, Lovelace \\& Belinski~(2002); while still simplified, our present treatment represents a significant improvement over our earlier work, mainly because of its two-dimensional character which is essential in capturing the structure and evolution of the magnetic field on the poloidal plane.   The two-dimensional models confirm our previous results. When we use appropriate values of the plasma conductivity and the proper magnetospheric boundary conditions, these models produce configurations in which the magnetic flux interior to the accretion disk's inner edge $r_{\\rm in}$ increases linearly with time. At the same time, the present results complement our previous investigation and delineate conditions under which such field growth is possible. We have thus found that, for steady magnetic-field growth near the central object of an accretion disk, we need large diffusivity in the disk and a region near the central object in which the diffusivity practically drops to zero. The zero-diffusivity condition in the central region proves to be necessary because then the magnetic flux that crosses into radii $r < r_{\\rm in}$ cannot diffuse back out and becomes effectively trapped near the central object (a finite diffusivity in the central region would eventually allow for the flux to leak out to $r>r_{\\rm in}$, saturating thus the field growth, as was argued in the Appendix of Bisnovatyi-Kogan, Lovelace \\& Belinski~[2002]). We believe that the zero-diffusivity condition is reasonable in the central region where the accretion disk does not extend and the associated dissipative processes are thus not operative.  In our present simplified picture, we have ignored the transition region between the conducting interior (where $\\eta$ is effectively zero and the field is greatly amplified from a field a fraction of a Gauss to a field six or more orders of magnitude higher) and the diffusive exterior. In reality, any small amount of diffusivity present will `smooth out' the abrupt field transition. Magnetic flux will continue to accumulate as long as the thickness of the transition region remains smaller than $r_{\\rm in}$. Were the thickness of the transition region to become of the same order as $r_{\\rm in}$, the mechanism would then saturate, and the field would diffuse outward accross $r_{\\rm in}$ at the same rate that it is brought in. If we take into account the Spitzer collisional diffusivity estimated as $\\eta_{\\rm Spitzer}\\sim 10^3 T_6^{1.5}\\mbox{cm}^2\\mbox{s}^{-1}$, where $T_n$ is the plasma temperature in units of $10^n$~K (e.g.~Zombeck~1992), the field growth would saturate when the ratio $B(r<r_{\\rm in})/B(r>r_{\\rm in})$ reaches a value of the order of $v_r r_{\\rm in}/\\eta_{\\rm Spitzer}> 10^{12}$. We conclude that the Spitzer resistivity does not place any practical limit to the growth of the field interior to $r_{\\rm in}$.    Finally, we conclude that this battery mechanism could be important in accounting for the observed magnetic fields in several astrophysical sites as outlined in CK. Around a black hole in particular, our mechanism generates a magnetic field of the right large--scale dipolar topology required for the Blandford-Znajek mechanism of electromagnetic energy extraction from a spinning black hole to work (Blandford \\& Znajek~[1977]; Blandford~[2002]). In this respect we would like to point the reader's attention to the intriguing possibility that these considerations may in fact apply to understanding the low frequency cyclic variability observed in the galactic microquasar GRS 1915+105 (e.g. Pooley \\& Fender [1998]; Ueda {\\em et al.} [2002]). A distinguishing characteristic of this source are its X-ray/IR/radio flares that repeat over time scales of $\\sim 20-40$ minutes and lead to outflows with  $\\Gamma \\simeq 2-3$. It is interesting to note that if magnetic fields of order $B_z \\simeq 0.1$ G generated by the mechanism proposed herein in one accretion time of $\\sim$0.01 sec (in a model with $M=10M_\\odot$, $L=0.1 L_{\\rm Eddington}$, and $\\alpha=0.1$) were to increase linearly with time as envisioned in this work, then after a period of $\\sim$1000 seconds they would achieve values which could easily account for the observed flares and outflows.        DK would like to acknowledge support by INTEGRAL and HST GO grants."
    },
    "0608/astro-ph0608192_arXiv.txt": {
        "abstract": "We use high-resolution $N$-body simulations to investigate the optical depth of giant arcs with length-to-width ratio larger than 7.5 and 10 in the `standard' $\\LCDM$ model with $\\sigma_8=0.9$ and $\\Omega_{\\rm m,0}=0.3$ and a model based on three-year Wilkinson Microwave Anisotropy Probe (WMAP) data. We find that, in dark-matter only simulations, the lensing probability in the three-year WMAP model (with $\\sigma_8=0.74$ and $\\Omega_{\\rm m,0}=0.238)$ decreases by a factor of $\\sim 6$ compared with that in the `standard' $\\LCDM$ model. The effects of baryonic cooling, star formation and feedbacks are uncertain, but we argue that baryons will only increase the  the lensing cross-section by a moderate factor, $\\sim 2$. We conclude that the low central value of $\\sigma_8$ and $\\Omega_{\\rm m,0}$ preferred by the WMAP three-year data may be too low to be compatible with observations if conventional assumptions of the background source population are correct. ",
        "introduction": "In the hierarchical scenario of structure formation, structures in the universe are assumed to have grown from tiny quantum fluctuations generated during an inflation period through gravitational instability. Inflationary models predict that the initial fluctuations are Gaussian and have a roughly scale-free power-spectrum: $P(k) \\propto k^n$ with $n \\sim 1$. The amplitude of the power spectrum, for which a reliable prediction from inflationary theory is still lacking, has to be determined from observations. This amplitude is usually represented by a quantity, $\\sigma_8$, which is the {\\it RMS} of the linear density perturbations within a spherical window of radius $8h^{-1}\\mpc$, where $h$ the Hubble constant in units of 100$\\kms\\mpc^{-1}$.  The linear power spectrum can be probed with a variety of observations, ranging from  abundances of clusters, weak gravitational lensing, large-scale structure in galaxy distribution, galaxy motions, Lyman-$\\alpha$ forests, and most importantly,  the temperature fluctuations in the cosmic microwave background (see Spergel et al. 2006, and references therein). Recent results from the three-year data of the Wilkinson Microwave Anisotropy Probe (WMAP) alone prefers a $\\sigma_8$ value of  $0.74^{+0.06}_{-0.05}$. Although the errorbar is still quite large, the median value is lower than the value, $\\sigma_8=0.9$, adopted in the `standard' $\\LCDM$ model. The lower value of $\\sigma_8$ is favored by a number of other observations related to the clustering of galaxies (e.g. Jing, Mo \\& B\\\"orner 1998; Yang et al. 2004, 2005; van den Bosch, Mo \\& Yang 2003; Tinker et al. 2006; S\\'{a}nchez et al. 2006), but weak lensing results (e.g. \\citealt{Bacon02,Van Waerbeke02} ) and data on Lyman alpha forests (Viel, Haehnelt \\& Lewis 2006; Seljak, Slosar \\& McDonald 2006) seem to prefer a higher value.  As the largest bound structures in the universe, clusters of galaxies provide a sensitive probe for $\\sigma_8$, because their abundance in the universe depends sensitively on the amplitude of the density perturbations (and also the matter content in the universe, $\\Omega_{\\rm m,0}$). The abundance of clusters can be determined from X-ray (e.g. \\citealt{Rosati02,Pop04, RB02, R06}) and optical surveys (e.g. \\citealt{Gladders06}). However, in order to use the observed abundance to constrain model of structure formation, one has to know the masses of these objects accurately. This is in general difficult to do.  For example, X-ray studies usually assume hydrostatic equilibrium in order to derive the cluster mass, which may be invalid for many (such as merging) clusters (e.g. Gao \\& White 2006). Another probe of the abundance of clusters is  giant arcs that are produced when background galaxies are tangentially stretched by foreground clusters (e.g. \\citealt{Luppino99, ZG03, Gladders03, Sand05}). A number of recent investigations made comparisons between the observed abundance of giant arcs with that expected in the `standard' $\\LCDM$ model with $\\sigma_8=0.9$ (\\citealt{Dalal04}; \\citealt{Li05}; \\citealt{Hen05}; \\citealt{Horesh05}), finding agreement between model and observation (\\citealt{Dalal04, Wambsganss04}). In light of the change in cosmological parameters preferred by  the WMAP three-year data, it becomes important to re-evaluate how the predicted lensing probability changes.  In this paper, we study the number of giant arcs expected in the new cosmological model, using high resolution $N$-body simulations. We compare the results with those obtained from similar simulations of the standard $\\Lambda$CDM model, and with current observations. The plan of the paper is as follows. In \\S2, we describe the simulations we use and the analysis to predict giant arcs in the $N$-body simulations. Our main results are presented in \\S3, and we finish with a discussion in \\S4. ",
        "conclusions": "We have compared, using  high-resolution $N$-body simulations, the lensing probabilities of giant arcs in two cosmological models: the old `concordance' model with $\\sigma_8=0.9$, and the model with parameters given by the three-year WMAP data. We find that the lensing probability decreases by a factor of $\\sim 6$ in the 3-year WMAP model compared to that in the $\\LCDM$0 concordance model. This decrease is largely a result of the much lower number density of massive haloes in the WMAP3 model (see Fig. 1).  There has been a long debate whether the number of observed giant arcs is consistent with theoretical predictions. Early comparisons use simple analytical models (\\citealt{Kormann94, WH93, Wu96}; see also \\citealt{Oguri03}), which were later shown to severely under-estimate the lensing optical depth (Bartelmann \\& Weiss 1994; Bartelmann, Steinmetz \\& Weiss 1995; Meneghetti et al. 2003a; Torri et al. 2004). Many recent studies use N-body simulations, but with only dark matter (e.g. \\citealt{Wambsganss04}; \\citealt{Dalal04}, see \\citealt{Puchwein05} for an exception). Below we discuss the uncertainties in both observations and predictions, and examine the consistency found in several previous works in the $\\LCDM$0 cosmology.  \\subsection{Consistency in the $\\LCDM$0 cosmology}  Several previous studies concluded that the observational results of giant arcs are compatible with theoretical predictions (\\citealt{Oguri03, Wambsganss04, Dalal04}) in the $\\LCDM$0 cosmology. However, there are uncertainties in such conclusion.  The study by \\citet{Oguri03} used the axis ratio distributions in the tri-axial model of Jing \\& Suto (2002), and assumed a central profile, $\\rho \\propto r^{-1.5}$, that are steeper than the \\citet{NFW97} form.  Gas cooling can steepen the halo density profile, but it also makes the central mass distribution rounder. In this case, the shape distributions obtained in dark matter-only simulations are not suitable for clusters where gas cooling  steepens the density profile dramatically. Since a shallower profile or a more spherical shape will reduce the giant arc cross-sections, it is unclear whether or not the conclusion of \\citet{Oguri03} still holds if consistent models for density profile and halo shape are used. \\citet{Wambsganss04} assumes that the $L/W$ ratio can be approximated by the magnification. This assumption is valid for isothermal spheres, but not for real clusters which appear to be well fit by the \\citet{NFW97} profile (\\citealt{vdM00, Com06, VF06}). This assumption over-estimates the optical depth by a factor of few (\\citealt{Dalal04, Li05}). Furthermore, they also assumed a slightly higher normalization, $\\sigma_8 (0.95)$, which also increases the optical depth. The study by \\citet{Dalal04} assumes that arcs can be approximated as rectangles, which may overestimate the total cross-section by a factor of two compared to the perhaps more realistic assumption of elliptical arcs. In summary, even in the $\\LCDM$0 cosmology, the consistency between observations and predictions requires somewhat optimistic assumptions. The WMAP3 model, which predicts an optical depth by a factor of $\\sim 6$ smaller, makes it even harder to explain the observed giant arcs.  \\subsection{Effects of baryons}  In the real universe, baryons account for roughly 20\\% of the total mass, which can cool (form stars) and sink toward the centres of clusters. The radiative cooling likely has two effects: it will increase the concentration of baryons at the center of clusters, and at the same time, make the clusters more spherical (e.g. \\citealt{Dubinski94, Kaz04}). The former increases while the latter decreases the lensing cross-sections, and so the overall influence depends on which effect dominates.  Puchwein et al. (2005) studied the lensing cross-sections of clusters with $M \\ga 10^{15}h^{-1} M_\\odot$ with different prescriptions of numerical viscousities. They concluded that the higher concentrations due to baryons dominate and the baryons increase the lensing cross-sections by a factor of $\\sim 2$. However, they cautioned that their clusters may suffer from overcooling, as the stellar density in the core is larger than observed, and so the effect of baryons may be over-estimated. Notice also that their simulations did not include feedbacks from active galactic nuclei, which may further decrease the cooling of baryons at cluster centres and the corresponding lensing cross-sections. As the star formation treatment in hydrodynamical simulations is uncertain, \\citet{Mene03b} adopted an alternate empirical approach. They studied the effect of baryons by including a central cD galaxy in clusters. They found that the increase in the lensing cross-section due to cD galaxies is again quite moderate, by a factor of $ \\la 2$. Such an  enhancement cannot compensate the reduction in the optical depth due to the lower abundance of clusters in the WMAP3 model compared with the $\\LCDM$0 model.  \\subsection{Uncertainties in the cluster and background source populations}  The cluster abundance (but not internal structures, particularly when star formation is included) can now be reliably predicted from both numerical simulations and the extended Press-Schechter analytic formalism. In the $\\LCDM$0 cosmology, various investigations, using similar assumptions of the background source population and definition of the length-to-width ratio, predict lensing optical depths that agree with one another within a factor of 1.5 (e.g. \\citealt{Dalal04, Li05}; see \\citealt{Li05} for a detailed discussion). We caution, however, that even our $300 h^{-1}\\mpc$ box size still appears to be somewhat too small to sample the tail of the cluster mass function well, as is shown by the large noise in the right panel of Fig. 2. We have 158 clusters with $M>10^{14} h^{-1}M_\\odot$ (7 clusters with $M>3.8\\times10^{14}h^{-1} M_\\odot$). At redshift $z=0.3$ for $\\zs=3$, 50\\% of the giant arc cross-section is contributed by clusters with $M>3.8\\times10^{14}h^{-1} M_\\odot$. The larger noise may also be partly due to our inadequate sampling of merger events as we only dump the simulation data with a redshift interval of $d z \\approx 0.1$. As shown by a number of previous works (Torri et al 2004; Hennawi et al. 2005; Fedeli et al. 2006), merger events can boost the lensing cross-sections significantly. To examine the importance of this effect, we ran a lower resolution simulation in the WMAP3 model which evolves $N_{\\rm DM}=512^3$ dark matter particles in a box with sidelength of $600 h^{-1} \\mpc$. Due to the larger volume, the new simulation better samples the high mass tail of the cluster mass function and merger events. The optical depth decreases by a factor of $\\sim 2$. The larger optical depth in our first WMAP3 simulation is partly because it has relatively more massive clusters than the larger simulation due to cosmic variance. Furthermore, the poorer resolution of the larger simulation may have also reduced the optical depth, and so the real change should be $\\la 2$.  Substructures along the line of sight may be important for the anomalous flux ratio problem (\\citealt{Metcalf05}), but are unlikely to be important for the lensing cross-sections as they are expected to  contribute only a few per cent of the surface density. Notice that we use a cube of $4h^{-1}\\mpc$ to evaluate the lensing cross-sections.  In \\citet{Li05} the sidelength is chosen to be $2r_{\\rm vir}$. These two choices give almost identical optical depths (compare Fig. 2 with Fig.7 in \\citealt{Li05}). This demonstrates that the influence of matter (including substructures) in the outer skirts of clusters is not important.  A much larger uncertainty concerns the background source population, including their ellipticity, size and redshift distributions. While the former two appear to have modest effects on the lensing cross-sections (e.g. \\citealt{Oguri02, Li05}), the source redshift distribution has quite dramatic effects on the lensing cross-section (Wambsganss et al. 2004), as can be seen in Fig. \\ref{fig:tau}. This is the biggest uncertainty in the predictions of giant arcs. More redshift measurements of giant arcs will be particularly useful to clarify the situation (see \\citealt{Cov06} for a recent effort).  Currently the giant arc samples are still small. The largest dedicated search for giant arcs in X-ray clusters was performed by Luppino et al. (1999) who found strong lensing in 8 out of 38 clusters. In the optical, Zaritsky \\& Gonzalez (2003) found two giant arcs using the Las Campanas Distant Cluster Survey between redshift 0.5 and 0.7 in an effective search area of 69 square degrees while Gladders et al. (2003) found eight clusters with giant arcs using the Red-Sequence Cluster Survey over $\\sim 90$ square degrees. The largest arc sample found from HST archives was presented by Sand et al. (2005) with 116 arcs from 128 clusters, although its selection function is likely heterogeneous. Clearly the current giant arc samples are somewhat limited, although future weak lensing surveys (e.g. \\citealt{Wittman06}) may yield a large number of giant arcs as a by-product. A combined analysis of strong and weak lensing will be particularly useful for constraining $\\sigma_8$ and the cluster inner mass profiles.  Until larger arc samples become available and our understanding of the background source population is improved, it is difficult to reach firm conclusions concerning the consistency of the WMAP three-year model with observations. Nevertheless, it appears difficult to reconcile the giant arc statistics with the low central $\\sigma_8$ value (0.74) preferred by the WMAP three-year data."
    },
    "0608/astro-ph0608471_arXiv.txt": {
        "abstract": "We present semi-analytic dynamical models for giant molecular clouds evolving under the influence of HII regions launched by newborn star clusters. In contrast to previous work, we neither assume that clouds are in virial or energetic equilibrium, nor do we ignore the effects of star formation feedback. The clouds, which we treat as spherical, can expand and contract homologously. Photoionization drives mass ejection; the recoil of cloud material both stirs turbulent motions and leads to an effective confining pressure. The balance between these effects and the decay of turbulent motions through isothermal shocks determines clouds' dynamical and energetic evolution. We find that for realistic values of the rates of turbulent dissipation, photoevaporation, and energy injection by HII regions, the massive clouds where most molecular gas in the Galaxy resides live for a few crossing times, in good agreement with recent observational estimates that large clouds in local group galaxies survive roughly $20-30$ Myr. During this time clouds remain close to equilibrium, with virial parameters of $1-3$ and column densities near $10^{22}$ H atoms cm$^{-2}$, also in agreement with observed cloud properties. Over their lives they convert $5-10\\%$ of their mass into stars, after which point most clouds are destroyed when a large HII region unbinds them. In contrast, small clouds like those found in the solar neighborhood only survive $\\sim 1$ crossing time before being destroyed. ",
        "introduction": "Giant molecular clouds (GMCs) are the primary reservoirs of molecular gas within spiral galaxies.  Star formation is tightly correlated with the molecular column density within spiral galaxies \\citep{wong02}, and is therefore controlled by the formation and evolution of these giant clouds. In this and a subsequent paper we develop the theory of GMC dynamics and present semi-analytical models for GMC evolution.  We rely on simplifying assumptions about the structure of the cloud and the properties of the surrounding interstellar medium in order to focus on clouds' global energy budget and dynamical state.  Observations show that stars form much more slowly than the free-fall rate in GMCs \\citep{zuckerman74,rownd99,wong02,gao04a,wu05}, and any successful GMC model must explain why this should be. It is now widely held that collapse is inhibited primarily by intensely supersonic motions \\citep{vazquezsemadeni03,maclow04}, rather than magnetic fields \\citep{mouschovias76,shu87,mckee89}. (Star formation is strongly suppressed in low-extinction regions of molecular clouds, beyond what one would expect if the star formation rate were simply following the column density to some power, suggesting that magnetic fields may play a secondary role -- \\citealt{onishi98,onishi99,onishi02,johnstone04} -- though see \\citealt{hatchell05} who find that low column densities reduce but do not completely prevent star formation.) Simulations and analytic theory indicate that the observed level of turbulence in GMCs is sufficient to produce the observed rate of star formation \\citep{krumholz05c}.  However, undriven supersonic turbulence decays via radiation from isothermal shocks with an $e$-folding time of roughly one cloud crossing time \\citep{maclow98, stone98, maclow99, padoan99}, so undriven turbulence alone is not sufficient to prevent global collapse. Instead, GMCs must either be destroyed before their turbulent motions decay, or the turbulence must be continually driven. The mode of destruction is intimately related to the clouds' dynamical state.  Unless a cloud is destroyed all at once, any internal agent of destruction is also an internal source for turbulent energy -- and one strong enough to balance turbulent decay \\citep{matzner02}.  Destruction from within therefore favors models that achieve energetic and dynamical equilibrium \\citep[e.g.][]{mckee89}, if only briefly.  The alternative -- that clouds are disrupted entirely by outside agents \\citep[e.g.][]{bonnell06b} -- requires most of the cloud mass to remain gravitationally unbound, as bound regions rapidly collapse to higher density and become impervious to external forces. This is very difficult to reconcile with observational estimates of GMCs lifetimes and ratios of kinetic to potential energy. We critique the hypothesis of unbound clouds in greater detail in \\S~\\ref{discussion}.  In the following sections, we investigate the properties of molecular clouds both stirred and destroyed by HII regions within the cloud volume. The models we present here are improved in several ways relative to prior work: \\begin{itemize} \\item[-] Rather than enforcing strict mechanical or energetic equilibrium, we solve for the time evolution of a cloud's radius and turbulent velocity dispersion according to the time-dependent virial and energy equations. In an early discussion of this problem, \\citet{mckee89} allowed for time dependence in the energy equation, but assumed virial equilibrium. \\citet{matzner99a} and \\citet{matzner99b} followed the evolution of GMCs using time dependent virial and energy equations, but neglected several terms in these equations that we model here. While our approach is still not a full numerical solution of the equations of gravity, radiation, and magnetohydrodynamics, this approach enables us to study cloud dynamics without ignoring the effects of feedback on GMC evolution, as most numerical simulations to date have done \\citep[e.g.][]{clark05}. \\item[-] We account self-consistently for the recoil of cloud matter from the sites of mass ejection.  In addition to driving turbulence, inward recoil confines the cloud.  Recoil confinement is equivalent to an additional (and variable) external pressure, which becomes dynamically significant when cloud destruction is rapid. Corresponding terms appear in the virial (\\S~\\ref{equationofmotion}, Appendix \\ref{virialderivation}) and energy (\\S \\ref{energyeqn}, Appendix \\ref{energyderivation}) equations. \\item[-] Our model for HII regions (\\S \\ref{hiiregions}) accounts for the scale-dependence of density and velocity structures within GMCs. Although this does not fully account for the three-dimensional structure of the turbulent cloud medium, it is a significant improvement over the uniform cloud model employed by \\cite{whitworth79}, \\cite{williams97}, and \\cite{matzner02}. \\item[-] We apply the \\cite{krumholz05c} prescription for the star formation rate within turbulent clouds.  This formula accurately predicts the star formation rate on a variety of scales, from starburst galaxies to the dense precursors of individual star clusters \\citep{tan06a, krumholz06c}.  We use it to govern the birth rate of ionizing associations (\\S~\\ref{starformation}). \\item[-] Our dynamical simulations (\\S~\\ref{method}) track the formation and evolution of many individual HII regions.  This approach accounts for the finite lifetime of ionizing stars and the time delay associated with the deceleration of shells driven by HII regions, neither of which is negligible compared to GMC dynamical times. \\end{itemize}  We analyze the results of our models in \\S~\\ref{results}, discuss the implications of our findings in \\S~\\ref{discussion}, and summarize our conclusions in \\S~\\ref{conclusions}. In a future work (Matzner, Krumholz, \\& McKee, 2006, in preparation, hereafter Paper II) we apply this model to the problem of GMC formation and evolution in the galactic environment, including the effects of spiral density waves.  We do make several approximations in our work (\\S~\\ref{limitations}). We assume that the clouds are spherical and that they expand or contract homologously. We assume that the clouds are sufficiently massive that the energy injection is dominated by HII regions, not protostellar outflows, which based on the models of \\citet{matzner02} should be true for clouds of mass $\\sim 10^5$ $\\msun$ or more. Such clouds contain most of the molecular mass in the Galaxy. We neglect the possibility that the column density of the cloud must exceed a threshold in order for stars to form \\citep[e.g][]{mckee89}. We neglect energy  injection by HII regions after they reach pressure equilibrium with the surrounding medium. Finally, we neglect possible time dependent effects due to the ambient medium of the GMC: no mass can be added to the cloud, and the ambient pressure remains constant. Despite these approximations, the models we present in this Paper illustrate the degree to which GMC properties can be understood in terms of internal dynamics.  Obviously, real GMCs are not homologously expanding or contracting spheres with smooth density distributions, so the approximations we make to render the problem analytically tractable are quite limiting. For this reason one might be tempted give up on analytic treatment altogether and simply attempt to solve the problem numerically. Unfortunately, full numerical simulation of a giant molecular cloud, including the formation of multiple star clusters and the effects of their feedback, is not feasible with current codes and computers. Instead, one is forced to make numerous approximations regardless of whether one takes a numerical or analytic approach. Many numerical simulations of GMC evolution simply ignore feedback altogether \\citep[e.g.][]{clark05}, include it only from a single source and/or focus on size scales much smaller than an entire GMC \\citep[e.g.][]{dale05,li06b}, or focus on the galactic scale and lack the resolution to say anything about individual GMCs \\citep[e.g.][]{slyz05, tasker06}. For this reason, an analytic approach that allows us to include feedback provides a valuable complement to numerical results, and points out areas for future simulation in which new effects might be discovered by a more thorough treatment of the physics. ",
        "conclusions": "\\label{conclusions}  In this Paper we derive the basic equations governing the evolution of the mass, radius, and velocity dispersion of giant molecular clouds under the approximation of homologous motion. We construct simple models for the rate at which turbulent motions decay due to radiation from isothermal shocks and for the rate at which HII regions driven by massive stars within the cloud drive turbulent motions, and we use these to study the global evolution and energetic balance of GMCs. This enables us to build GMC models in which we neither assume energetic and virial balance, nor neglect the effects of feedback. Thus, we are able for the first time to address critical questions such as whether GMCs are in virial or energetic equilibrium, how long they live, what determines their lifetimes and star formation histories, and what determines properties such as column density and virial parameter that appear to be roughly the same for all GMCs.  Our primary conclusion is that giant molecular clouds with observed properties are indeed quasi-equilibrium objects. The rate at which they lose energy via radiation from isothermal shocks is roughly balanced by the rate at which HII regions from star formation inject energy back into the cloud. This feedback keeps them virialized, $\\avir = 1-2$, and keeps them at the observed column density $N_H\\approx 1.5\\times 10^{22}$ cm$^{-2}$. Our results suggest that this column density, which appears to be generic for GMCs both in the Milky Way and in other galaxies where GMCs are observable, is in fact the result of feedback, since the efficiency of star formation feedback varies with column density, and the observed column density corresponds to the one required for equilibrium. Whether the GMC formation process, particularly the passage through an overdense and overpressured spiral shock, also plays a role in selecting the GMC column density is at this point an open question, one we will address in Paper II.  The duration of the equilibrium state of a GMC depends on cloud mass, but for clouds $\\gtsim 10^6$ $\\msun$ in mass it lasts for $2-3$ crossing times, or roughly 30 Myr. The dominant destruction mechanism for GMCs is dynamical unbinding by the momentum delivered by an expanding HII region, but for massive clouds this unbinding does not happen until HII regions have photoionized away $\\sim 90\\%$ of the GMC mass. Smaller clouds are dynamically disrupted by HII regions in $\\sim 10$ Myr while $\\sim 50\\%$ of their mass is still in molecular form, but it is unclear if this disruption actually destroys their molecular material or simply breaks it into smaller, mutually unbound clouds.  Overall, our models produce a picture of GMCs that is quite different than that suggested by models in which GMCs are gravitationally unbound or in which feedback from star formation is neglected. Models that neglect feedback appear incapable of reproducing the observed lifetimes and properties of the GMC population. We conclude that any reasonable model of GMC evolution cannot neglect the effects of star formation feedback, and that even a simple model of feedback such as the one we use here produces good agreement with the observed properties of GMCs."
    },
    "0608/astro-ph0608647_arXiv.txt": {
        "abstract": "We have analyzed HST/NICMOS2 F110W, F160W, F165M, and F207M band  images covering the  central 1\\mbox{$^\\prime$}$\\times$1\\mbox{$^\\prime$} region of the cluster associated with Mon R2 in order to constrain the Initial Mass Function (IMF)  down to  20 M$_\\mathrm{jup}$. The flux ratio between the F165M and F160W bands  was used to measure the strength of the water band absorption feature and select a sample of 12 out of the total sample of  181 objects that have effective temperatures between 2700 K and 3300 K. These objects are placed in the HR diagram together with sources observed by \\citet{carpenter} to estimate an  age of $\\sim$1 Myr for the low mass cluster population. By constructing extinction limited samples, we are able to constrain the IMF and the fraction of stars with a circumstellar disk in a sample that is 90\\%\\ complete for both high and low mass objects. For stars with estimated masses between 0.1 M$_\\odot$ and 1.0 M$_\\odot$ for a 1 Myr population with A$_\\mathrm{V} \\le 19$ mag, we find that $27\\pm9\\%$  have a near--infrared excess indicative of a circumstellar disk. The derived  fraction is similar to, or slightly lower than, the fraction found in other  star forming regions of comparable age. We constrain the number of stars in the mass interval 0.08--1.0 M$_\\odot$ to the number of objects in the mass interval 0.02--0.08 M$_\\odot$ by forming the ratio, $R^{**}$=N(0.08--1 M$_\\odot$)/N(0.02--0.08 M$_\\odot$) for objects in an extinction limited sample complete for A$_\\mathrm{V}\\le 7$ mag. The ratio is found to be  $R^{**}=2.2\\pm1.3$ assuming an age of 1 Myr, consistent with  the similar ratio predicted by the system IMF proposed by \\citet{chabrierreview}. The ratio is  similar to the ratios observed towards the  Orion Nebula Cluster and IC 348 as well as the ratio derived in the 28 square degree survey of Taurus by \\citet{guieu}. ",
        "introduction": "The shape of the IMF and whether or not it is  universal have been central  questions in astrophysics for more than 50 years. The shape of the stellar part of the IMF has been examined observationally in some detail. Above 0.5 M$_\\odot$, the IMF derived for field stars is well characterized by a Salpeter power law, $dN/dM \\propto M^{-\\alpha_1}, \\alpha_1=2.35$ \\citep{kro02}. At this characteristic mass, the slope of the power law changes to $\\alpha_2=1.35$ and for masses  below 0.1 M$_\\odot$ the IMF flattens further \\citep{kro02,chabrierreview}, although the slope is very uncertain \\citep{allen}. To search for variations in the IMF as a function of environment, one has to look at individual star--forming episodes, such as star clusters or associations. Young clusters in particular are useful for these studies because the more massive stars are still present in the cluster and the clusters have not experienced significant dynamical mass segregation. Recently, the IMF studies have been extended deep  into the brown  dwarf regime in several star forming regions, most notably  the Orion Nebula Cluster \\citep[ONC, e.g.][]{carphil,muench,slesnick}, IC 348 \\citep[e.g.][]{najita00,luhman2}, and Taurus \\citep[e.g.][]{briceno,luhman,guieu}.   In order to search for variations in the sub--stellar IMF as a function of environment, a larger ensemble of young clusters is needed. Here, we present a study of the IMF in the Mon R2 cluster. The embedded cluster associated with Mon R2 was originally discovered by \\citet{beckwith} and is  located at a distance of 830$\\pm$50 pc \\citep{herbst}. % The most massive star in the cluster is $\\sim$ 10 M$_\\odot$ \\citep{carpenter,massi}. \\citet{carpenter} estimated  Mon R2 to contain at  least 475 stars within a 3.6$\\times$3.6 pc region, significantly less than the $\\sim$ 3500 stars found within a 2.5 pc radius of the ONC \\citep{hillenbrand98}. IC 348 contains $\\sim$ 400 members  including brown dwarfs \\citep{luhman2}. Thus, Mon R2 has a richness  intermediate between the  two well studied clusters. The molecular cloud has been shown by \\citet{choi} to have a central density of $n_c\\sim10^7\\mathrm{cm}^{-3}$, much higher than typical densities in the Taurus molecular cloud cores ($n_c\\sim10^5\\mathrm{cm}^{-3}$, see e.g. \\citet{onishi}). If the density of the molecular cores have an impact on the IMF as suggested by e.g. \\citet{batebonnell} and \\citet{goodwin}, we would expect the IMF in Mon R2 to have relatively more brown dwarfs than Taurus.  Previously, \\citet{carpenter} imaged   Mon R2  down to the brown dwarf limit and obtained spectra for a subset of the stars. They established the  age of the cluster to be $\\le$ 3 Myr, and  that the  majority of  the low mass  stars  clustered around the  1 Myr isochrone. Adopting an age of 1 Myr for the stellar population, they found the ratio of high mass to low mass stars N(1--10 M$_\\odot$)/N(0.1--1 M$_\\odot$)$=0.1$, which is   consistent with other nearby star forming regions \\citep{meyer00} and  with a \\citet{miller} field star IMF. Despite the fact that Mon R2 is relatively nearby, the brown dwarf content  has not  been investigated in a systematic manner. Here, we present HST/NICMOS 2 near--infrared imaging of the central 0.24pc$\\times$0.24pc (1\\mbox{$^\\prime$} $\\times$ 1\\mbox{$^\\prime$}) of the cluster. Our main goal is to constrain the low--mass stellar and  brown dwarf content in the cluster and compare the results to those  derived for  other  young clusters and the field. In addition, we estimate the fraction of stars with disks within our surveyed region.   The paper is structured as follows: In Section 2, we present the observations, the data reduction and the photometry. Section 3 presents the  results from the color--magnitude and color--color diagrams for the cluster. We  place objects in the HR diagram based on the estimate of the effective temperature derived from the F160W and F165M band observations. In Section 4, we  discuss our constraints on the IMF and disk frequency based on  extinction limited samples constructed from our observations. % Section 5 compares our derived ratios with the similar ratios for other star forming regions. Finally, we conclude in Section 6 with a summary of our results. ",
        "conclusions": "We have presented the results from HST/NICMOS2 F110W, F160W, F165M, and F207M band imaging of  the inner 1$\\arcmin\\times1\\arcmin$ of the  embedded cluster associated with  Mon R2. Our results are as follows: \\begin{itemize} \\item{The effective temperature has been estimated for a small set of stars in the temperature range $2700-3300$ K based on a water vapor index and these objects have been placed in the HR diagram. A 1 Myr \\citet{baraffe} isochrone is consistent with the lower mass objects placed in the HR diagram in agreement with \\citet{carpenter}.} \\item{We find that for stars with spectral type M6 or earlier (0.1--1M$_\\odot$) a disk fraction of $27\\pm9$\\% based on an extinction limited sample of 43 stars A$_\\mathrm{V} \\le 19$ mag. } \\item{We created three extinction limited samples complete for A$_\\mathrm{V}=13$, 10, and 7 mag, and containing 34, 19, and 13 objects, respectively. We have calculated the  three ratios of low mass stars to brown dwarfs $R=N(0.08-1.0 \\mathrm{M}_{\\odot})/(0.04-0.08 \\mathrm{M}_\\odot)=10.3\\pm5.8$, $R^{*}=N(0.08-1.0\\mathrm{M}_{\\odot})/N(0.03-0.08 \\mathrm{M}_{\\odot})=8.5\\pm6.4$ and $R^{**}=N(0.08-1.0 \\mathrm{M}_{\\odot})/N(0.02-0.08 \\mathrm{M}_{\\odot})=2.2\\pm1.3$. } \\item{The derived ratios are consistent with the similar ratios for  Taurus, IC 348, the ONC, and the system field IMF of \\citet{chabrierreview}. There is thus no compelling evidence for variations in the relative brown dwarf content between Mon R2 and other nearby star forming regions. } \\end{itemize}"
    },
    "0608/astro-ph0608121_arXiv.txt": {
        "abstract": "{} {We have observed the 6030 and 6035 MHz transitions of OH in high-mass star-forming regions to obtain magnetic field estimates in both maser emission and absorption.} {Observations were taken with the Effelsberg 100 m telescope.} {Our observations are consistent with previous results, although we do detect a new 6030 MHz maser feature near $-70$~km~s$^{-1}$\\ in the vicinity of W3(OH).  In absorption we obtain a possible estimate of $-1.1 \\pm 0.3$~mG for the average line-of-sight component of the magnetic field in the absorbing OH gas in K3-50 and submilligauss upper limits for the line-of-sight field strength in DR 21 and W3.} {These results indicate that the magnetic field strength in the vicinity of OH masers is higher than that of the surrounding, non-masing material, which in turn suggests that the density of masing OH regions is higher than that of their surroundings.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608317_arXiv.txt": {
        "abstract": "We discuss the possible appearance of strange exotic multi-quark states in the interior of neutron stars and signals for the existence of strange quark matter in the core of compact stars. We show how the in-medium properties of possible pentaquark states are constrained by pulsar mass measurements. The possibility of generating the observed large pulsar kick velocities by asymmetric emission of neutrinos from strange quark matter in magnetic fields is outlined. ",
        "introduction": "In the spirit of this session on exotic quark states, strange quark matter and astrophysics, we are going to discuss the interrelation of these seemingly different topics. First, we give a brief discussion on multiquark states, in detail we address the properties of multiquark states with strangeness and charm. We consider light pentaquarks with strangeness then heavy pentaquarks and multi-quark state with charm. We summarise the current understanding of the in-medium properties of pentaquarks in cold dense matter and extrapolate it to the case of neutron star matter in beta-equilibrium. The possible appearance of various pentaquark states and its implications for the global properties of neutron stars is analysed and confronted with recent pulsar mass measurements.  In the second part of this talk, we present newly suggested signals for the possible presence of strange quark matter in the core of compact stars. In particular, we focus on the phenomenon of pulsar kicks, which has received considerable attention recently. The recent detailed measurements of velocities of nearby pulsars in our galaxy clearly shows, that speeds in excess of 1000 km/s are involved, much higher than the typical velocity of ordinary nearby stars. So far a generally accepted explanation for this effect is not at hand. In addition, we briefly comment on recent developments on the relation of gamma-ray bursts with strange quark matter and the recent gravitational wave signals derived for the presence of strange quark matter in compact stars. ",
        "conclusions": "We have discussed multiquark states, with strangeness and with charm, and used pulsar mass measurements to derive constraints on the in-medium properties of light pentaquark states. Multiply charmed multiquark states have unusual properties and can be bound by strong colour-electric forces. Double prong decay patterns can signal their existence, if bound, and can be measured with the future tracking devices for charmed particles of the PHENIX and STAR detectors at RHIC. One should not miss this oportunity and look for it! Light pentaquarks, i.e.\\ five-quark states with strangeness, can be present in the core of neutron stars and reduce the maximum mass. Hence, their in-medium properties can be constrained by pulsar mass measurements. Present limits give only loose bounds, which become much tighter with only a slightly increased pulsar mass limit. We are looking forward to ongoing and future pulsar scans and continued observations of binary pulsars which will reduce the uncertainty considerably in the near future!  The presence of strange quark matter in compact stars can signal itself in various astrophysical observables. The puzzling observation of high velocity pulsars might be related to the presence of strange quark matter in pulsars. The conditions for a successful neutrino rocket, the asymmetric emission of neutrinos due to polarised electrons in magnetic fields, have been examined in detail.  Large magnetic fields are needed, which might be present in the quark core of compact stars. Enough kick is reached if the asymmetric emission of neutrinos starts at high enough temperature which is only reached in the early evolution of the proto-neutron star stage, i.e.\\ within a minute after the supernova collapse. At these temperatures, however, the neutrino mean-free path is shorter than the radius of the compact star, in particular in neutron matter. Sufficiently large values are reached for colour-superconducting quark matter in the CFL phase only, if effects from a massless excitation, the so called H, are ignored. However, then the specific heat and the available thermal energy is reduced exponentially so that not enough kick can be produced. Hence, our present understanding of high-density QCD and colour-superconductivity does not give large enough mean-free paths or sufficiently large heat capacities to allow for a successful neutrino rocket. More detailed investigations are needed to pursue this issue in more detail.  The burst pattern of gamma-ray bursters with long quiescence times might indicate the presence of two separate phase transitions in compact stars: one from nuclear matter to normal quark matter then the second one to colour-superconducting quark matter. Also gravitational wave patterns can exhibit the presence of strange quark matter in compact stars either from compact star mergers or from collapsing neutron stars.  Last but not least much more astrophysical data is on the horizon poised for discovery. There are lots of opportunities for strange quark matter physics. Certainly, the cross-talk with observers will be crucial and vital for a successful advancement of strange astrophysics!    \\ack This work is supported in part by the Gesellschaft f\\\"ur Schwerionenforschung mbH, Darmstadt, Germany."
    },
    "0608/astro-ph0608584_arXiv.txt": {
        "abstract": "We have redefined group membership of six southern galaxy groups in the local universe (mean cz $< 2000 \\; \\mbox{km} \\, \\mbox{s}^{-1}$) based on new redshift measurements from our recently acquired Anglo--Australian Telescope 2dF spectra. For each group, we investigate member galaxy kinematics, substructure, luminosity functions and luminosity-weighted dynamics. Our calculations confirm that the group sizes, virial masses and luminosities cover the range expected for galaxy groups, except that the luminosity of NGC 4038 is boosted by the central starburst merger pair. We find that a combination of kinematical, substructural and dynamical techniques can reliably distinguish loose, unvirialised groups from compact, dynamically relaxed groups.  Applying these techniques, we find that Dorado, NGC 4038 and NGC 4697 are unvirialised, whereas NGC 681, NGC 1400 and NGC 5084 are dynamically relaxed. ",
        "introduction": "Galaxy groups are important structures, containing a significant fraction of the galaxy population in the universe. In November 2004 (ATAC/2004B/19) and April 2005 (PATT/2005A/13), we observed six nearby southern galaxy groups (Dorado, NGC 681, NGC 1400, NGC 4038, NGC 4697 and NGC 5084) with the 2dF multi-object spectrograph at the 3.9-m Anglo--Australian Telescope. The primary purpose of these observations was to locate ultra-compact dwarf galaxies (UCDs). In this paper we combine non-UCD galaxy redshift measurements from these observations with previously catalogued data to create revised group membership lists. We investigate luminosity-weighted group kinematics and dynamics based upon the virial method employed in \\citet{Ferguson..1990}, using Monte Carlo modelling to estimate uncertainties due to observational errors.  A meaningful kinematical and dynamical analysis of galaxy groups requires a clear definition of group membership -- but galaxy groups are variously defined depending on the group detection algorithm selected. Researchers have historically relied on positional, morphological and surface brightness information to define galaxy groups, but increasing access to redshift data has improved group membership definition. For example, \\citet{Huchra..1982} applied a ``friends of friends'' percolation algorithm to search for iso-density contours in redshift space, and \\citet{Maia..1989} used the Southern Sky Redshift Survey to produce a Catalog of Southern Groups of Galaxies. \\citet{Ferguson..1990} used galaxy morphology, together with available but limited redshifts, to create catalogs of nearby galaxy groups, including Dorado and NGC 1400. \\citet{Garcia..1993} combined the percolation method of \\citet{Huchra..1982} with a hierarchical method derived from \\citet{Tully..1987} to create a large catalogue of galaxy groups, including NGC 681, NGC 4038, NGC 4697 and NGC5084. Starting with group membership as defined by \\citet{Ferguson..1990} and \\citet{Garcia..1993}, our results more accurately define these six galaxy groups as gravitationally-bound systems in redshift space.  The structure of our paper is as follows -- Section (2) describes our 2dF observations and redshift results obtained; Section (3) details the methods, results and interpretation we have used in our analysis of group kinematics, substructure and luminosity-weighted dynamics; and Section (4) summarises our findings. ",
        "conclusions": "Our key findings are as follows:  \\begin{itemize} \\item Target selection by morphological parameters (image area and peak intensity against magnitude) has resulted in 12 to 22 per cent of the successfully measured redshifts being identified as group member galaxies. \\item We have redefined the non-dwarf galaxy membership of the six groups with more accurate redshift--velocity measurements, eliminating several background galaxies. However, significant numbers of less luminous galaxies in Dorado and NGC 1400 groups still require redshift measurements. \\item Several techniques confirm that the six galaxy groups can be separated into loose group (Dorado, NGC 4038 and NGC 4697) and compact group (NGC 681, NGC 1400  and NGC 5084) subsets in both kinematical and dynamical terms. Kinematical techniques include the comparison of sky position plots converted to a common group distance, and the Dressler-Shectman substructure test. Dynamical tests (luminosity-weighted) include comparison of harmonic mean radii at a common group distance, and testing for virialisation through the crossing time. \\item We find substructure to be evident from the Dressler-Shectman test in Dorado, NGC 4038 and NGC 4697. A more refined application of the crossing time test for virialisation also confirms these three groups to be still undergoing dynamical relaxation. The luminosity-weighted virial mass estimate and mass-to-light ratio for these groups may therefore be questionable. \\end{itemize}"
    },
    "0608/astro-ph0608067_arXiv.txt": {
        "abstract": "The spin temperature of neutral hydrogen, which determines the 21 cm optical depth and brightness temperature, is set by the competition between radiative and collisional processes.  In the high-redshift intergalactic medium, the dominant collisions are typically those between hydrogen atoms.  However, collisions with electrons couple much more efficiently to the spin state of hydrogen than do collisions with other hydrogen atoms and thus become important once the ionized fraction exceeds $\\sim 1 \\%$.  Here we compute the rate at which electron-hydrogen collisions change the hydrogen spin.  Previous calculations included only $S$-wave scattering and ignored resonances near the $n=2$ threshold.  We provide accurate results, including all partial wave terms through the $F$-wave, for the de-excitation rate at temperatures $T_K \\la 1.5 \\times 10^4 \\kel$; beyond that point, excitation to $n \\ge 2$ hydrogen levels becomes significant.  Accurate electron-hydrogen collision rates at higher temperatures are not necessary, because collisional excitation in this regime inevitably produces \\lya photons, which in turn dominate spin exchange when $T_K \\ga 6200 \\kel$ even in the absence of radiative sources. Our rates differ from previous calculations by several percent over the temperature range of interest.  We also consider some simple astrophysical examples where our spin de-excitation rates are useful. ",
        "introduction": "\\label{intro}  The 21 cm transition is potentially a powerful probe of the pre-reionization intergalactic medium (IGM) because of the enormous amount of neutral hydrogen in the Universe at that time \\citep{field58, scott90, madau97}.  It can teach us about reionization, the formation of the first structures and the first galaxies, and even the ``dark ages\" before these objects formed (\\citealt{furl06-review}, and references therein).  It is therefore crucial to understand the fundamental physics underlying the 21 cm transition.  One critical aspect is the spin temperature, which is determined by the competition between the scattering of cosmic microwave background (CMB) photons, the scattering of \\lya photons \\citep{wouthuysen52, field58}, and collisions.  When CMB scattering dominates, the IGM remains invisible because the spin temperature approaches that of the CMB (which is used as a backlight).  Before star formation commences, collisions are the only way to break this degeneracy.  The total coupling rate is determined by collisions both with other hydrogen atoms and with electrons.  At the low residual electron fraction expected after cosmological recombination \\citep{seager99}, H--H collisions dominate.  Spin exchange in such interactions has received a great deal of attention over the years \\citep{purcell56, smith66, allison69, zygelman05, hirata06}.  Such collisions suffice to couple the spin temperature $T_S$ to the kinetic temperature $T_K$ at high densities, when $z \\ga 50$ \\citep{loeb04}.  However, the spin exchange rate in these collisions falls dramatically at $T_K \\la 100 \\kel$; coupled with the declining density, such interactions become ineffective at lower redshifts. H--e$^-$ collisions have received much less attention. However, at fixed temperature, free electrons move at much higher velocities than do hydrogen atoms, and are thus much more efficient at spin exchange on a per-particle basis.  If the mean ionized fraction is $\\bxion \\ga 0.01$, they can play a significant role in the coupling. This situation can occur if, for example, X-rays from the first quasars or star-forming galaxies create a warm, partially-ionized medium \\citep{oh01, venkatesan01, ricotti04_a, ricotti05, nusser05-xray, kuhlen05-sim, kuhlen06-21cm}, or if the decay or annihilation of exotic particles deposits a significant amount of energy into the IGM \\citep{chen04-decay, pierpaoli04, shchekinov06, furl06-dm}.  The spin de-excitation rate in H--e$^-$ collisions, which we will denote $\\kappa_{10}$, obviously depends on the detailed scattering properties of these collisions.  Over the past several decades, this interaction has received a great deal of attention as a model problem for electron-atom interactions (e.g., \\citealt{massey51, schwartz61, temkin62, burke62a, poet78, bray02}).  At IGM temperatures $T_K \\la 10^4 \\kel$, most electrons have energies below the threshold to excite the $n=2$ level, and the dominant mechanism of spin de-excitation is electron-electron spin exchange.  Although this regime has been particularly well-studied, it has only been solved in full nonrelativistic detail relatively recently.  The latest spin exchange calculation \\citep{smith66} used the scattering phase shifts computed numerically by \\citet{schwartz61}.  While accurate, these included only $S$-wave scattering and neglected both higher-order partial waves and H$^-$ resonances near the $n =2$ threshold (which were not resolved by the numerical algorithm). Both of these issues have been solved over the intervening four decades (see, e.g., \\citealt{wang93, wang94} for a recent calculation).  Our goal is to recalculate $\\kappa_{10}$ in light of these newer phase shifts.  Higher-order partial waves become significant at $T_K \\ga 1000 \\kel$; if X-ray heating in the early Universe is strong, then these temperatures can be achieved easily \\citep{furl06-glob}.  We therefore examine the high-temperature limit in some detail.  Our calculation breaks down once collisional electronic excitation (and ionization) become important.  We will show that, above $T_K \\sim 1.5 \\times 10^4 \\kel$, direct spin de-excitation by scattering above the $n=2$ threshold cannot be ignored.  Although such interactions are difficult to model -- because of the wide array of possible transitions -- in practice this regime is relatively unimportant. In the low-density IGM, collisional excitation to the $2p$ and higher levels is followed rapidly by radiative de-excitation, producing a \\lya background.  We show that this background inevitably dominates the spin temperature coupling at $T_K \\ga 6200 \\kel$ (assuming a thermal distribution of electrons).  The crossover can occur at smaller temperatures if a non-thermal population of fast electrons exists (produced, for example, by X-rays; \\citealt{chen06, chuzhoy06-first}).  The remainder of this paper is organized as follows.  In \\S \\ref{21cm}, we briefly review the 21 cm transition.  Our main results are contained in \\S \\ref{ehcoll}, where we examine various mechanisms for spin de-excitation, consider the high-temperature limit, and calculate the spin de-excitation rates for H--e$^-$ collisions.  We present some simple astrophysical applications in \\S \\ref{astro}, and we conclude in \\S \\ref{disc}.  In our numerical calculations, we assume a cosmology with $\\Omega_m=0.26$, $\\Omega_\\Lambda=0.74$, $\\Omega_b=0.044$, $H=100 h \\hunits$ (with $h=0.74$), $n=0.95$, and $\\sigma_8=0.8$, consistent with the most recent measurements \\citep{spergel06}, although we have increased $\\sigma_8$ from the best-fit \\emph{WMAP} value in order to improve agreement with weak-lensing data. ",
        "conclusions": "\\label{disc}  We have computed the spin de-excitation rates for hydrogen in the ground state during electron-hydrogen collisions.  Previous calculations assumed that spin exchange through elastic $S$-wave scattering was the only relevant mechanism.  Our analysis showed that while spin exchange does dominate spin de-excitation, higher partial waves contribute significantly to the total de-excitation rate.  We used newer calculations of the elastic scattering phase shifts for the $L=0$--$3$ partial waves \\citep{wang94} to re-compute $\\kappa_{10}$.  Our main results are presented in Table~\\ref{tab:kappa}.  They remain accurate up to $T_K \\la 1.5 \\times 10^4 \\kel$, where collisions above the $n=2$ threshold begin to become important at the percent level.  Our results differ from \\citet{smith66} and especially from the widely-used fit in \\citet{liszt01} by several percent over the range $1 \\kel < T_K < 10^4 \\kel$.  We have also shown that spin coupling produced directly through H--e$^-$ collisions will become a secondary effect at $T_K \\ga 6200 \\kel$, because \\lya photons produced through collisional excitations easily dominate the spin coupling in this regime.  Although the excitation rate to the $2p$ level (and to other levels that cascade through Ly$\\alpha$) can be much smaller than the collisional spin de-excitation rate, each \\lya photon scatters $\\sim 10^5$ times before redshifting out of resonance.  This dramatically boosts the efficiency of the radiation background relative to collisions.  In practice, the \\lya background may be important at even lower temperatures.  In our calculations, we have assumed a Maxwell-Boltzmann distribution of electron velocities, but even a small population of nonthermal high-energy electrons can induce significant \\lya coupling.  For example, X-rays produce fast secondary electrons able to excite hydrogen collisionally, which can easily produce such a background \\citep{chen06, chuzhoy06-first}.  We also considered some simple applications of our results to the high-redshift IGM.  Electron-hydrogen collisions are likely to be most important if the radiation background is dominated by X-rays, so that the IGM becomes warm and weakly ionized.  Such scenarios are particularly common when black hole accretion provides a significant fraction of the ionizing flux \\citep{ricotti04_a, ricotti05, kuhlen06-21cm}.  We have seen that  the typical 21 cm fluctuation amplitude can reach the milliKelvin level even without a significant ultraviolet background.  Our rate coefficients will be useful in assessing the implications of such scenarios for the 21 cm sky.  \\vspace{0.1cm}  SRF thanks the Tapir group at Caltech for their hospitality while much of this work was completed.  This publication has been approved for release as LA-UR-06-5395.  Los Alamos National Laboratory, an affirmative action/equal opportunity employer, is operated by Los Alamos National Security, LLC, for the National Nuclear Security Administration of the U.S. under contract DE-AC52-06NA25396."
    },
    "0608/astro-ph0608298_arXiv.txt": {
        "abstract": "The host stars of extrasolar planets tend to be metal-rich. We have examined the data for these stars for evidence of trends in other galactic parameters, without success. However, several ESP hosts are likely to be members of the thick disk population, indicating that planet formation has occurred throughout the full lifetime of the Galactic disk. We briefly consider the radial metallicity gradient and age-metallicity relation of the Galactic disk, and complete a back-of-the envelope estimate of the likely number of solar-type stars with planetary companions with $6 < R < 10$ kpc. ",
        "introduction": " ",
        "conclusions": "More than 130 stars are now known to harbour planetary systems. The correlation between planetary frequency and metallicity is now well established, and probably reflects nature rather than nurture. We have analysed the statistical properties of the current sample of ESP hosts, looking for other possible trends and biases. With the possible exception of a higher mean velocity perpendicular to the Plane, the planetary hosts appear to be unremarkable members of the Galactic Disk.  Several ESP host stars show enhanced $\\alpha$-element abundances and high velocities relative to the Sun, strongly suggesting that they are members of the thick disk. If so, this indicates that planet formation has been a constant presence in the Galactic Disk.  We briefly considered the radial abundance gradient and age-metallicity distribution of disk stars. Recent results suggest that the gradient is flatter in the vicinity of the Sun than the canonical value, but probably steepens at larger radii; there is relatively little reliable data for the inner disk, $R < 6$ kpc. The age-metallicity distribution of the Berkeley/Carnegie RV sample shows substantial dispersion at all ages. While the uncertainties in these parameters limit our ability to model the full Galaxy, we can use local statistics to estimate the planetary population in the vicinity of the Sun. We estimate that there are over $3.5 \\times 10^7$ `RV-detectable' planetary systems with Galactocentric radii in the range 6 to 10 kpc."
    },
    "0608/astro-ph0608251_arXiv.txt": {
        "abstract": "s{ For several years the only experiments sensitive to astrophysical gamma rays with energies beyond the reach of EGRET but below that of the Cherenkov imaging telescopes have been the \"solar tower\" detectors. They use $>2000$ m$^2$ mirror areas to sample the Cherenkov wavefront generated by $<100$ GeV gamma rays, obtaining Crab sensitivities of more than 6$\\sigma$ in one ON-source hour. I will review the history of the solar tower Cherenkov experiments from 1992 to the present and their key design features. I will describe some successful analysis strategies, then summarize the principal results obtained. \\\\ \\em{Contribution to the proceedings of \"Towards a Network of Atmospheric Cherenkov Detectors VII\", 27-29 April 2005 - Ecole Polytechnique, Palaiseau, France.} http://polywww.in2p3.fr/actualites/congres/cherenkov2005/ } ",
        "introduction": " ",
        "conclusions": "In 1992 we thought that opening the 100 GeV energy range, between EGRET and Whipple/CAT/Hegra, would give fast access to the hundreds of sources seen by EGRET. Now, over a decade later, HESS has shown us that the key is {\\em not} just to shift in energy range but especially to improve in sensitivity. The many sources we hoped to see are indeed there, but at the $< {1 \\over 20}$ Crab level rather than $> {1 \\over 5}$. In 1992 we also thought that wavefront sampling might have advantages to complement imagers at low energies. HESS has since confirmed that imaging is {\\em very} powerful.  The solar tower experiments have made the first flux meaurements of the three canonical Cherenkov sources in this new energy window (Crab, Mrk 421, Mrk501). Upper limits (weak detections?) for a few interesting blazars, amongst which are W Com and 1ES 1426+428, have been obtained. M31 provided a rare constraint on astrophysical SUSY. Our Crab pulsar limit is likely to remain the best until GLAST is launched.  The solar tower experiments were successful in spite of a range of adverse conditions, ranging from the political (STACEE's efforts to convince Commonwealth Edison to let them stay at Solar-II were nearly as epic as VERITAS's fight to stay on Mt. Hopkins), to the environmental (the number of cloudy nights at Th\\'emis, but also at Sandia, has been getting consistently worse), and to the astrological (the conjunction of Saturn with the Crab was annoying).  A success not to be overlooked is the training acquired by our students and postdocs, some of whom have gone on to key roles in HESS and VERITAS, or other fields of physics. Pioneering the solar tower approach required deep thought about the atmospheric Cherenkov technique, and the ability to tackle tough experimental problems. Roughly 20 PhD theses have been written.  What future for wavefront sampling, especially with the new imager arrays making such a splash? The 2000 heliostats at Solar-II still make some of us dream, and the (small!) CACTUS team is implementing some new ideas, but imager arrays will surely be the linchpin of future detectors. Could imager sensitivity be enhanced by adding ASGAT-like telescopes? The limiting factors are i) primary electrons (or cosmic rays fragmenting almost completely to a $\\pi^o$) and ii) angular resolution. Timing arrays give a good measure of the shower core position, which is on average 1 radiation length higher in the sky for electrons than for gamma rays. There could be ideas to explore here."
    },
    "0608/hep-ph0608078_arXiv.txt": {
        "abstract": "We show that, as a result of non-linear self-interactions, it is feasible, at least in light of the bounds coming from terrestrial tests of gravity, measurements of the Casimir force and those constraints imposed by the physics of compact objects, big-bang nucleosynthesis and measurements of the cosmic microwave background, for there to exist, in our Universe, one or more scalar fields that couple to matter much \\emph{more strongly} than gravity does. These scalar fields behave like chameleons: changing their properties to fit their surroundings. As a result these scalar fields can be not only very strongly coupled to matter, but also remain relatively light over solar system scales. These fields could also be detected by a number of future experiments provided they are properly designed to do so. These results open up an altogether new window, which might lead to a completely different view of the r\\^{o}le played by light scalar fields in particle physics and cosmology. ",
        "introduction": "} There is wide-spread interest in the possibility that, in addition to the matter described by the standard model of particle physics, our Universe may be populated by one or more scalar fields. These are a general feature in high energy physics beyond the standard model and are often related to the presence of extra-dimensions, \\cite{morecham3}.  The existence of scalar fields has also been postulated as means to explain the early and late time acceleration of the Universe \\cite{sn1a,inflation,quint,quint1,quint2,easson}. It is almost always the case that such fields interact with matter: either due to a direct Lagrangian coupling or indirectly through a coupling to the Ricci scalar or as the result of quantum loop corrections, \\cite{damour,carroll1,carroll2,anupam1,anupam2}.  If the scalar field self-interactions are negligible, then the experimental bounds on such a field are very strong: requiring it to either couple to matter much more weakly than gravity does, or to be very heavy \\cite{uzan,bounds,damour1,damour2}. Recently, a novel scenario was presented by Khoury and Weltman \\cite{cham1} that employed self-interactions of the scalar-field to avoid the most restrictive of the current bounds.  In the models that they proposed, a scalar field couples to matter with gravitational strength, in harmony with general expectations from string theory, whilst, at the same time, remaining very light on cosmological scales.   In this paper we will go much further and show, contrary to most expectations, that the scenario presented in \\cite{cham1} allows scalar fields, which are very light on cosmological scales, to couple to matter much \\emph{more} strongly than gravity does, and yet still satisfy \\emph{all} of the current experimental and observational constraints.  The cosmological value of such a field evolves over Hubble time-scales and could potentially cause the late-time acceleration of our Universe \\cite{chameleoncosmology}.  The crucial feature that these models possess is that the mass of the scalar field depends on the local background matter density. On Earth, where the density is some $10^{30}$ times higher than the cosmological background, the Compton wavelength of the field is sufficiently small as to satisfy all existing tests of gravity.  In the solar system, where the density is several orders of magnitude smaller, the Compton wavelength of the field can be much larger.  This means that, in those models, it is possible for the scalar field to have a mass in the solar system that is much smaller than was previously thought allowed.   In the cosmos, the field is lighter still and its energy density evolves slowly over cosmological time-scales and it could function as an effective cosmological constant. While the idea of a density-dependent mass term is not new~\\cite{added,pol,others,others1,others2,others3}, the work presented in \\cite{cham1,chameleoncosmology} is novel in that the scalar field can couple directly to matter with gravitational strength. If a scalar field theory contains a mechanism by which the scalar field can obtain a mass that is greater in high-density regions than in sparse ones, we deem it to possess a \\emph{chameleon mechanism} and be a \\emph{chameleon field theory}.   When referring to chameleon theories, it is common to refer to the scalar field as the \\emph{chameleon}.   We start this article by reviewing the main features of scalar field theories with a chameleon mechanism. Afterwards, this paper is divided into roughly two parts: in sections \\ref{sing}, \\ref{effmacr} and \\ref{forcetwo} we study the behaviour of chameleon theories as field theories, and derive some important results. From sections \\ref{exper} onwards, we combine these results with a number of experimental and astrophysical limits to constrain the unknown parameters of these chameleon theories $\\{\\beta, M, \\lambda\\}$. We shall show how the non-linear effects, identified in sections \\ref{sing}-\\ref{forcetwo}, allow for a very large matter coupling, $\\beta$, to be compatible with all the available data.  We also note that some laboratory based tests of gravity need to be redesigned, if they are to be able to detect the chameleon. If the design of these experiments can be adjusted in the required way, and their current precision maintained of its current level, then a large range of sub-Planckian chameleon theories could be detected, or ruled out, in the near future.  In section \\ref{sing}, we study how $\\phi$ behaves both inside and outside an isolated body and derive the conditions that must hold for such a body to have a thin-shell.  In this section, we show how non-linear effects ensure that the value that the chameleon takes far away from a body with a thin-shell is \\emph{independent} of the matter-coupling, $\\beta$. Whilst such $\\beta$-independence as been noted before for $\\phi^4$-theory in \\cite{nelson}, this is the first time that it has been shown to be a generic prediction of a large class of chameleon theories.  In section \\ref{effmacr} we show the internal, \\emph{microscopic}, structure of macroscopic bodies can unexpectedly alter the \\emph{macroscopic} behaviour of the chameleon. Using the results of section \\ref{sing} and \\ref{effmacr} we are then able to calculate the $\\phi$-force between two bodies; this is done in section \\ref{forcetwo}.  In each of these sections we take care to note precisely when linearisation of the chameleon field equation is invalid.  Laboratory bounds on chameleon field theories are analysed in section \\ref{exper}.  We focus mainly on the E\\\"{o}t-Wash experiment reported in \\cite{EotWash, Eotnew}, which tests for corrections to the $1/r^2$ behaviour of gravity, and experimental programmes that measurement the Casimir force \\cite{Lamcas,Decca,othercas}. We also look at the variety of laboratory and solar-system based tests for violations of the weak equivalence principle (WEP) \\cite{LLR, WEP,WEP1,WEP2}.  The extent to which proposed satellite-based searches for WEP violation will aide in the search for scalar fields with a chameleon-like behaviour is considered in this section.  We shall see that for a large range of values of $M$ and $\\lambda$, laboratory tests of gravity at unable to place \\emph{any} upper-bound on $\\beta$.  In sections \\ref{compact} and \\ref{cosmo} we show how the stability of white-dwarfs and neutron stars, as well as requirements coming from big bang nucleosynthesis and the Cosmic Microwave Background, can be used to bound the parameters of chameleon field theories.  We shall see that such considerations do result in upper-bounds on $\\beta$.  Finally, in sections \\ref{allbounds} and \\ref{conclude} we collate all of the different experimental and astrophysical restrictions on chameleon theories, use them to plot the allowed values of $\\beta$, $M$ and $\\lambda$, and discuss our results and their implications.  We include a summary of the main results at the end of sections \\ref{sing}-\\ref{forcetwo} for easy reference. This allows the reader, less interested on the detailed derivation of the formulae, to follow the whole article. Throughout this work we take the signature of spacetime to be $(+\\,-\\,-\\,-)$ and set $\\hbar = c =1$; $G= 1/M_{pl}^2$ where $M_{pl}$ is the Planck mass. ",
        "conclusions": "}  In this article we have investigated scalar field theories in which the field is strongly coupled to matter. In particular, we have studied the so called chameleon scalar fields. A scenario presented by Khoury and Weltman \\cite{cham1} that employed self-interactions of the scalar-field to avoid the most restrictive of the current bounds on such fields and its coupling. In the models that they proposed, these fields would couple to matter with gravitational strength, in harmony with general expectations from string theory, whilst, at the same time, remaining very light on cosmological scales. In this work we went  much further and show, contrary to most expectations, that the scenario presented in \\cite{cham1} allows scalar fields, which are very light on cosmological scales, to couple to matter much \\emph{more} strongly than gravity does, and yet still satisfy \\emph{all} of the current experimental and observational constraints.   Previous investigations on such scenarios, \\cite{chama, chameleoncosmology, nelson,phi4}, noted that an important feature of chameleon field theories is that they make unambiguous and testable predictions for near-future tests of gravity in space. This is timely as there are currently four satellite experiments either in the proposal stage or due to be launched shortly (SEE~\\cite{SEE}, STEP~\\cite{STEP}, GG~\\cite{GG} and MICROSCOPE~\\cite{MICRO}).  A reasonably sized region of the parameter space of the chameleon theories considered here will be visible to these missions. Theories with very large couplings $\\beta \\gg 1$ will, however, remain undetectable.   The ability of these planned missions to detect large $\\beta$ theories could, however, be exponentially increased if the experiments they carry were to be redesigned slightly in the light of our findings.  Previous studies claimed that typical test masses in the above satellite experiments do not have a thin shell. Therefore, the extra force is comparable to their gravitational interaction. The chameleon model hence predicts that MICROSCOPE, STEP and GG could measure violations of the weak equivalence principle that are stronger than currently allowed by laboratory experiments. Furthermore, the SEE project could measure an effective Newton's constant different, by order of unity, from that measured on Earth.  We have seen, in this paper, that both of these features are very much properties of chameleon theories with runaway ($n>0$) potentials. They will not, in general, occur if the chameleon potential has a minimum (e.g. $n<-4$ theories). These features are also very much associated with a gravitational strength chameleon coupling i.e. $\\beta \\sim {\\mathcal{O}}(1)$, and will not occur if $\\beta \\gg 1$, or $\\beta \\ll 1$.  The major result, presented in this work, is that current experiments do \\emph{not} limit the coupling of the chameleon to matter, $\\beta$, to be order ${\\mathcal{O}}(1)$ or smaller. Indeed, if we wish to have a `natural' value of $M$  in a $V\\propto \\phi^{-1}$ theory then we must require $\\beta \\gtrsim 10^{4}$ (or $\\beta \\lesssim 10^{-3}$.  If $\\beta \\gg 1$, the test-masses in the planned satellite experiments will still have thin-shells. As such SEE, STEP, GG and MICROSCOPE, as they are currently proposed,  will be \\emph{unable} to detect the chameleon and place an upper-bound on $\\beta$.  We have shown that upper-bounds on the matter coupling, $\\beta$, can be derived from astrophysical and cosmological considerations. Also, if $\\beta$ is very large, of the order of $10^{17}$ or greater, then it might even be possible to detect the effect on the chameleon on scattering amplitudes in particle colliders.  This possibility is one avenue that requires further in depth study.  We noted in the introduction that $\\beta \\gg 1$ could be seen as being pleasant in light of the hierarchy problem \\cite{HP,HP1,HP2}. If a chameleon with a large $\\beta$ where detected, it would imply that new physics emerges at a sub-Planckian energy-scale: $M_{\\phi} = M_{pl}/\\beta$.  A large value of the matter coupling is also preferable to an order unity value in that it leads to the late time behaviour of the chameleon being much more \\emph{weakly} dependent on the initial conditions, than it would be if $\\beta \\lesssim {\\mathcal{O}}(1)$.  The magnitude of $\\beta (\\Delta\\phi = \\phi_1 - \\phi_2)/M_{pl}$ quantifies the relative amount by the particle masses differ between a region where $\\phi = \\phi_1$ and one where $\\phi=\\phi_2$. The larger the coupling is then, the easier it is for there to have been a very large difference between the current values of particle masses, and the values that they had in the very early universe (i.e. pre-BBN). If the particle masses were very different from their present values at, say, the epoch of the electroweak phase transition, then the predictions of electroweak baryogenesis could be significantly altered.  In this paper have taken the chameleon potential to have a power-law form i.e $V\\propto \\phi^{-n}$.   This is certainly \\emph{not} the only class of chameleon potentials that it is possible to have. In general, any potential that satisfies $\\beta V_{,\\phi} < 0$, $V_{,\\phi\\phi} > 0$ and $V_{,\\phi\\phi\\phi}/V_{,\\phi} > 0$ in a region near some $\\phi = \\phi_{0}$ will produce a chameleon theory.  In fact, a generic potential might contain many different regions in which chameleon behaviour is displayed. In some of these regions, the potential may appear to have a runaway form, and so behave qualitatively as an $n>0$ theory. In other regions, the potential might have a local minimum, leading to $n \\leq -4$ type behaviour.   The existence of the matter coupling provides one with a mechanism, along the lines of that considered by Damour and Polyakov \\cite{pol}, by which the scalar field $\\phi$ can, during the radiation era, be moved into a region where it behaves like a chameleon field.   The larger $\\beta$ is, the more effective this mechanism becomes. Given this mechanism, one important avenue for further study is to see precisely how general late-time chameleon behaviour is of a generic scalar field theory with a strong coupling to matter.  In this paper we have avoided the temptation to linearise the chameleon field equation, eq. (\\ref{micro}), when it is not valid to do so.  We have, instead, combined matched asymptotic expansions with approximate analytical, and exact numerical, solution of the full non-linear equations to study the behaviour of chameleon field theories in more detail.  The main results of this analysis were: \\begin{itemize} \\item We found the conditions under which a body would have a thin-shell, and noted that the development of a thin-shell is related to the onset of non-linear behaviour. \\item We have shown that the far field of a body with a thin-shell is \\emph{independent} of the coupling strength $\\beta$: this is a generic property of all the chameleon field theories with a power-law potential.  This $\\beta$-independence was seen, in section \\ref{exper}, to have important consequences for the design of experiments that search for WEP violations, and was seen to be vital in allowing theories with $\\beta \\gg 1$ to be compatible with the current experimental bounds. The $\\beta$ independence of the far-field of thin-shelled bodies was seen to cause the $\\phi$-force between two such bodies to be $\\beta$-independent also. \\item For $\\beta \\sim O(1)$ the best bounds on $M$ and $\\lambda$ currently come from the E\\\"{o}t-Wash experiment \\cite{EotWash,Eotnew}  For $\\beta \\gg 1$, the best bounds on $M$ and $\\lambda$ were found to come from measurements of the Casimir force. \\item Non-linear effects were shown to limit the magnitude of the average chameleon mass in a thin-shelled body to be smaller than some critical value, $m^{crit}_c$. Intriguingly, $m^{crit}_c$ is independent of $\\beta$, $M$ and $\\lambda$, and depend only on $n$ and the \\emph{microscopic} properties of the thin-shelled body. \\item The experimental constraints on the coupling of chameleon fields to matter are much weaker than those on non-chameleon fields.  In fact when $n \\neq -4$, the constraints on large $\\beta$ theories are weaker than those in which the scalar field couples to matter with gravitational strength ($\\beta \\sim O(1)$). Almost paradoxically, strongly-coupled scalar fields are actually harder to detect than weakly coupled ones. \\item Perhaps the most important result, though, is that the ability of table-top gravity tests to see strongly coupled, chameleon fields could be exponentially increased if certain features of their design could be adjusted in the appropriate manner.  The detection, or exclusion, of chameleon fields with $\\beta \\gg 1$ represents a significant but ultimately, we believe, achievable challenge to experimentalists. Searches for large $\\beta$ chameleon fields represent one way in which table-top tests of gravity could be used to probe for new physics beyond the standard model. Whether, or not, chameleon fields actually exist, it is important to note those areas into which our current experiments cannot see and, if possible, design experiments to probe those areas. \\end{itemize}  We have shown, in this paper, that scalar field theories that couple to matter much more strongly than gravity are not only viable but could well be detected by a number of future experiments provided they are properly designed to do so.  This result opens up an altogether new window which might lead to a completely different view of the r\\^{o}le played by light scalar fields in particle physics and cosmology."
    },
    "0608/astro-ph0608605.txt": {
        "abstract": "{We present the first-ever simulations of non-ideal magnetohydrodynamical (MHD) stellar magnetospheric winds coupled with disc-driven jets where the resistive and viscous accretion disc is self-consistently described.}{ These innovative MHD simulations are devoted to the study of the interplay between a stellar wind (having different ejection mass rates) and an MHD disc-driven jet embedding the stellar wind.}{The transmagnetosonic, collimated MHD outflows  are investigated numerically using the VAC code. We first investigate the various angular momentum transports occurring in the magneto-viscous accretion disc. We then analyze the modifications induced by the interaction between the two components of the outflow.}{Our simulations show that the inner outflow is accelerated from the central object's hot corona thanks to both the thermal pressure  and the Lorentz force. In our framework, the thermal acceleration is sustained by the heating produced by  the  dissipated magnetic energy due to the turbulence. Conversely, the  outflow launched from the resistive accretion disc is mainly accelerated by the magneto-centrifugal force.}{The simulations show that the MHD disc-driven outflow extracts angular momentum more efficiently  than do viscous effects in near-equipartition, thin-magnetized discs where turbulence is fully developed.  We also show that, when a dense inner stellar wind occurs, the resulting disc-driven jet has a different structure, namely a magnetic structure where poloidal magnetic field lines are more inclined because of the pressure caused by the stellar wind. This modification leads to both an enhanced mass-ejection rate in the disc-driven jet and a larger radial extension that is in better agreement with the observations, besides being more consistent.} ",
        "introduction": "%%%%%%%%%% Observations %%%%%%%%%%%%%   Accreting stellar objects are often associated with collimated jets or  winds from accretion  discs. Most of those objects also show evidence of winds originating in  a corona surrounding the central object. These accretion-ejection phenomena are observed in different astrophysical sources ranging from young stellar  objects (YSOs), X-ray binaries, planetary nebulae to active galactic nuclei (AGNs) (see, e.g. \\citep{Livio07} and references therein). The outflow provides an efficient way of extracting  angular momentum and converting  gravitational energy from the accretion disc or from the central object into outflow kinetic energy.  Observations show that most of the jets are launched very close to the central engine. In the case of YSOs, there is direct observational evidence \\citep{Burrowsetal96} as well as in the case of some microquasars \\citep{Fenderetal97, Mirabel03}. For instance, it has been suggested that, in microquasars, the fastest components of the outflow are launched in the vicinity of the black hole \\citep{meier03}. Another piece of evidence that the outflow may originate in a region relatively close to the central  object is that the observed asymptotic velocity of the jet is close to the escaping speed from the central engine. Thus there is a direct relation between the  asymptotic speed and the depth of the gravitational potential \\citep{Mirabel99, Livio99, Pringle93}. Moreover, the high-resolution images of $H_{\\alpha}$ and $[OI]$ \\citep{Bacciottietal02} reveal a continuous transverse variation of the jet's velocity, where the fastest and densest components are closer to the central  axis.\\\\ The high velocity of the observed jet in YSOs suggests that they originate in a region that is no larger than one astronomical unity (AU) in extent \\citep{KwanTademaru88} and between 0.3 to  4.0 AU from the star in the case of the LVC of DG Tau \\citep{Andersonetal03}. This theoretical prediction is supported, in the case of a disc wind, by  observations of the rotation of several jets associated with TTauris \\citep{Coffeyetal04}. Moreover, in the case of classical TTauris (CTTS), UV observations \\citep{Beristainetal01,Dupreeetal05} reveal the presence of a warm wind whose temperature  is at least of $3\\times 10^5 {\\rm K}$. It appears that the source of this wind is restricted to the star itself, where X-ray observations support the presence  of a hot corona in CTT stars \\citep{Feigelson&Montmerle99}. These observations also suggest the existence in CTTS of stellar winds comparable to the solar wind. These winds may be both thermally and magneto-centrifugally accelerated.  %%%%%%% Model & Simulation %%%%%%%% Since the discovery of the existence of winds and jets in astrophysics, enormous progress has been made regarding the understanding of these phenomena. At the same time, we still do not know precisely how the wind from the central corona of the star or the compact object interacts with the disc outflow and the  respective roles and differences between these two types of flows.\\\\ Several works have analytically and numerically studied the formation of outflows launched from the accretion disc  \\citep{BP82, Cao&Spruit94, ContopoulosLovelace94, Ustyugovaetal95, Ouyedetal97, VlahakisTsinganos98, CasseFerreira00, Casse&Keppens2002, Casse&Keppens2004, Andersonetal05, Pudritzetal06}. Other works have focused on the outflows  from the hot  corona of the central objects  \\citep{Sakurai85, SautyTsinganos94, Fendt03,  Matt&Balick04}. \\\\ In models dealing with outflows launched from accretion discs, the magnetic field plays a  key role  in the accretion, the acceleration, and the collimation of the associated vertical wind, which is also supported by recent observations  \\citep{Donatietal05}.  The detection of rotation signatures in TTauri jets gives extra strong support to the magneto-centrifugal launching from the accretion disc. However in stellar ouflows, the wind must be thermally accelerated because of the strong heating of viscous and non-ideal magnetohydrodynamical (MHD) mechanisms. This acceleration increases and approaches the magneto-centrifugal acceleration at least near the polar  axis. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Some models have already investigated diffusive disc-driven jet launching. In some simulations the accretion disc was considered as a fixed, time-independent boundary condition, while a constant magnetic resistivity prevails through the entire outflow \\citep{Fendt&Cemeljic02}. \\citet{Kuwa05} have included an accretion disc in their resistive simulations but only considered an uniform magnetic resistivity everywhere in their computational domain. Because of the very different physical conditions prevailing in the disk, the jet, and the external medium, it seems very unlikely that this kind of resistivity is physically relevant. %include an accretion disc in their resistive MHD simulations but %were not able to go beyond one inner disc rotation, which is too small a %timescale to investigate any flow dynamics. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Models involving two component bipolar outflows have been proposed in the case of AGN as such as in Sol et al. (1989) and  Renaud \\& Henri (1998), where an electron-positron central wind component is surrounded by an external, ideal MHD disc-driven jet. Another two-component outflow model has been proposed by \\citep{Bogovalov&Tsinganos05} where a relativistic pulsar wind is surrounded and collimated by an ideal MHD disc-driven wind. In the case of YSOs, two-component models were considered, such as  in X-wind outflows  \\citep{SautyTsinganos94, Ferreiraetal00}. The inner component extracts its energies from the corona around the central region (the central object and the inner part of the accretion disc where an  advection-dominated accretion flow may exist), while the second component is  launched from the thin  accretion disc.\\\\ %%%%%%%% Present the plan of the paper %%%%%%%% The aim of this paper is to investigate the formation of a two-component outflow around YSOs, one coming from the thin accretion disc and the other one injected from the hot corona of the central star. This work is developed on the base of the  disc wind simulations of Casse \\& Keppens (2002, 2004) (CK02, CK04). Its goal is to study the influence of the stellar wind on the structure and the dynamics of the jet around YSOs. Furthermore, we investigate the consequences of the  energy dissipation in the outflow close to the polar axis. Before that, we present simulations of the outflow launched from a resistive and viscous accretion disc where magnetic Prandtl number (ratio of the anomalous viscosity to the anomalous resistivity) equals unity.  This is the first time that viscosity and resistivity are considered together in the disc and included in an MHD simulation involving both the accretion disc and the related jet. In previous works, the viscous accretion disc was examined without taking the disc wind into account or with an imposed internal structure \\citep{VonRekowskietal06}, which does not enable the authors to study the complete angular momentum transfer. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Thus, the first part investigates the relative role of angular-momentum transport by viscosity and by the outflowing plasma and its influence on the formation of the outflow from the disc. Then in a second part, we present the results of our simulations of ideal MHD outflows launched from resistive, viscous, accretion discs surrounding the turbulent wind accelerated from the hot corona of the central star. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "In this paper we present numerical simulations of the interaction between an  accretion-ejection structure launching a disc-driven jet and a stellar wind emitted either from the central object and/or from its magnetosphere, in particular for the case of YSOs. In our framework, the accretion disc is near equipartition between thermal pressure and magnetic pressure where turbulence is believed to occur. This turbulence is characterized by a time and space-dependent anomalous  resistivity and viscosity set by using an $\\alpha$ description. The origin of the turbulence is still unknown but is not likely to arise from magneto-rotational instability since an equipartition disc is inconsistent with the development of such instability \\citep{Ogilvie&Livio01}. \\\\ The properties of both the accretion disc and  outflow were investigated in this paper.  In a first stage, we analyzed the contribution of the various hydrodynamical and magnetohydrodynamical mechanisms to the angular momentum transport in the thin accretion disc including, for the first time, both anomalous viscosity and resistivity, with a magnetic Prandtl  number equal to unity.  We demonstrated that the MHD Poynting flux associated with the disc-driven jet plays a major role in the removal of the angular momentum from the thin accretion disc. The angular momentum radial transport provided by the anomalous viscosity inside the  disc remains weak and contributes only  $1\\%$ of the total angular momentum transport (this value  agrees with analytical estimates depending on disc thickness and $\\alpha$ value). This is consistent with the viscous torque depending upon the radial derivative of angular velocity, while the magnetic torque is mainly controlled by the vertical derivative of the toroidal component of the magnetic field. In a thin accretion disc, the toroidal magnetic field varies from zero to an equipartition value on a disc scale height $\\epsilon R, \\varepsilon \\ll 1$, leading to a much more efficient extraction of the rotational energy from the magnetic torque into the MHD Poynting energy flux feeding the jet. Basically, with our simulation and despite the disc viscosity, we have reproduced very similar results to those obtained by CK04 where the viscous torque was neglected.\\\\ %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure*} {\\rotatebox{0}{\\resizebox{16cm}{5cm}{\\includegraphics{5484F15.ps}}}} \\caption{Plot of the temporal  evolution of the ejection  mass-loss rate from the accretion disc in a) the ratio of the stellar mass-loss rate to the ejection  mass-loss rate from the accretion disc as a function of time in b)  for the simulation with a stellar mass-loss rate of $\\dot{M}=10^{-7}M_\\odot/yr$.} \\label{Fig15} \\end{figure*} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure*} {\\rotatebox{0}{\\resizebox{18cm}{5.5cm}{\\includegraphics{5484F16.eps}}}} \\caption{The transverse variation of different physical quantities as magnetic-field components, velocity components, and the fast-magnetosonic Mach number %(left) the poloidal and toroidale magnetic field (right) vertical velocity at Z = 100, for the simulation with stellar mass loss $\\dot{M}=10^{-7}M_\\odot/yr$}. \\label{Fig16} \\end{figure*} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% In the second stage of the present paper, we studied the effects induced  by the launching of a stellar wind inside the hollow jet arising from the accretion disc. The stellar outflow is thermally driven by the turbulent viscous and resistive stresses in a region close to the polar axis. In order to mimic the non-ideal effects believed to occur inside the stellar wind, we prescribed anomalous viscosity and resistivity in the wind region, leading to a turbulent heating of the plasma near the polar axis.\\\\ We performed various simulations using different stellar mass ejection rates from the central objects. These stellar ejection-mass rates range from a Sun-like star ($\\dot{M}=10^{-9}M_{\\odot}/yr$) to O and B type stars ($\\dot{M}=10^{-7}M_{\\odot}/yr$). The influence of the stellar wind on the dynamics and the structure of jet and the accretion-ejection structure around the stellar object gets stronger with higher stellar mass ejection rates. As an example in the simulations where stellar mass ejection rates are $\\dot{M}=10^{-9}M_{\\odot}/yr$, we obtained a very similar disc-driven jet to the one in CK04. The only difference lies in the presence of the internal, fast, hot plasma coming from the central object. In this simulation, the ejection-mass rate in the disc-driven jet is similar to the one in CK04, on the order of $15\\%$ of the inner disc accretion mass rate, while the stellar ejection mass rate represents $1\\%$  of the total mass loss in the outflow.  In our simulations, the collimation of the stellar outflow takes place once the jet from the accretion disc is launched and has reached a significant spatial extension. Its collimation is provided mainly by pinching the toroidal magnetic field in equilibrium with the thermal pressure gradient. Conversely the collimation of the jet from the accretion disc is induced by the poloidal magnetic field pressure gradient balancing the centrifugal force. Furthermore, in all the simulation the stellar wind keeps having a more or less conical expansion up to the asymptotic region where the disc-driven jet acts to collimate the stellar flow into a cylindrical flow.  These important results are  self-consistently obtained, in contrast to simulations of \\citet{Bogovalov&Tsinganos05} where a relativistic wind was collimated by a jet but which did not consider either the accretion disc, the jet launching, or the stellar wind acceleration. Indeed in our model, we  self-consistent describe the launching and the collimation of disc and stellar wind. In particular, our simulations completely describe both the stellar flow acceleration (the stellar flow is injected with sub-Alfv\\`enic velocity) and the launching mechanism of the jet from the accretion disc.  We have shown that the contribution of  non-ideal MHD mechanisms in the acceleration of the stellar outflow can be significant since turning on this dissipative mechanism leads, for instance, to higher terminal velocities of the stellar jet-collimated flow. Our prescription of these dissipative mechanisms is of course subject to improvements, but our goal was to show that they  enable an increase in the efficiency of both the thermal and  magnetic acceleration of the stellar wind. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% The turbulence may be produced by the interaction of the stellar wind with the disc-driven jet. Moreover, as in the solar wind, the turbulence in the stellar wind may also have a stellar origin and/or  a possible connection to the accretion occurring near the surface of the star. In this scenario, a part of the energy released by accretion is carried away in the wind by outwardly propagating Alfv\\'en waves inducing  turbulence. This scenario is a variant of models where a significant part of the accretion energy is converted into thermal energy in the star corona (see Matt \\& Pudritz 2005 and references therein). By performing simulations with no stellar wind turbulence but with a large amount of thermal energy at the base of the wind, we found quite similar result except for the velocity field, the resistive continuous heating of the stellar wind being more efficient in providing higher velocity. It is noteworthy that, similarl to models depositing thermal energy near the stellar corona, the amount of energy released by the turbulent heating is a significant fraction of the accretion energy (in the particular case of our simulations, it represents near $35\\%$ of the accretion energy). %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Note also that similar double-layer jets were found by Koide and collaborators in  various simulations (e.g. Koide et al. 1998; Koide et al. 1999) but those simulations were devoted to rapidly variable jets (with only a few disc rotations) and not to steady structures.\\\\ We performed simulations with higher stellar ejection-mass  rate, typically with $\\dot{M}=10^{-7}M_{\\odot}/yr$ (compatible with O-B type stars). The increase in the stellar mass-loss rate induces a faster and larger expansion of the jet. Indeed the enhanced pressure provoked by the stellar wind tends to bend the disc magnetic field lines over a larger radial extension, leading to a larger disc-driven jet. The corresponding disc-driven jet mass-ejection rate is much higher than in previous simulations since it reaches $50\\%$ of the disc inner accretion rate (stellar ejection-mass rate on the order of $10\\%$). The simulations give a quantitative threshold beyond which the stellar jet gives a significant extra expansion of the disc jet. Typically a mass loss rate from the star on the order of $\\dot{M}=10^{-7}M_{\\odot}/yr$ gives a factor two in the radial expansion of the disc-driven jet. Although the total jet remains small in cross section, as in CK04, and compared to self-similar disc wind models. Moreover the stellar jet play a very important role, also, in the formation of the disc wind. %Note at this stage that We verify that,  part of the external disc wind may look quasi radially-self-similar in nature, the most inner part of the stellar wind is quasi meridionally self-similar. In our work we neglected all radiative losses coming from the central star or from the plasma itself. The implementation of these terms and the study of their impact on the outflow structure is postponed to future work."
    },
    "0608/astro-ph0608245_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\subsection{The status quo}  Visualization is a key tool in astronomy for discovery and analysis tool that is used throughout the discipline.  In observational astronomy, data display applications are used from the observation planning stage through data collection and reduction phases, to production of figures for journal articles.  In theoretical astrophysics, data display is an essential process in comprehending nearly every simulation data set, and in summarising results for publication.  Most existing astronomy display tools operate in the two-dimensional (2d) paradigm.  That is, multi-dimensional data is explicitly reduced to a data set having at most two ``look-up'' coordinates (or indices) into the data prior to being presented to a display device.   Examples of 2d display packages widely adopted by the astronomy community include the {\\sc DS9} image display tool\\footnote{{\\sc DS9} -- \\url{http://hea-www.harvard.edu/RD/ds9}}, the \\pgplot\\ programming library\\footnote{\\pgplot\\ -- \\url{http://www.astro.caltech.edu/~tjp/pgplot}}, and third-party commercial packages such as IDL\\footnote{IDL -- \\url{http://www.rsinc.com/idl}} and MONGO/SuperMONGO\\footnote{SuperMONGO -- \\url{http://www.astro.princeton.edu/~rhl/sm}}.  The dominance of the 2d display paradigm is a simple consequence of both the primary publishing medium -- paper -- and display device -- computer monitor -- being inherently flat and 2-dimensional. Additionally, the computational demands of 2d graphics display are typically much smaller than those of three-dimensional (3D) display. In this paper, by ``3D display'' we mean that a set of geometrical objects (or ``geometry'') is described -- by the programmer and/or software user -- to the underlying graphics system using a full three-dimensional coordinate system.  The device itself may then produce a 2d view of the content (e.g. on a standard desktop monitor); it might render the content to an immersive display (e.g. a dome) that gives implicit depth perception cues; or it might produce a genuine {\\em stereoscopic}\\/ visualization that presents slightly different views of the geometry to the viewer's left and right eyes.  Historically, the few 3D visualization tools that were developed did not achieve widespread use because processing speed limited their interactivity.  For example, the Karma {\\sc xray}\\footnote{Karma {\\sc xray} -- \\url{http://www.atnf.csiro.au/computing/software/karma}} package, developed more than ten years ago for volume rendering tasks, was not always able to reach usable rendering speeds on the desktop workstations available at that time, as it relied on the processing power of early Sun SPARC processors.  This is no longer the case: the astronomers' publishing medium now enables web-based colour graphics, animation and even interactive content; cheap 3D display devices are becoming available; and the computational demands of 3D graphics are now easily handled by off-the-shelf, hardware-accelerated graphics cards that ship with nearly every desktop and laptop computer sold today.  The explosion of availability of very fast graphics cards, fed by the massive computer game and entertainment market, now means that the most powerful processor in a desktop computer is often the graphics processor, not the CPU.  With specialized circuitry for placing graphics primitives into a 3D virtual space and rendering a view of that space directly to the screen buffer, displaying 3D environments on desktop computers is now practical and effective.  Evolution of the academic publishing medium, together with the wider availability of advanced computer graphics hardware, has led us and our collaborators to pursue the 3D display, analysis and publication paradigm.  Specifically: \\begin{itemize} \\item Beeson et al.\\ (2003) report on the development and implementation of a distributed volume rendering technique with stereoscopic support; \\item Beeson et al.\\ (2004) present a web-based tool for presenting multi-dimensional catalogues in 3D form via VRML technology; and \\item Fluke et al.\\ (2006) report on new, economical versions of traditional display techniques such as dome and tiled wall displays. \\end{itemize}  The Virtual Observatory paradigm (Quinn et al.\\ 2004) is also suggesting new ideas, and VOPlot3D\\footnote{VOPlot3D -- \\url{http://vo.iucaa.ernet.in/~voi/VOPlot3D_UserGuide_1_0.htm}} is an interesting development.  It is a Java applet that presents a 2d projection of a 3D point-based data set (extracted from a user-supplied multi-column catalogue) and enables rotation, panning and zooming of the view.  In Rixon et al.\\ (2004), the re-casting of the distributed volume renderer [dvr, Beeson et al.\\ (2003)] as a Grid-enabled application for remote server-based visualization was presented and discussed.  The Remote Visualization Service\\footnote{RVS -- designed and developed by the CSIRO Australia Telescope National Facility, see \\url{http://www.atnf.csiro.au/vo/rvs/}} is a nice example of the adaptation of sophisticated, legacy astronomy visualization software to the server-side visualization model, and while it does not (presently) offer 3D viewing capabilities, the plans for RVS always included the eventual incorporation of volume rendering (via dvr) as an advanced feature.  A high-profile example of a multi-dimensional data visualization system is provided by the UK National Cosmology Supercomputer (COSMOS -- \\url{http://www.damtp.cam.ac.uk/cosmos/Public/index.html}).  This system builds on the server-side visualization model to provide access to multiple visualization services from low-end, graphically ``primitive'' desktop computers.  The user (client) runs one of IRIS Explorer, Partiview or Vis5d+ on their workstation, and controls the rendering parameters locally, but the visualization itself is generated on the server, which provides substantial processor, memory, disk and graphics resources.  The rendered image is compressed and transferred back to the client machine for display.  This is a nice but central-resource-hungry system that provides non-stereoscopic 3D display services.  The visualization field in astronomy has become very active recently.  Yet on the whole, the uptake of the new programs for displaying and analysing data is exceedingly low, and it is fair to say that many in the community would view them as ``toys'' - they demonstrate a neat or useful idea, and might be useful occasionally.  They generally do not yet consider any of these systems as standard, in the league of \\pgplot, {\\sc miriad}, {\\sc IDL} and so on.  Because of this the new tools are not installed and potential users are not exposed to them.  Consequently  there is no {\\em de facto}\\/ standard library for 3D visualization and display.  As an aside, many of the above tools are actually very simple to use once installed, and on the whole are very well documented.  \\subsection{Community behaviour}  In view of the above, it was felt that consulting the community to assess their {\\em current}\\/ practices in astronomy data visualization and analysis might assist in the future development of visualization tools that astronomers will find useful and {\\em want}\\/ to adopt.  To this end, Fluke et al.\\ (2006) report on their Advanced Image Displays for Astronomy (AIDA) survey of Australian astronomers.  While based on a moderately small sample of forty-one astronomers, it is the best information we have on the behaviour and habits of the community in this domain.  {\\em Visualization and analysis environment.}  Excluding wavelength-specific software from Fluke et al.\\ (2006) Appendix~A Table~4, we find that the most common analysis tools are: Custom \\pgplot\\ tools (regularly used by 44\\% of respondents), Karma (39\\%), Mongo/SuperMongo (34\\%), IDL (29\\%) and other locally developed software (27\\%).  With the exception of Karma all of these are actually programming or scripting environments or libraries that provide functionality to use from program code. In fact, Karma {\\em does}\\/ provide a programming library and interface, but we know of no current Australian astronomers using the Karma library as opposed to its pre-packaged tools such as {\\sc kvis}, {\\sc xray}, etc.  This is a stark contrast to the recent developments described earlier, which were all pre-packaged, ``out-of-the-box'' applications for displaying or analysing data,  none of which provide anything like a well-documented {\\em programming}\\/ interface or scripting language. Yet the results of the AIDA survey suggest a preference amongst astronomers for this approach!  {\\em Visualization and analysis plot type.}  From Table~5 of Fluke et al.\\ (2006), we learn that the dominant graph/plot types used by astronomers are 2d graphs, histograms and plots (used by 95\\% of respondents).  Beyond that, 85\\% regularly display/produce 2d images, and 27\\% 3D images.  We expect that the respondents using ``3D images'' as a visualization method are including both 2d slices from 3D datasets, and also volume rendering.  There are no major surprises here.  {\\em Display device.} Participants in the AIDA survey were asked about their experience of display devices, ranging from the ubiquitous CRT/LCD desktop monitor and single-screen stereoscopic projection, to the sophisticated and expensive Virtual Room system.  While very few had used anything beyond the standard issue desktop monitor, $\\sim 50$\\% of respondents had seen a single-screen stereoscopic projection system in action but had not used such a device personally.  Approximately 70\\% had seen a digital dome projection system, as used in planeteriums, but again, had not used the device for their own work.  Participants were then asked what factors prevented them from using advanced image displays in their research.  Lack of knowledge of available displays was the most common response ($\\sim 70$\\% of respondents), followed by lack of appropriate software, lack of local facilities and cost (each nominated by $\\sim 30$\\% of respondents).  Comments made by participants highlighted ``the time required to develop appropriate tools to take advantage of these displays'' and ``[the] lack of knowledge of available software'' as key factors preventing higher uptake of advance display use in astronomy in Australia.  \\subsection{Responding to the community}  Based on the outcomes of the AIDA survey, we set about deducing what capabilities were most needed in a package that would positively encourage astronomers to begin using advanced, 3D plotting techniques. With recent experience of developing several advanced visualization tools, and the skills and knowledge in advanced displays built up at our institute, we have set about sharing this capability with the astronomy community in a form that we hope will experience significant uptake.  That is, we have created an advanced visualization library with the following key features and benefits: \\begin{enumerate} \\item{\\em provides three-dimensional plotting functions} --- dramatically increases capacity and options for display of multi-dimensional data. \\item{\\em provides a mechanism for navigating the 3D world} --- drastically improves comprehension of 3D geometry, and perception of depth on 2d devices. \\item{\\em is a programming library} --- can be called from existing and new C, FORTRAN and C++ codes. \\item{\\em has a PGPLOT-like interface} --- many of the functions will look familiar to PGPLOT users. \\item{\\em works on GNU/Linux and Apple/OSX} --- available for the most common astronomer workstations/laptops. \\item{\\em works on standard monitors} --- enables users to create and explore 3D geometry at their desktop, without requiring expensive hardware. \\item{\\em works on advanced displays} --- if a stereoscopic or dome display is available, it can be used with no change to the code. \\item{\\em comes with documented sample programs} --- allowing new users to learn by example. \\item{\\em is a binary distribution} --- no compilation is necessary to install the library or run the sample programs. \\item{\\em can save hardcopy output} --- for production of figures and movies for journals and websites.  The TGA format is used. \\item{\\em can save geometry files} --- allowing views to be saved to disk and displayed later without re-running code. \\end{enumerate}  VTK\\footnote{VTK -- \\url{http://public.kitware.com/VTK}} offers a number of the above features, but it is an extremely heavyweight solution, with an extended learning curve.  Additionally, VTK is written in C++, which is not the most popular language of the astronomy community.  IDL (versions 5 and above) now has support for 3D graphics, and {\\em in principle}\\/ with some clever coding can generate stereoscopic pairs for projection.  We have not seen IDL produce stereographics though, and would be interested to hear from astronomers who have done so.  Tipsy\\footnote{Tipsy -- \\url{http://www-hpcc.astro.washington.edu/tools/tipsy/tipsy.html}} is an application specific to N-body simulations: it can display particle positions and velocity vectors from an arbitrary viewpoint.  The user can follow selected particle(s) as they evolve, and calculate bulk properties of particle groups.  This specialized tool is well-used throughout the N-body community, but cannot be programmed to do new things.  It would be moderately simple to reproduce the capabilities of Tipsy using the \\splot\\ library, thus gaining support for dome and stereoscopic displays, and the capability for extensions to the application.  This paper briefly describes the architecture and implementation of \\splot\\ --- a library that meets the above specifications and is being made available (in binary form only) to the astronomical community. We then give some specific usage examples, consider some future extensions and uses, and close the paper with directions on obtaining the \\splot\\ distribution. ",
        "conclusions": ""
    },
    "0608/gr-qc0608034_arXiv.txt": {
        "abstract": "The retarded Green function of the electromagnetic field in spacetime of a straight thin cosmic string is found. It splits into a geodesic part (corresponding to the propagation along null rays) and to the field scattered on the string. With help of the Green function the electric and magnetic fields of simple sources are constructed. It is shown that these sources are influenced by the cosmic string through a self-interaction with their field. The distant field of static sources is studied and it is found that it has a different multipole structure than in Minkowski spacetime. On the other hand, the string suppresses the electric and magnetic field of distant sources---the field is expelled from regions near the string. ",
        "introduction": "\\label{sc:intro}  One of the simplest mass sources in general relativity are \\emph{cosmic strings}---linear objects with a given linear mass density and a linear tension. Cosmic strings arise as topological defects in various gauge theories (see, e.g., \\cite{Garfinkle:1985}), or as a macroscopic variant of the fundamental strings (e.g., \\cite{SarangiTye:2002}). Thin cosmic strings can be phenomenologically described by conical singularities of the spacetime metric---one-dimensional \\vague{objects} the angle around which is different than ${2\\pi}$. The static straight thin cosmic string is thus represented by a locally flat spacetime with a conical singularity along an axis. The deficit angle is proportional to the linear mass density which is equal to the linear tension \\cite{Vilenkin:1985}. A string with a deficit angle has a positive mass density and it is stretched; a string with an excess angle has negative mass density and is squeezed. For an overview of physics of cosmic strings see, e.g., \\cite{VilenkinShellard:book} and references therein; for more recent developments see~\\cite{DavisKibble:2005}.  Beside the empty spacetime with a single string, cosmic strings appear in a wide variety of solutions of Einstein equations. Any axially symmetric spacetime can be trivially modified to contain a string on the axis of symmetry. However, there exist also solutions where the string play a key physical role as, for example, C-metric. Here the string is a physical agent causing the motion of black holes---see, e.g., \\cite{KinnersleyWalker:1970,Bonnor:1982,Krtous:2005}. In a wide class of boost-rotation symmetric spacetimes \\cite{BicakSchmidt:1989} the strings accelerate even more general sources.  In the 80's the cosmic strings were considered as candidates for a mechanism of galaxy formation. This possibility was abandoned mainly because of inconsistencies with cosmic microwave background observations. Recently, however, the interest in cosmic strings reappeared in the context of the \\vague{brane-world} scenarios of the superstring theory. These suggest the existence of macroscopic fundamental strings behaving as \\vague{old-fashioned} cosmic strings (e.g., Refs.~\\cite{DavisKibble:2005,Kibble:2004} and references therein). There were also tentative hints of a detection of cosmic strings \\cite{SazhinEtal:2003,SchildEtal:2004} based on specific gravitational lensing effects, but the explanation in terms of cosmic strings was not confirmed by subsequent observations \\cite{AgolHoganPlotkin:2006}.  In the case of a single straight thin string the curvature is localized only on the world-sheet of the string. Such spacetime thus represents a very simple but non-trival example which can serve as a toy model for studying various phenomena due to curvature.  In this paper we investigate the behavior of the electromagnetic field in the background of a non-charged cosmic string. In Sec.~\\ref{sc:retgrfc} we find the retarded Green function and with its help we demonstrate the general behavior of the propagation of the electromagnetic field \\cite{DeWittBrehme:1960}: the field propagates (i) along null geodesics (on the light cone of the source), and (ii) it is scattered on the curvature (so-called \\vague{tail} term of the field). In our case the field propagates on the deformed light cones of the source and it is scattered on the cosmic string.  In Sec.~\\ref{sc:coulomb} we use the retarded Green function to derive the electric field of the static monopole and dipole sources. We obtain the field equivalent to that of Ref.~\\cite{Linet:1986}, which was derived by a different method. We also recover that the source is self-interacting: a monopole charge is repelled from the string by its own field, a dipole is forced to align \\vague{around} the string.  A magnetic field of the current flowing along the line parallel with the string is derived in Sec.~\\ref{sc:magfield}.  In Secs.~\\ref{sc:farfield} and \\ref{sc:homfield} we study the behavior of the electric field of strong static charges at large distances and near the string. We find an interesting effect: the asymptotic field has a different monopole structure from that in empty space, and the field of large charges which are localized far from the string is suppressed near the string. The field which would be nonvanishing and homogeneous in an empty space is expelled from the region near the string. It means that the string \\vague{shields} its neighborhood from the influence of distant sources. ",
        "conclusions": "We have discussed the electromagnetic field in the spacetime of a cosmic string. The retarded Green function \\eqref{grfcexpl} for the vector potential was found. It can be split into a \\vague{geodetical} part \\eqref{grfcgeod} corresponding to the empty space propagation of the field, and into a \\vague{scattered} part \\eqref{grfcscat} corresponding to scattering on the curvature located within the string.  Using this Green function electric and magnetic fields of various simple static sources were constructed. It was demonstrated that although the spacetime is locally flat (outside the string), the global deficit angle causes several interesting phenomena. We found a self-interaction of the point charge (repulsive away from the string), of the dipole (turning the dipole into the ${\\tens{e}_{\\ph}}$ direction \\vague{around} the string), or of the line current (attractive toward the string). In general, the self-interaction of charges and currents near the string implies that the string participates on electromagnetic processes even if it is not charged itself. It can have consequences when studying the motion of cosmic strings through, e.g., ionised plasma in an early Universe.  Further, we found that the field of a static source at large distances has a different multipole structure, namely, there is no dipole term corresponding to the dipole perpendicular to the string.  Near the string, the electric and magnetic fields orthogonal to the string are suppressed. For the same reason, there exists no homogeneous static field perpendicular to the string.  Most of the discussed effects can easily be understood in the special case ${\\iconpar=2}$ (respectively, for ${\\iconpar}$ being any integer). In this case the scattered part of the field vanishes and the resulting field can be obtained as the field in the half of Minkowski spacetime given by the real source and by a fictitious source obtained as a reflection of the real source with respect to the axis (these two contributions correspond to different terms in the sum over geodetics in \\eqref{scpot}). The self-interaction is thus given by the interaction of the real and the fictitious sources. A suppression of the dipole field at large distances arises because a combination of the dipole (orthogonal to the string) and of its reflected image gives a quadrupole configuration. And finally, the expulsion of the field away from the string arises as a cancelation of the fields of the source and of its image in middle between them.     After deriving the general form of the Green function we concentrated mainly on static phenomena. As a next step it would be interesting to study similar questions in dynamical contexts; for example, the scattering of plane waves on the string or the field of a strong oscillating dipole located far from the string.   \\vspace{-1.5ex}"
    },
    "0608/astro-ph0608559_arXiv.txt": {
        "abstract": "{The atmospheric parameters and iron abundance of the \\textit{Sloan Digital Sky Survey} (SDSS) spectrophotometric standard star BD~+17~4708 are critically examined using up-to-date Kurucz model atmospheres, LTE line formation calculations, and reliable atomic data. We find $\\teff=6141\\pm50$~K, $\\logg=3.87\\pm0.08$, and $\\feh=-1.74\\pm0.09$. The line-of-sight interstellar reddening, bolometric flux, limb-darkened angular diameter, stellar mass, and the abundances of Mg, Si, and Ca are also obtained: $E(B-V)=0.010\\pm0.003$, $\\fbol=(4.89\\pm0.10)\\times10^{-9}$~erg~cm$^{-2}$~s$^{-1}$, $\\theta=0.1016\\pm0.0023$~mas, $M=0.91\\pm0.06M_\\odot$, $\\mathrm{[Mg/Fe]}=0.40\\pm0.10$, $\\mathrm{[Si/Fe]}=0.35\\pm0.11$, $\\mathrm{[Ca/Fe]}=0.36\\pm0.11$. This star is a unique example of a moderately metal-poor star for which the effective temperature ($\\teff$) can be accurately constrained from the \\textit{observed} spectral energy distribution (corrected for reddening). Such analysis leads to a value that is higher than most spectroscopic results previously reported in the literature ($\\sim$5950~K). Interstellar reddening was estimated using various prescriptions, including an analysis of interstellar lines. The surface gravity of the star was inferred from the fitting of the wings of the \\ion{Mg}{i}\\,\\textit{b} lines. We used transition probabilities measured in the laboratory and reliable damping constants for unblended Fe lines to derive the iron abundance using both \\ion{Fe}{i} and \\ion{Fe}{ii} lines. We find that the ionization balance of Fe lines is satisfied only if a low $\\teff$ ($\\sim5950$~K) is adopted. The mean iron abundance we obtain from the \\ion{Fe}{ii} lines corresponds to $A_\\mathrm{Fe}=5.77\\pm0.09$ ($\\feh=-1.74$ for our derived $A_{\\mathrm{Fe},\\odot}=7.51$) while that from the \\ion{Fe}{i} lines is $A_\\mathrm{Fe}=5.92\\pm0.11$, and therefore with our preferred $\\teff$ ($6141$~K), the discrepancy between the mean iron abundance from \\ion{Fe}{i} and \\ion{Fe}{ii} lines cannot be explained by overionization by UV photons as the main non-LTE effect. Interestingly, the \\ion{Fe}{i} excitation balance is satisfied with a $\\teff$ only slightly warmer than our preferred solution and not with the lower value of 5950~K. We also comment on non-LTE effects and the importance of inelastic collisions with neutral H atoms in the determination of oxygen abundances in metal-poor stars from the 7774~\\AA\\ \\ion{O}{i} triplet. ",
        "introduction": "The derivation of stellar chemical compositions relies on the accurate determination of the atmospheric parameters $\\teff$ (effective temperature) and $\\logg$ (surface gravity). These quantities may be inferred either from the stellar spectrum or by semi-empirical methods that are normally based on photometric and astrometric measurements. Often, the photometric and astrometric parameters are used as first approximations and the final solution is found iteratively with the help of the spectrum. This tuning of parameters is, however, model-dependent, and may lead to erroneous conclusions if the models are inadequate.  Most abundance analyses of FGK stars are made using homogeneous plane-parallel model atmospheres and LTE (local thermodynamic equilibrium) line formation. However, recent abundance analyses using hydrodynamical model atmospheres and non-LTE line formation have demonstrated that the effects of surface inhomogeneities and departures from LTE on abundance analyses are not negligible in the Sun and solar-type stars of different metallicities (e.g., Asplund \\& Garc\\'{\\i}a~Perez 2001, Korn et~al. 2003, Allende~Prieto et~al. 2004b, Asplund 2005), being probably dramatic for very metal-poor stars (Shchukina et~al. 2005).  The F8-type star BD~+17~4708 has been chosen as a spectrophotometric standard, either primary or secondary, for various systems (e.g., Oke \\& Gunn 1983, Rufener \\& Nicolet 1988, J\\o rgensen 1994, Zhou et~al. 2001). In particular, this star is the primary standard of the \\textit{Sloan Digital Sky Survey} (SDSS) photometric system (Fukugita et~al. 1996; Gunn et~al. 1998, 2006; Smith et~al. 2002). BD~+17~4708 is one of the few stars, and the only subdwarf, with very accurate absolute fluxes (Bohlin \\& Gilliland 2004b), which allows us to test, in an independent way, the models and different atmospheric parameters that have been derived for it. In particular, it is a unique case in which the effective temperature may be well constrained. Also, it can be quite useful to know, with high accuracy, what the fundamental parameters and overall chemical composition of this star are, given that its model atmosphere and predicted fluxes may be used to complement the observed spectral energy distribution in the transformation of observed magnitudes into physical fluxes in the SDSS.  BD~+17~4708 has been studied by several groups (Table~\\ref{t:literature}), who have derived effective temperatures between 5800~K and 6200~K, $\\feh$ values between $-2.0$ and $-1.4$ and a surface gravity of $\\logg\\simeq4.0$ (Fig.~\\ref{f:literature}).\\footnote{Throughout this paper, we use the standard definitions $\\mathrm[X/Y]=\\log(N_X/N_Y)-\\log(N_X/N_Y)_\\odot$ and $A_\\mathrm{X}=\\log(N_X/N_H)+12$, where $N_X$ is the number density of the element $X$. The surface gravity $g$ in $\\logg$ is given in cgs units.} In this paper, we present a detailed determination of the atmospheric parameters and the iron abundance of BD~+17~4708. A critical comparison with values previously reported in the literature is provided in the Appendix. Due to their relevance for studies of stellar interiors, we also derive the abundances of Mg, Si, and Ca, and discuss the uncertainties in the determination of the oxygen abundance from the 7774~\\AA\\ triplet, in particular the importance of inelastic collisions with neutral H atoms in the non-LTE computations involved. Our analysis is based on the observed spectral energy distribution and a high resolution, high signal-to-noise (S/N) spectrum of the star.  \\begin{table} \\centering \\begin{tabular}{lccc} \\hline\\hline Reference & $\\teff$ (K) & $\\logg$ & $\\feh$ \\\\ \\hline Peterson (1981)                & 5800 & 4.00 & $-1.95$ \\\\ Rebolo et~al. (1988)           & 5890 & 4.00 & $-1.70$ \\\\ Magain (1989)                  & 5960 & 3.40 & $-1.93$ \\\\ Axer et~al. (1994)             & 6100 & 4.40 & $-1.42$ \\\\ Spite et~al. (1994)            & 5950 & 3.30 & $-1.50$ \\\\ Th\\'evenin \\& Idiart (1999)    & 5929 & 4.02 & $-1.54$ \\\\ Boesgaard et~al. (1999, K93)   & 6091 & 3.81 & $-1.73$ \\\\ Boesgaard et~al. (1999, C83)   & 5956 & 3.65 & $-1.81$ \\\\ Fulbright (2000)               & 6025 & 4.00 & $-1.63$ \\\\ Mishenina et~al. (2000)        & 6000 & 4.00 & $-1.65$ \\\\ Ryan et~al. (2001)             & 5983 & ---- & $-1.86$ \\\\ Simmerer et~al. (2004)         & 5941 & 3.98 & $-1.60$ \\\\ Nissen et~al. (2004)           & 5943 & 3.97 & $-1.64$ \\\\ Mel\\'endez \\& Ram\\'{\\i}rez (2004) & 6154 & 3.93 & $-1.64$ \\\\ Asplund et~al. (2006)          & 6183 & 4.11 & $-1.51$ \\\\ \\hline \\end{tabular} \\caption{Literature data for BD~+17~4708. Note that there are two sets of parameters given by Boesgaard et~al. (1999): those given in the King (1993, K93) scale, and those given in the Carney (1983a,b; C83) scale. In this Table, only values derived for the first time (although mostly taken from the more recent papers) by each author(s) are shown to avoid duplicity.} \\label{t:literature} \\end{table}  \\begin{figure} \\centering \\includegraphics[width=7.0cm]{5647fig1.eps} \\caption{Distribution of the $\\teff$, $\\feh$, and $\\logg$ values found in the literature for the star BD~+17~4708, as given in Table~\\ref{t:literature}.} \\label{f:literature} \\end{figure} ",
        "conclusions": "The high accuracy with which the spectral energy distribution of the SDSS standard BD~+17~4708 has been measured has allowed us to provide a reliable estimate of its effective temperature. We have then used spectral line analysis to infer consistent parameters for this star.  Once the degeneracy between $\\teff$ and $E(B-V)$ in the model fits to the observed flux distribution is broken by independent estimates of $E(B-V)$, which included a detailed modeling of interstellar absorption features in the observed spectrum, we obtain the following parameters, given here along with reasonable estimates of the ($1\\sigma$) error bars: $\\teff=6141\\pm50$~K, $\\feh=-1.74\\pm0.09$, $\\logg=3.87\\pm0.08$, $E(B-V)=0.010\\pm0.003$. The spectral energy distribution also allowed us to obtain reliable values for the bolometric flux, $f_\\mathrm{bol}=4.89\\pm0.10\\times10^{-9}$~erg~cm$^{-2}$~s$^{-1}$, and angular diameter, $\\theta=0.1016\\pm0.0023$~mas, of the star. We thus provide accurate (in an absolute sense) parameters for the spectrophotometric standard of the \\textit{Sloan Digital Sky Survey}.  Compared to previous spectroscopic studies, our $\\teff$ is higher by about 190~K, which has a severe impact on classical abundance analyses of moderately metal-poor stars. For example, this increase ruins the good agreement between the iron abundance derived from \\ion{Fe}{i} and \\ion{Fe}{ii} lines found when a lower $\\teff$ is used. Despite this, all other features on the spectrum (e.g., the Balmer lines or the strong Mg lines used to constrain the $\\logg$ value) seem to be more consistent with the models when a high $\\teff$ is adopted. In particular, the excitation balance of \\ion{Fe}{i} lines is satisfied with the high $\\teff$ but not with the lower value.  We also determine the mean abundance of $\\alpha$-elements ($\\mathrm{[\\alpha/Fe]}=0.37\\pm0.06$). The non-LTE modeling of the permitted oxygen triplet lines should include the effect of collisions with neutral H to reasonably reproduce the observations but a better physical treatment of H collisions is needed."
    },
    "0608/astro-ph0608073_arXiv.txt": {
        "abstract": "{A mosaic of six, deep, CCD images in the light of the \\oiii\\ nebular emission line has been obtained with the 1.3--m Skinakas (Crete) telescope of the filamentary nebulosity surrounding P~Cygni. The \\oiii\\ line discriminates against confusing galactic H{\\sc ii} regions along the same sight--lines and the new mosaic did not include the 4.8 mag. central star; a source of artifacts in the previous lower angular resolution observations. New giant `lobes' and `shells' are found to be clustered around P~Cygni which must be the relics of historic eruptions between 2400 and up to $\\approx$~10$^{5}$~yr ago. ",
        "introduction": "The circumstellar environment of the proto-typical Luminous Blue Variable star (LBV - Conti 1984; Humphreys 1989; Davidson, Moffat \\& Lamers 1989) P~Cygni has been revealed at optical wavelengths in the work presented in a sequence of seven papers (Johnson et al. 1992; Barlow et al 1994; Meaburn et al. 1996; O'Connor, Meaburn \\& Bryce 1998; Meaburn, L\\'{o}pez \\& O'Connor 1999; Meaburn et al. 2000; Meaburn et al. 2004).  In these papers the discovery of a 22\\arcsec\\ diameter inner shell (IS) was found to be surrounded by a 1.6\\arcmin\\ diameter, outer shell (OS). More controversially it was suggested that a large region of filamentary nebulosity which surrounds the star could have been ejected in the early stages of its evolution and not simply be an unrelated supernova remnant along the same sight--line.  The latest kinematical and morphological evidence in support of this ejection proposition was presented in Meaburn et al. (2004) as a consequence of a deep, very wide--field, image in the light of the \\oiii\\ nebular emission line, though still some reservations remained about the association with P~Cygni. The `giant lobe' (GL) projecting to the east of P~Cygni (O'Connor, Meaburn \\& Bryce 1998) was shown to have southern and eastern counterparts (here these will still be referred to as GLs). The wide--field CCD \\oiii\\ image by Meaburn et al. (2004) was restricted in its sensitivity by the inclusion of the 4.8--mag star P~Cygni in the same field; artifacts due to gross signal saturation and diffraction became unacceptable in places and areas of the final image had to be blanked out. Incidentally the \\oiii\\ line was chosen for this imagery as this is dominant in the emission from the GL nebulosity; consequently considerable discrimination was achieved against the predominantly \\ha\\ and \\nii\\ emitting, diffuse, galactic H{\\sc ii} regions in the broader field.  This saturation/artifact problem has since been overcome by obtaining a mosaic of deeper \\oiii\\ CCD images at $\\approx$ 5 times higher angular resolution than previously covered fields adjacent to P~Cygni but where none includes the star itself. New GL features that are revealed help to confirm -- and see Sect.3 -- the association of all of the extensive, filamentary GL nebulosity with P~Cygni. ",
        "conclusions": "The wider discussion involving both the morphology and kinematics of the GL nebulosity around both P~Cygni and related LBV stars in Meaburn et al. (2004) will not be repeated here. Most of the same conclusions stand. The purpose of the present paper is to report new observations, that help to confirm the direct relationship between the surrounding, extensive, high-excitation filamentary GL nebulosity and P~Cygni. This association was always left conservatively in doubt after the previous more limited observations.  Here (Figs. 1 \\& 2), P~Cygni appears to be completely surrounded by overlapping, elongated, shell--like filamentary structures over NS and EW diameters of 16\\arcmin\\ (8.4 pc) and 20\\arcmin\\ (10 pc), respectively, where the distance to this star is 1.8 kpc (van Schewick 1968; Lamers et al. 1983). The \\nii\\ and \\ha\\ line profiles (Meaburn et al. 1999, 2000 \\& 2004) over these GLs have many of the velocity characteristics of those of the optical filaments of evolved supernova remnants and, in the limited high spectral resolution observations so far obtained, exhibit predominantly receding radial velocities with respect to \\vsys\\ though coming to this value close to the star. However, it is this distribution of lobes and shells, apparently centered on P~Cygni that is compelling evidence for a direct association. The positive radial velocities found in the previous limited kinematical observations are still explicable (Meaburn et al. 2004) as a 'breeze' as P~Cygni moves towards the observer at high velocity with respect to it local ambient medium.  The images in Figs. 1 \\& 2 reveal the multiple lobes and shells, each $\\approx$~5pc across, that must be a consequence of continuous eruptions of P~Cygni over a period of 2400--10$^{5}$ yr ago. The lower limit is the well--determined (Meaburn et al. 2000) dynamical age of the OS (shown in the inset in Fig. 1) and the upper limit is that estimated (Meaburn et al. 2004) for the extremities of the bright eastern GL also prominant in Fig. 1. All of these features then precede the observed 1600 AD `great outburst' of P~Cygni (de Groot 1969) but are within the duration of the LBV phase expected for a 50\\msun\\ star (Humphreys \\& Davidson 1994 and note that Turner et al. 2001 suggested P~Cygni had a progenitor mass as low as 25\\msun ). It is now thought most unlikely that the alternative suggestion, that unrelated SNR optical filaments along the same sight--line are being observed. It is the symmetry of the \\oiii\\ shells and lobes, and of the general \\oiii\\ emission, around P~Cygni combined with the previous kinematical associations (Meaburn et al 2004) that eliminates this possibility.  Spatially--resolved profiles of the \\oiii\\ line are now required over the whole field to investigate the dynamics of these ejected phenomena more thoroughly. So far only receding motions have been detected. Also low--dispersion spectra are required to measure the velocities of the shocks that are undoubtably exciting these extensive \\oiii\\ emitting filaments. Incidentally, the rather unstructured and inhomogeneous nature of the nebulosity around P~Cygni (compared with planetary nebulae, for example) is reminiscent of other high--mass evolved stars; WR 124 is surrounded by a clumpy ring of ejecta that may have originated in a preceeding LBV phase (Grosdidier et al. 1998; van der Sluys \\& Lamers 2003)."
    },
    "0608/astro-ph0608590_arXiv.txt": {
        "abstract": "Galaxies are not ``closed-boxes''.  Both gas accretion (infall) and winds (outflow) change a galaxy's metallicity and gas fraction, lowering the galaxy's effective yield ($\\equiv\\!Z/\\ln{(f_{gas}^{-1})}$).  Low effective yields are seen in galaxies with rotation speeds less than $\\sim\\!120\\kms$, and have been widely interpreted as the onset of supernova-driven winds below a characteristic galaxy mass.  However, accretion of metal-poor gas is also a viable explanation for the low effective yields observed in low mass galaxies.  Analytic calculations presented in this paper prove: (1) that neither infall nor unenriched outflows can produce effective yields that are as low as observed; (2) that metal-enriched outflows are the only mechanism that can significantly reduce the effective yield, but only in gas-rich systems; (3) that it is nearly impossible to reduce the effective yield of a gas-poor system, no matter how much gas is lost or accreted; and (4) that any subsequent star formation rapidly drives the effective yield back to the closed-box value.  Given that massive galaxy disks are systematically more gas-poor than low mass dwarf irregulars, these calculations imply that: (1) massive galaxies may have experienced substantial infall and/or outflow in spite of their high effective yields; (2) only gas-rich systems can have low effective yields; and (3) only systems with low star formation rates can maintain low effective yields in between episodes of mass loss.  The drop in effective yield seen in low mass galaxies is therefore not necessarily due to the onset of supernova-driven winds.  Instead, it is likely due to the fact that galaxies with $V_c\\lesssim\\!120\\kms$ have surface densities lying entirely below the Kennicutt star formation threshold, as previously noted by Verde et al. and Dalcanton et al.  The resulting low star formation efficiency keeps these galaxies sufficiently gas-rich that their effective yields can be reduced by outflows, and keeps their effective yields low between episodes of gas loss.  Additional calculations confirm that there is no galaxy mass at which the metal-loss becomes dramatically stronger.  The fraction of baryonic mass lost through winds varies only weakly with galaxy mass, shows no sharp upturn at any mass scale, and does not require that more than 15\\% of baryons have been lost by galaxies of any mass. Supernova feedback is therefore unlikely to be an effective mechanism for removing large amounts of gas from low mass disk galaxies.  In addition, the dependence between metal-loss and galaxy mass is sufficiently weak that massive galaxies dominate metal enrichment of the IGM.  The calculations in this paper are based on impulsive accretion, outflow, and star formation events that provide limiting cases for more arbitrary chemical evolution histories, as proven in an Appendix. ",
        "introduction": "\\label{introsec}  Since the seminal work of \\citet{larson74}, supernova-driven winds have been a favored mechanism for explaining the properties of low mass galaxies.  Modern theories of galaxy formation frequently invoke gas outflows (or ``feedback'') to explain the lack of gas in dwarf spheroidals \\citep[e.g.,][]{sandage65,dekel86,lanfranchi04}, the paucity of low mass galaxies compared to CDM simulations \\citep[e.g.,][]{white91}, the hot gaseous halos around dwarf starbursts \\citep[e.g.,][]{marlowe95, martin98, martin02, ott05}, and the metal enrichment of the intergalactic medium \\citep[e.g.,][]{silk87,madau01,mori02}.  Outflows may also be partially responsible for low metallicities in disk galaxies, particularly in low mass dwarf irregulars.  For many years there has been evidence for a mass-metallicity relationship among galaxies, with low mass galaxies having systematically lower metallicities \\citep[see Figure~7 of the review by ][ for an early plot of this relation\\footnote{While outflow was initially considered as a possible explanation for the mass-metallicity relationship in disks, most initial theoretical work on outflows focused primarily on elliptical systems \\citep[e.g.][]{hartwick80}.}]{pagel81}.  While some of the correlation between mass and metallicity is certainly due to less past enrichment in gas-rich low mass galaxies, some may also be due to a direct loss of metals through winds.  Recently, after bringing abundance measurements to a common metallicity scale and fiducial radius, \\citet{garnett02} found a tight relationship between metallicity and galaxy mass.  By combining measurements of the metallicity with estimates of the gas richness, he showed that galaxies systematically depart from the closed-box model of chemical evolution below a characteristic mass scale near $V_c\\sim125\\kms$.  These results have since been qualitatively confirmed by \\citet{pilyugin04}, using an alternative abundance calibration, and by \\citet{tremonti04}, using a much larger sample of metallicities from the Sloan Digital Sky Survey (SDSS), although with less accurate gas mass fractions.  These results have widely been interpreted as evidence for the onset of strong galactic winds at a specific galaxy mass scale, below which gas and metals can easily escape the shallow potential well.  Although most theoretical attention has focused on outflow, infall is an equally appealing mechanism for producing low metallicities. Outflow reduces the metallicity by preferentially driving metals out of a system, while infall reduces the metallicity by diluting a galaxy's interstellar medium (ISM) with fresh, low metallicity gas. As shown by \\citet{koppen99}, the metallicity drops if rate of infall is higher than the rate of star formation, thus allowing the accreted metal-poor gas to dilute the ISM faster than it can be enriched by evolving stars.  Given the very low star formation rates and star formation efficiencies seen in the majority of low mass disk galaxies \\citep[e.g.][]{hunter04}, infall is therefore a viable mechanism for producing low metallicities in low mass galaxies.  The conclusion that low mass galaxies have not evolved as ``closed-boxes'' is based on measurements of their effective yields. The effective yield measures how a galaxy's metallicity deviates from what would be expected for a galaxy that had the same gas mass fraction, but that had been evolving as a closed-box, i.e.\\ with no inflow or outflow of gas.  A system that evolves as a closed-box obeys a simple analytic relationship between the metallicity of the gas and the gas mass fraction (see reviews by Pagel 1997 and Tinsley 1980).  As gas is converted into stars, the gas mass fraction $f_{gas}$ ($\\equiv M_{gas}/(M_{gas}+M_{stars})$) decreases and the metallicity $Z_{gas}$ of the gas increases according to \\begin{equation}    \\label{closedboxeqn} Z_{gas} = y_{true} \\, \\ln{(1/f_{gas})} \\end{equation} \\citep{searle72} where $y_{true}$ is the true nucleosynthetic yield, defined as the mass in primary elements freshly produced by massive stars, in units of the stellar mass that is locked-up in long-lived stars and stellar remnants.  For closed-box evolution, the metallicity increases without limit, and equals the nucleosynthetic yield when $f_{gas}=\\exp{(-1)}$.  Equation~\\ref{closedboxeqn} assumes that the metals produced by a generation of stars are instantly returned to the ISM and are well-mixed with existing gas.  This ``instantaneous recycling'' approximation is likely to be valid for measurements of galaxies' current gas-phase metallicities, which typically use oxygen abundances measured in HII regions.  Since the production of oxygen is dominated by winds and SN from massive stars ($>\\!8\\msun$; see Figures~5~\\&~7 from Maeder 1992), prompt return of oxygen to the local ISM is a reasonable assumption.  If a galaxy evolves as a closed-box, the ratio of $Z_{gas}/\\ln{f_{gas}^{-1}}$ should be a constant equal to the nucleosynthetic yield.  However, this ratio will be lower if metals have been lost from the system through winds, or if the current gas has been diluted with fresh infall of metal-poor gas.  One may therefore define the above ratio as the ``effective yield'' $y_{eff}$ \\begin{equation}    \\label{peffeqn} y_{eff} \\equiv \\frac{Z_{gas}}{\\ln{(1/f_{gas})}}, \\end{equation} which will be constant ($y_{eff}=y_{true}$) for any galaxy that has evolved as a closed-box, assuming that the nucleosynthetic yield is invariant.  In contrast, if any gas has either entered or left the galaxy, the measured effective yield will drop\\footnote{Note that there is no plausible process that can increase the effective yield, as proved by \\citet{edmunds90} and shown graphically in an elegant paper by \\citet{koppen99}.  The only exception is accretion of metal-rich gas -- a highly unlikely possibility that will not be considered further in this paper.  Thus, the largest observed value of $y_{eff}$ is a lower limit to the true nucleosynthetic yield.} below the closed-box value of $y_{eff}=y_{true}$, due to changes in the metallicity $Z_{gas}$ and/or the gas mass fraction $f_{gas}$.  The effective yield is therefore an observationally determined quantity that can be used to diagnose departures from closed-box evolution.  Unfortunately, low values of $y_{eff}$ alone do not immediately reveal {\\emph{why}} a system has departed from closed-box evolution.  Either the addition of metal-poor gas, or the removal metal-rich gas could suppress the measured effective yield.  It is therefore premature to assume that supernova-driven outflows alone can explain the low effective yields seen for galaxies with $V_c\\lesssim 100-120\\kms$.  This paper explores ways to distinguish between infall and outflow using the effective yield and the gas mass fraction. \\S\\ref{effectiveyieldobssec} summarizes the current observations of effective yields.  Calculations in \\S\\ref{gasflowsec} show how infall (\\S\\ref{infallsec}), outflow (\\S\\ref{outflowsec}), and the subsequent return to closed-box evolution (\\S\\ref{evolutionsec}) change the effective yield and the gas mass fraction. A comparison with observational data in \\S\\ref{pilyuginsec} leaves metal-rich outflows as the only viable mechanism for producing the low effective yields observed in gas-rich galaxies. \\S\\ref{thresholdsec} stresses the importance of high gas-richness in allowing enriched outflows to suppress the effective yield and suggests that the observed mass-dependent threshold in the effective yield is more closely linked to low star formation efficiencies than to the depth of the potential well.  Simple models in \\S\\ref{modelsec} derive the dependence of mass and metal-loss as a function of galaxy rotation speed.  Following the conclusions in \\S\\ref{conclusionsec}, Appendix~\\ref{proofsec} presents a proof showing how the effective yields calculated for simple ``impulsive'' gas flows are limiting cases of those that result from arbitrary chemical evolution histories.  Appendix~\\ref{garnettsec} re-examines the conclusions of the paper using the smaller sample of abundances calculated by \\citet{garnett02}. ",
        "conclusions": "  \\begin{enumerate}  \\item There is a minimum effective yield that can be produced by gas accretion.  The minimum is significantly higher than the effective yields observed in low mass galaxies.  Thus low effective yields cannot be due to gas infall.  \\item Outflows that drive out gas with the mean metallicity of the ISM can only produce low effective yields if they remove nearly the entire ISM.  Because the same galaxies that show low effective yields also show high gas fractions, outflows of unenriched gas cannot produce low effective yields.  \\item Outflows that consist primarily of escaped SN ejecta are extremely efficient at reducing the effective yield.  Metal-enriched outflow is therefore the only viable mechanism for producing galaxies with low effective yields.  \\item Metal-enriched outflows are ineffective at reducing the effective yields of gas-poor systems.  Thus, only galaxies that have maintained high gas mass fractions will show depressed effective yields in response to outflow.  Only galaxies with inefficient star formation are capable of remaining gas-rich to the present day, and thus they are the only systems which can show significantly depressed effective yields.  \\item Any star formation that takes place after gas accretion or gas loss will increase the effective yield back to the true nucleosynthetic yield expected for closed-box evolution.  Thus, only galaxies with inefficient star formation can maintain low effective yields long after infall or outflow has finished.  \\item Points 4~\\&~5 suggest that in addition to the ability to drive winds, galaxies with low effective yields must also have high gas mass fractions and low star formation rates.  Thus, the drop in effective yield observed at $V_c\\!\\sim\\!120\\kms$ ($M_{baryon}\\!\\sim\\! 10^{10}\\msun$) cannot be interpreted solely as a mass threshold for escape of SN-driven winds.  More massive galaxies may have also experienced significant outflows, but will show no reduction in their effective yields due to their low gas mass fractions and efficient star formation.  \\item Points 1~\\&~2 should not be taken as evidence that infall and unenriched outflows have not occurred.  They only indicate that these two processes produce little noticeable change in a galaxy's effective yield.  \\item At high redshifts, a larger fraction of galaxies are likely to be gas-rich.  Therefore a larger fraction of high redshift galaxies will show reduced effective yields, even if the rates and sizes of typical outflows are identical to those at the present day.  \\item The current data on the effective yield and gas mass fraction as a function of galaxy rotation speed suggest that the fraction of metals lost due to outflows increases steadily toward lower mass galaxies, reaching 50\\% at $\\lesssim30\\kms$.  However, the fraction of baryonic mass lost is quite modest ($\\lesssim 15$\\%) at all galaxy masses.  Neither the fraction of metals lost or the fraction of baryons lost shows a significant feature at $V_c\\!\\sim\\!120\\kms$.  \\item If either the initial mass function or the nucleosynthetic yield depends strongly on metallicity, then the amount of gas and metal-loss needed to explain the effective yield data would be different than calculated above.  \\end{enumerate}"
    },
    "0608/astro-ph0608629_arXiv.txt": {
        "abstract": "We present phase-referenced VLBI results on the radio continuum and the OH 18~cm megamaser emission from the Ultra-Luminous Infrared Galaxy, IRAS~17208--0014. The observations were carried out at 1599~MHz using the Very Long Baseline Array, the phased VLA, and the Green Bank Telescope. The highest resolution radio continuum results show several compact sources with brightness temperatures on the order of $10^{6}$~K. These sources are more likely to be clustered supernova remnants and/or luminous radio supernovae. However, the agreement between the number of observed and expected compact sources above the 5~$\\sigma$ level supports the possibility that each one of the compact sources could be dominated by a recently detonated luminous radio supernova. The continuum results suggest that there is no radio-loud AGN in the nuclear region of this galaxy. The OH 18~cm megamaser emission in IRAS~17208--0014 is detected at various angular resolutions. It has an extent of $170 \\times 110$~pc, and is mostly localized in two regions separated by 61~pc. The structure and dynamics of the maser emission seem to be consistent with a clumpy, rotating, ring-like geometry, with the two dominant maser regions marking the tangential points of the proposed rotating-ring distribution. Assuming Keplerian motion for the rotating maser ring, the enclosed dynamical mass and the mass density within a radius of 30.5~pc, are about {$3 \\times 10^7~({\\rm sin}^{-2}i)~M{_\\odot}$}, and $281~({\\rm sin}^{-2}~i)~M{_\\odot}~{\\rm pc}^{-3}$, respectively. ",
        "introduction": "At bolometric luminosities greater than $10^{11}~L{_\\odot}$, infrared galaxies become the most numerous objects in the local universe ($z \\leq 0.3$) \\citep {SM96}. The trigger for the intense infrared emission appears to be the strong interaction or merger of molecular gas-rich spirals. Galaxies with the highest infrared luminosities ($L{_{\\rm IR}}(8-1000~{\\mu}m) \\geq 10^{12} L{_\\odot}$), known as Ultra-Luminous Infrared Galaxies (ULIRGs), appear to be advanced merger systems powered by a nuclear starburst and/or an AGN. These may represent an important stage in the formation of quasi-stellar objects \\citep{SAN88}.  The extreme merging conditions in ULIRGs could, in some cases, induce nuclear OH 18~cm masing several orders of magnitude stronger than for typical Galactic masers, i.e., megamasers ($10^6$ times), or even gigamasers ($10^9$ times) \\citep{B89,B92}. Recent surveys by \\citet{DG02} report an OH megamaser (OHM) detection rate of up to 33\\% for ULIRGs at $z>0.1$.  Sensitive Very Long Baseline Interferometry (VLBI) observations provide the most direct means to study OHM and radio continuum emission in the innermost, dust-obscured, regions of merging systems. While radio continuum studies play a key role to directly image and determine the nature of the nuclear power sources (AGN and/or starburst) in these galaxies, the OHM emission studies can enhance our understanding of the kinematics on parsec scales in their nuclear regions, and be used as probes of AGNs or intense nuclear starbursts.  In the present paper, we report VLBI observations of the 18 cm radio continuum and OHM emission from the galaxy IRAS~17208--0014 at $z=0.04281$. This object is one of the few strong OHM emitters at $z < 0.1$ that can be imaged at milliarcsecond (mas) resolution, and contribute to a better understanding of the nature, extent, and dynamics of extragalactic OHM emission.  This galaxy, which is classified as a ULIRG, has an infrared luminosity of $L_{\\rm IR}= 2.5 \\times 10^{12} L{_\\odot}$, as defined in \\citet {GOL95}. We adopt a distance of 171~Mpc, assuming ${H_\\circ=75}$~km~s$^{-1}$~Mpc$^{-1}$. At this distance, 1~mas corresponds to 0.76~pc.  Optical images of IRAS~17208--0014 at 6550 {\\AA} show two tidal tails from a merger \\citep {MM90,MUR96}. Near-IR images show a very disturbed morphology and an extended, but single, nucleus, suggesting a very advanced merger \\citep {ZL93,MUR96}. Infrared observations suggest that this galaxy represents the extreme of starburst-dominated sources of this type \\citep {SOI00}, and an observational proof that a collision of galaxies can lead to a mass distribution similar to elliptical galaxies \\citep {ZL93}.  The radio continuum emission of this galaxy at milliarcsecond (mas) resolution was first detected by the VLBI observations of \\citet {MOM03} at 21~cm. These observations revealed that the radio emission is dominated by an extreme nuclear starburst, with no indication of a compact high-brightness temperature source (i.e., a radio-loud AGN).  IRAS~17208--0014 exhibits OHM activity at 1665, 1667, and 1720~MHz \\citep{MAR89}. The strongest emission is in the 1667~MHz line, with a luminosity of $L_{\\rm OH}=1054~L_{\\odot}$ \\citep{MAR89}. Very Large Array (VLA) and single-dish Nan\\c{c}ay observations showed the existence of two separate velocity structures with similar widths for both the 1667 and 1665~MHz OH main lines \\citep {MAR89}. However, the 18~cm VLBI observations of \\citet{DLLS99} did not detect the 1665~MHz line, nor the continuum emission, and recovered only 20\\% of the 1667~MHz line. ",
        "conclusions": "We have presented the results of sensitive, phase-referenced VLBI observations of the 18 cm continuum and OH megamaser emission in the central 0.4~kpc of the ULIRG IRAS~17208--0014.  The continuum emission from this galaxy was imaged at a higher resolution than has been previously reported. Several compact sources with flux densities greater than 5$\\sigma$ were detected. These sources have brightness temperatures on the order of $10^{6}$~K, and are more likely to be clustered RSNe and SNRs. However, we cannot rule out the possibility that each of the compact sources is mainly powered by an individual bright RSN nested in a region that contains faded SNRs. This is supported by the agreement in the number of the observed and expected continuum sources above the 5~$\\sigma$ level. The continuum results do not reveal a radio-loud AGN in the nuclear region of this galaxy.  The OHM results show the detection of both the 1665 and 1667~MHz mainlines in IRAS~17208--0014. The OHM emission in this ULIRG has an extent of $170 \\times 110$~pc, and is mostly localized in two regions separated by 80~mas (61~pc). The overall structure and dynamics of the maser emission seem to be consistent with a clumpy, rotating, ring-like geometry, with the two dominant maser regions marking the tangential points of the proposed ring structure. However, our OH results reveal another component extended to the south. Considering the ongoing intense starburst in this advanced merger system, the southern extension could to be associated with a gaseous inflow or outflow. Assuming Keplerian motion for the observed OH ring structure, the enclosed dynamical mass and the mass density within a radius of 30.5~pc, are about {$3 \\times 10^7~({\\rm sin}^{-2}i)~M{_\\odot}$}, and $281~({\\rm sin}^{-2}~i)~M{_\\odot}~{\\rm pc}^{-3}$, respectively.  Our results show that the starburst activity, as seen in the radio continuum images, has a larger extent than the OHM emission. Moreover, the low-resolution images show that at larger scales, the OHM emission peaks do not coincide with the peak of the continuum emission. This could be the results of free-free absorption and/or possibly an indication that the merger is incomplete. At the full angular resolution, the brightest continuum and OHM spots are not correlated. The derived lower limits on the maser amplification are up to 130.  The observed line-to-continuum ratios, 1667 to 1665~MHz line ratios, and the 1667~MHz line widths suggest that the strong OHM emission in various regions can be explained by unsaturated masing clouds overlapping in both space and velocity."
    },
    "0608/astro-ph0608303_arXiv.txt": {
        "abstract": "The cosmic neutrino background (CNB) consists of low-energy relic neutrinos which decoupled from the cosmological fluid at a redshift $z \\sim 10^{10}$. Despite being the second-most abundant particles in the universe, direct observation remains a distant challenge. Based on the measured neutrino mass differences, one species of neutrinos may still be relativistic with a thermal distribution characterized by the temperature $T\\sim 1.9$K. We show that the temperature distribution on the sky is anisotropic, much like the photon background, experiencing Sachs-Wolfe and integrated Sachs-Wolfe effects. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608135_arXiv.txt": {
        "abstract": "{}{This paper reports the photospheric, magnetic and circumstellar gas characteristics of the magnetic B8p star 36 Lyncis (HD 79158).}{Using archival data and new polarised and unpolarised high-resolution spectra, we redetermine the basic physical properties, the rotational period and the geometry of the magnetic field, and the photospheric abundances of various elements.}{Based on magnetic and spectroscopic measurements, we infer an improved rotational period of $3.83475\\pm 0.00002$~d. We determine a current epoch of the longitudinal magnetic field positive extremum (HJD 2452246.033), and provide constraints on the geometry of the dipole magnetic field ($i\\geq 56\\degr$, $3210~{\\rm G}\\leq B_{\\rm d}\\leq 3930$~G, $\\beta$ unconstrained). We redetermine the effective temperature and surface gravity using the optical and UV energy distributions, optical photometry and Balmer line profiles ($T_{\\rm eff}=13300\\pm 300$~K, $\\log g=3.7-4.2$), and based on the Hipparcos parallax we redetermine the luminosity, mass, radius and true rotational speed ($L=2.54\\pm 0.16~L_\\odot, M=4.0\\pm 0.2~M_\\odot, R=3.4\\pm 0.7~R_\\odot, v_{\\rm eq}=45-61.5$~\\kms). We measure photospheric abundances for 21 elements using optical and UV spectra, and constrain the presence of vertical stratification of these elements. We perform preliminary {preliminary} Doppler Imaging of the surface distribution of Fe, finding that Fe is distributed in a patchy belt near the rotational equator. Most remarkably, we confirm strong variations of the H$\\alpha$ line core which we interpret as due to occultations of the star by magnetically-confined circumstellar gas.}{}{} ",
        "introduction": "The magnetic Bp stars exhibit a rich variety of magnetic and wind phenomena, many of which are unique to this class of star. Landstreet \\& Borra (1978), detecting for the first time a magnetic field in a helium strong star, proposed that the photometric and spectroscopic variations of $\\sigma$~Ori E resulted from hot gas trapped above the magnetic equator. Brown et al. (1985) discovered unexpectedly strong C~{\\sc iv}, Si~{\\sc iv} and N~{\\sc v} ultraviolet resonance line absorption and emission that varied with the rotational period of the magnetic helium weak star HD 21699. Subsequent work revealed similar phenomena in a number of other Bp stars, which was attributed to rotational modulation of magnetically-structured stellar winds (e.g. Shore 1987, Shore et al. 1990, Shore \\& Brown 1990, Bolton 1994, Donati et al. 2001, Smith \\& Groote 2001, Neiner et al. 2003). The additional discovery of non-thermal radio emission (Drake et al. 1987) from a subset of these objects resulted in the proposal that these stars are surrounded by an extended wind-generated magnetosphere (Linsky et al. 1992), in which the radio luminosity results from optically thick gyrosynchrotron emission and should therefore correlate with both magnetic field strength and effective temperature. According to Drake et al. { (2002)}, by 2002 about 120 magnetic CP stars had been observed in the radio, and of these about 25\\% had been confidently detected. The general existence of this phenomenon, and the establishment of correlations between effective temperature, magnetic field strength, rotational period and radio emission suggest that this phenomenon is intrinsic to Ap and Bp stars, and is related to their magnetic properties. Subsequently, X-ray emission was discovered (Drake et al. 1994) from a small number of these objects. As pointed out by Drake et al. { (2002)}, no clear correlation of X-ray brightness with other physical or observational characteristics (including radio luminosity) has been established, suggesting (at least in the A-type magnetic stars) that X-ray emission may be due to the presence of late-type companions, and unrelated to their magnetic properties\\footnote{There exists compelling evidence, in the form of variability of the X-ray flux according to the stellar rotational period, that some magnetic B and O stars (the Bp star AB Aur: Babel \\& Montmerle (1997); the magnetic O7V star $\\theta^1$~Ori C: Gagn\\'e et al. (1997), Donati et al. 2002) are intrinsic X-ray emitters. For the majority of magnetic upper-main sequence stars, however, this is not the case.}. More recently, sophisticated models describing the hydrodynamic interaction between the magnetic field and wind flow have been developed (Babel \\& Montmerle 1997, ud-Doula \\& Owocki 2002, Preuss et al. 2004, Townsend \\& Owocki 2005), leading to the possibility of combining the various individual kinds of observations (magnetic, radio, optical, UV, X-ray) into a  single coherent picture of the wind-magnetic field interaction and subsequent circumstellar structure of magnetic Bp stars.  36 Lyncis (HD 79158) is a bright ($V=5.3$), well-studied magnetic (Borra, Landstreet \\& Thompson 1983) helium weak B star. It is classified as B8IIIpMn in the Bright Star Catalogue (Hoffleit et al. 1995), as a Si star by Molnar (1972), as a Mn star by Cowley (1972), and as a Sr Ti helium weak star by Borra \\& Landstreet (1983). It was first identified as peculiar by Edwards (1932). The spectrum and abundances have been studied by Searle \\& Sargent (1964), Mihalas \\& Henshaw (1966) and Sargent, Greenstein \\& Sargent (1969). According to these studies, lines of neutral He and O are weak, whereas those of C, Ne, Si, Mg and Cr are only marginally anomalous. Lines of Ti and Fe are anomalously strong. HD 79158 is a member of the {\\em sn} class (Abt \\& Levato 1978), which means its Si~{\\sc ii} lines are sharp, whereas He~{\\sc i} lines are nebulous. Ryabchikova \\& Stateva (1995) { concluded} that 36 Lyn is a classical helium weak star.  Using IUE spectra, Sadakane (1984) reported the presence of strong lines coincident with the position of the C~{\\sc iv} $\\lambda\\lambda 1548, 1550$ and Si~{\\sc iv} $\\lambda\\lambda 1394, 1403$ resonance doublets. Shore, Brown \\& Sonneborn (1987) obtained additional magnetic field measurements and IUE spectra of 36 Lyn and two additional {\\em sn} stars, confirming the presence of the ``superionised'' C~{\\sc iv} and Si~{\\sc iv} resonance lines and detecting variability of these features. Shore et al. (1990) extended these observations and analysis, obtaining 22 new measurements of the longitudinal magnetic field and determining a rotational period of $3.8345\\pm 0.001$ days.  The goal of the present paper is to provide an accurate reassessment of the magnetic and photospheric properties of 36 Lyn as the basis for a detailed study of properties of the circumstellar matter (Smith et al. 2006, ``Paper II''). ",
        "conclusions": "We have performed a thorough reevaluation and analysis of the physical and spectroscopic characteristics of the B8p He-weak star 36 Lyncis. Based on measurements of the longitudinal magnetic field, as well as measurements of the core equivalent width of the H$\\alpha$ line, we have redetermined the rotational period $P_{\\rm rot}=3.83475\\pm 0.00003$ days. An examination of the variability of the H$\\alpha$ core equivalent width reveals strong changes coincident with magnetic crossover phases, which we interpret as due to occultation of the stellar disc by magnetically-confined circumstellar gas. We have determined the effective temperature, radius, mass, age, projected and true rotational velocities, the rotation axis inclination, and we have summarised evidence for binarity, radio and X-ray emission.  {\\rm We have found { systematic and significant} discrepancies between the temperature and gravity determined from the optical and UV energy distribution and from the Balmer lines, { concluding that they probably show true limitations of the ATLAS9 model atmospheres in describing the real photospheric structure of 36 Lyn.} The ATLAS9 models we have employed are calculated using scaled solar metallicities, while individual abundances in the atmosphere of 36 Lyn show large diversity. To check the influence of anomalous abundances on the model atmosphere structure, a few model atmospheres were calculated with {\\sc LLmodels} model atmosphere code (Shulyak et al. \\cite{LL04}) which takes into account the contribution of individual line opacities to the total continuum opacity through direct synthetic calculations. This code allows us to examine the influence of the particular abundance table of 36 Lyn on the model atmosphere structure. The flux distribution for an {\\sc LLmodels} model with $T_{\\rm eff}$=13300, $\\log g=4.0$ and the abundances of 36 Lyn actually provides a somewhat better simultaneous fit to both the energy distribution and Balmer lines. We conclude that a model atmosphere with individualised abundances improves the overall fit to the observables, { although discrepancies are still evident} (a surface gravity below 4.0 is still required to fit the H lines). { Possibly we are seeing the influence of peculiar and possibly stratified abundances on the photospheric temperature and pressure structure - effects which are not included in {\\sc LLmodels}.}  We have modeled Least-Squares Deconvolved Stokes $V$ profiles and the longitudinal field variation to constrain the magnetic field geometry. We have analysed both optical and UV spectra and, using spectrum synthesis techniques, we have determined the abundances of 21 elements. We find no evidence for vertical stratification of most elements, although for a few elements (e.g. Mn, Co, Ni, Ga) we have derived abundances in the optical which differ from those in the UV, suggesting that these elements may be stratified. {\\rm The influence of using an {\\sc LLmodels} individualised abundance model atmosphere on inferred abundances is small ($\\pm$0.1 dex) for all elements except He, for which the abundance is decreased by 0.4 dex. On the other hand, using the individualised abundance model decreases by 0.15 dex the difference between the optical and UV Ga abundances (Ga optical abundance is decreased by 0.1 dex and Ga UV abundance increased by 0.05 dex).}  We find convincing evidence exists for nonuniform surface (horizontal) distributions of some elements. Using the Doppler Imaging technique, we have constructed { an approximate} surface map of the Fe abundance distribution, which is characterised by a patchy ring-like structure between $-30\\degr$ and $+30\\degr$, and a contrast of somewhat more than 3 dex. The results reported in this paper set the stage for a detailed study of the structure and geometry of the circumstellar material in Paper II."
    },
    "0608/astro-ph0608653_arXiv.txt": {
        "abstract": "SAX J2103.5+4545 has been continuously monitored for $\\sim $ 900 days by Rossi X-ray Timing Explorer (RXTE) since its outburst in July 2002. Using these observations and previous archival RXTE observations of SAX J2103.5+4545, we refined the binary orbital parameters and find the new orbital period as P= (12.66536 $\\pm $ 0.00088) days and the eccentricity as 0.4055$\\pm$ 0.0032. With these new orbital parameters, we constructed the pulse frequency and pulse frequency derivative histories of the pulsar and confirmed the correlation between X-ray flux and pulse frequency derivative presented by Baykal, Stark and Swank (2002).  We constructed the power spectra for the fluctuations of pulse frequency derivatives and found that the power law index of the noise spectra is 2.13 $\\pm$ 0.6. The power law index is consistent with random walk in pulse frequency derivative and is the steepest among the HMXRBs. X-ray spectra analysis confirmed the inverse correlation trend between power-law index and X-ray flux found by Baykal, Stark and Swank (2002).  {\\bf{Keywords:}} X-rays:binaries; Stars:neutron; pulsars:individual:SAX J2103.5+4545; accretion, accretion disks ",
        "introduction": "The transient X-ray source SAX J2103.5+4545 was discovered by {\\it{BeppoSAX}} during its outburst between 1997 February and September with 358.61s pulsations and a spectrum consistent with an absorbed power law model with the photon index of $\\sim1.27$ and the absorption column density of $\\sim3.1\\times 10^{22}$cm$^{-2}$ (Hulleman, in't Zand, \\& Heise 1998)  Starting with the November 1999 outburst detected by the {\\it{all-sky monitor (ASM)}} on the {\\it{Rossi X-ray Timing Explorer (RXTE)}},  SAX J2103.5+4545 has been continuously monitored for more than a year through regular pointed {\\it{RXTE}} observations. From RXTE observations, the orbital period and eccentricity of the orbit were found to be 12.68 days and 0.4  (Baykal, Stark, \\& Swank 2000a,b). In the timing analysis, the source was initially found to be spinning up for $\\sim 150$ days, at which point the flux dropped quickly by a factor of $\\simeq 7$, and a weak spin-down began afterwards (Baykal, Stark, \\& Swank 2002). Significant correlation between X-ray flux and spin-up rate was explained by using the Ghosh \\& Lamb (1979) accretion disk model. The X-ray spectra well fitted the absorbed power law model with high energy cutoff and a $\\sim $6.4 keV fluorescent emission line (Baykal et al. 2002).  Orbital parameters found by using {\\it{RXTE}} observations of the source (Baykal et al. 2000a,2000b) indicated that the source has a high mass companion. The optical companion of SAX J2103.5+4545 was recently discovered to be a Be type star with a visual magnitude of $\\sim 14.2$ (Reig, Neguerela, Fabregat et al.2004; Flippova, Lutovinov, Shtykovsky et al. 2004).  SAX J2103.5+4545 had another outburst on July 2002 and started to be continuously monitored by {\\it{RXTE}}. Using $\\sim 2$ months of {\\it{RXTE}}-PCA observations around the simultaneous coverage of {\\it{RXTE}} and {\\it{XMM-Newton}} on January 6, 2003, spin period and spin-up rate of the source were found to be $(354.7940\\pm 0.0008)$s  and $(7.4\\pm 0.9)\\times 10^{-13}$ Hz.s$^{-1}$ (Inam, Baykal, Swank et al. 2004a). Using this simultaneous coverage, Inam et al. (2004a) discovered 22.7s quasi-periodic oscillations with an rms fractional amplitude of $\\sim6.6$\\% and a soft spectral component consistent with a blackbody emission with $kT\\sim 1.9$ keV and an emission radius of $\\sim 0.3$ km.  SAX J2103.5+4545 was also observed with the {\\it{INTEGRAL}} observatory in the 3-200 keV band resulting in a significant detection up to $\\sim 100$ keV (Lutovinov, Molkov,\\& Revnivtsev 2003; Flippova, Lutoviniv, Shtykovsky et al. 2004; Play, Reig, Martinez Nunez et al. 2004; Falanga, di Salva, Burderi et al. 2005; Sidoli, Mereghetti, Larsson et al. 2005). The spectral parameters found in the {\\it{INTEGRAL}} observations of the source were found to be compatible with those found by Baykal et al. (2002). ",
        "conclusions": "A correlation between pulse frequency derivative and X-ray flux has been observed in the outbursts of many X-ray pulsar transient systems. These systems are EXO 2030+375 (Parmar, White $\\&$ Stella 1989, Parmar et al. 1989, Reynolds et al. 1996), 2S 1417-62 (Finger, Wilson  $\\&$ Chakrabarty 1996, Inam et al., 2004b), GRO J1744-28 (Bildsten et al. 1997), XTE J1543-568 (In't Zand, Corbet $\\&$ Marshall 2001), KS 1947+300 (Morgan et al. 2002; Tsygankov \\& Lutovinov 2005) and SAX J2103.5+4545 (Baykal, Stark  $\\&$ Swank 2002). All of these sources have shown correlations between spin-up rate and X-ray flux. Among them only SAX J2103.5+4545 has been observed to have both spin-up and spin-down episodes. The correlation between spin-up/down and X-ray flux can be explained by accretion from an accretion disk. Both outbursts  of SAX J2103.5+4545 started with a spin-up trend, made a transition to a steady spin rate and then appeared to just begin a spin-down trend (see Figure 4 and Baykal et al. 2002). In Figure 5, we present the X-ray flux and pulse frequency derivative correlations including the results found in this paper and in Baykal et al. (2002).   If the accretion is via an accretion disk then the Keplerian rotation of the disk is disrupted by the the magnetosphere at the inner disk edge. Then the plasma is forced to accrete along the magnetic field lines.  The inner disk edge $r_{0}$ moves inward with increasing mass accretion rate. The dependence of the inner disk edge $r_{o}$ on the mass accretion rate $\\dot M$ may approximately be expressed as (Pringle $\\&$ Rees 1972, Lamb, Pethick, $\\&$ Pines 1973) \\begin{equation} r_{o}= K \\mu^{4/7}(GM)^{-1/7}\\dot M^{-2/7} \\end{equation} where $\\mu\\simeq BR^{3}$ is the neutron star magnetic moment with $B$ the magnetic field and  $R$ the neutron star radius, G is the gravitational constant, and $M$ is the mass of the neutron star. In this equation $K=0.91$ gives the Alfven radius for spherical accretion. Then the torque estimate is given by Ghosh \\& Lamb (1979) as \\begin{equation} 2\\pi I \\dot \\nu = n(w_{s}) \\dot M~l_{K}, \\end{equation} where $I$ is the moment of inertia of the neutron star, $l_{K} = (GM)r_{o})^{1/2}$ is the specific angular momentum added by a Keplerian disk to the neutron star at the inner disk edge $r_{o} $; \\begin{equation} n(w_{s}) \\approx 1.4 (1-w_{s}/w_{c})/(1-w_{s}) \\end{equation} is a dimensionless torque which is a measure of the variation of the accretion torque as estimated by the fastness parameter \\begin{equation} w_{s} =\\nu /\\nu _{K}(r_{o}) = (r_{o}/r_{co})^{3/2} = 2 \\pi K^{3/2} P^{-1}  (GM)^{-5/7} \\mu ^{6/7} \\dot M^{-3/7}, \\end{equation} where $r_{co}=(GM/(2\\pi\\nu)^{2})^{1/3}$ is the corotation radius at which the centrifugal forces balances the gravitational forces, $w_{c}$ is the critical fastness parameter at which the accretion torque is expected to vanish. The critical fastness parameter $w_{c}$ has been estimated to be $\\sim$ 0.35 and depends on the electrodynamics of the disk (Ghosh \\& Lamb 1979, Wang 1987, Ghosh 1993, Torkelsson 1998, Li $\\&$ Wickramasinghe 1998, Dai $\\&$ Li 2006).  The accreted material will produce X-ray emission at the neutron star surface with a luminosity expressed as \\begin{equation} L =  \\eta GM \\dot M /R \\end{equation} where $\\eta \\le 1 $ is the efficiency factor. From Equations 2,3 and 6, the rate of spin-up is related to the X-ray intensity through \\begin{equation} \\dot \\nu \\propto n(w_{s}) L^{6/7} =  n(w_{s}) (4 \\pi d^{2} F)^{6/7}, \\end{equation} where $d$ is the distance to the source and $F$ is the X-ray flux.  The distance of the source from optical observations was found to be $6.5\\pm 0.9$ kpc (Reig et al. 2004; Reig et al. 2005). Although cyclotron emission has not been seen from this source, the hard spectrum and cut off energy suggests that the  magnetic field is greater than $10^{12}$ Gauss (Lutovinov, Molkov,\\& Revnivtsev 2003; Flippova, Lutovinov, Shtykovsky et al. 2004; Falanga, di Salva, Burderi et al. 2005; Sidoli, Mereghetti, Larsson et al. 2005). We fitted the Ghosh \\& Lamb model setting power law index, magnetic field and distance as a free parameter. We obtained for the distance to the source  $4.5\\pm 0.5$ kpc and for the magnetic field $(16.5\\pm 2.5)\\times 10^{12}$ Gauss. We obtained a power-law index $0.96 \\pm 0.14$ which is consistent with the value 6/7 expected in the model. The large uncertainty in this power law index shows that radiation pressure dominated accretion disk model with power law index 0.925 is also possible (see Ghosh 1996). Since the peak outburst luminosity ($\\sim 2\\times 10^{36}$erg/sec, assuming a distance of 4.5kpc) is far below the Eddington limit ($\\sim 10^{38}$ erg/sec), we consider that gas pressure accretion disk with power law index 6/7 model is  the appropriate model for this source. It is also seen in Figure 5, the points corresponding to spin-up and spin-down are consistent with torque model.  Our distance estimate is lower by $\\sim 2 \\sigma$  than the value obtained from optical observations. The discrepancy could arise from the efficiency factor between accretion X-ray luminosity and observed luminosity. Another factor could be the uncertainty of interstellar absorption coefficient used in optical observations (see Reig et al. 2004 and references therein)  The noise torque fluctuations of  accreting X-ray pulsars are subject to both external and internal torques. The external torques arise from both the accretion of matter from the binary companion and the interaction of the pulsar magnetosphere with accretion disk. The internal torques arise from the response of the neutron star to these external torques and depend on the interior structure of neutron stars. However due to the pulse shape noise and consequent instrumental noise, torque fluctuations at short time scales (or analysis frequencies) could not be resolved so far (Baykal, Alpar, Kiziloglu 1991; Baykal 1997). Bildsten et al., (1997) constructed noise power spectra of HMXRBs using BATSE/CGRO pulsar observations. They found that for Vela X-1, Her X-1, 4U 1538-52 and GX 301-2 pulse frequency derivative fluctuations ($P_{\\dot \\nu}$) are consistent with a white torque noise model. For OAO 1657-415, GX 1+4 and Cen X-3, they showed that red torque noise for which $P_{\\dot \\nu}$  is proportional to $f^{-1}$. Our results indicate that for SAX J 2103.5+4545 $P_{\\dot \\nu}$ is proportional to  $f^{-2}$. In other words torque fluctuations are characterized as first order random walk in pulse frequency derivatives. The main reason of this might be the formation of transient accretion disks which leads to the formation of step like torque fluctuations on SAX J2103.5+4545. For higher analysis frequencies of $\\sim 1/13 d^{-1}$, instrumental noise due to the pulse shape fluctuations limits the pulse timing and hence cut-off the torque noise fluctuations.   It was shown for SAX J2103.5+4545 that there is $\\sim 50$\\% X-ray flux modulation within the time span of orbital period at the bright part of observations (see Figure 8 of Baykal, Stark, Swank 2000, and Figure 4 of Baykal, Stark, Swank 2002). Assuming that the pulse frequency and X-ray spectra correlations as seen in Figure 5 is valid at short time scales in single orbit we estimate the phase offsets using the Ghosh and Lamb torque model as in  $\\delta \\phi \\sim \\int \\int \\dot \\nu dt$ $\\sim \\int \\int  L^{6/7} dt \\sim 0.0027 $. This phase offset corresponds to the 1/0.0027=358 part of pulse phase. Typical uncertainty of pulse phases is $\\sim 0.04$ which is much larger than the phase offsets of torque fluctuations. Therefore we could not resolve the torque fluctuations within the orbital period time scale.  From the X-ray spectral analysis of SAX J2103.5+4545, we found that two spectral models fit almost equally well to data. The first spectral model consisting of absorbed power law with an high energy cut-off is a typical model for accretion powered pulsars (White et al. 1983). This spectral model has been found to well fit X-ray spectrum of SAX J2103.5+4545 before (Hulleman et al. 1998; Baykal et al. 2002; Lutovinov et al. 2003; Flippova et al. 2004; Play et al. 2004; Falanga et al. 2005; Sidoli et al. 2005)  The second spectral model has an additional soft component modeled as a blackbody. The model with a soft component was recently found to better fit the X-ray spectrum of SAX J2103.5+4545  when energies lower than 3 keV were included (Inam et al. 2004a). While fitting the spectra using this model, we fixed the blackbody peak energy at 1.9 keV (typical value found by Inam et al. 2004a) since our energy range does not cover around 1.9 keV and large uncertainties of this parameter might have arised. The blackbody component in accretion powered pulsars comes either from the reprocessed emission of the surrounding material or from the polar caps of the neutron stars (see Inam et al. 2004a and references therein). Inam et al. 2004a found that the polar cap interpretation is more likely for the blackbody emission of SAX J2103.5+4545. For the distance of 4.5kpc and considering the variation of the blackbody normalization, we estimated the radius of blackbody emitting region to be varying between $\\sim 0.01-0.47$ km which is consistent with $\\sim 0.3$ km value found by Inam et al. 2004a.  For both of the spectral models, we found that power law index anti-correlates with the unabsorbed X-ray flux  (see Figure 7). This anti-correlation indicates that the spectrum gets softer with the decreasing X-ray flux. This anti-correlation was previously found in the dataset after the 1999 outburst of SAX J2103.5+4545 (Baykal et al. 2002), and in the dataset after the 1999 outburst of 2S 1417-62 (Inam et al. 2004b). This type of spectral softening accompanied with decreasing flux was found to be primarily a consequence of mass accretion rate change and is not necessarily related to a significant accretion geometry change (Mezsaros et al. 1983; Harding et al. 1984).  {\\bf"
    },
    "0608/astro-ph0608379_arXiv.txt": {
        "abstract": "It is demonstrated here that if the prompt GRB emission is produced by the simplest version of the external shock model, a specific relation should prevail between the observed duration, isotropic equivalent energy, and photon peak energy. In essence, this relation arises because both the burst duration and the typical energy of the emitted synchrotron photons depend on the same combination of the, usually poorly constrained, external density at the deceleration radius, $n_{\\rm dec}$, and initial bulk Lorentz factor, $\\Gamma_0$. This has the fortunate consequence of making the relation independent of both $\\Gamma_0$ and $n_{\\rm dec}$.  Unless the efficiency of electron acceleration is very low, synchrotron gamma-rays from the external shock would fail to meet the current observational constraints for the vast majority of GRBs, including those with a smooth, single peak temporal profile. This argues either against an external shock origin for the prompt emission in GRBs or for changes in our understanding of the microphysical and radiation processes occurring within the shocked region. ",
        "introduction": "The simplest version of the standard fireball model for gamma-ray bursts (GRBs) involves a spherical explosion taking place in a uniform or a stratified surrounding medium. When an explosion deposits a large amount of energy into material with a much smaller amount of rest energy within a compact volume, an ultra-relativistic pair fireball is formed \\citep{CR78,Pac86,Goodman86}. The large pressure of the explosion causes the fireball to expand, and the thermal energy of the explosion is transformed into bulk kinetic energy due to strong adiabatic cooling of the particles in the comoving frame. Because of the Thomson coupling between the particles and photons, most of the original explosion energy is eventually carried by the baryons that were originally mixed into the explosion \\citep{SP90}.  This bulk kinetic energy cannot be efficiently radiated as gamma rays unless it is converted back to internal energy (i.e. the velocities of the protons must be re-randomized). This requires shocks, and in order to tap a reasonable fraction of the total kinetic energy, the shocks must be (at least mildly) relativistic.  Impact on an external medium would randomize about half of the initial energy merely by reducing the expansion Lorentz factor by a factor of $\\sim 2$ \\citep{rm92}. Alternatively, internal shocks may form within the outflow: for instance, if the Lorentz factor of the outflow varied by a factor $>2$, then the shocks that developed when fast material overtakes slower material would be internally (at least mildly) relativistic. There is a general consensus that the longer complex bursts must involve internal shocks, while simple smooth profiles could arise from an external shock interaction \\citep{SP97a,SP97b,rf00,np02,mkp04,rm01}. The latter would in effect be the beginning of the afterglow.  An external shock moving into a medium with a smooth density profile would naturally result in a burst with a simple time-profile. Angular variations within the outflow might still cause variability in the light curve, but variations on very small angular scales ($\\theta < \\Gamma_0^{-1}$, where $\\Gamma_0$ in the initial Lorentz factor) are required in order to produce the large variability of the prompt GRB emission \\citep{f99,dm99}.  A blobby external medium could produce significant variability only if the covering factor of blobs is low, implying modest efficiency. Furthermore, the resulting variability in the light curve would be small if produced close to or after the deceleration radius, or if the portion of the ejecta that collides with a blob is decelerated significantly \\citep{NG06}.  The purpose of this {\\it Letter} is to demonstrate that if the prompt emission is produced by the simplest version of the external shock model, this implies a specific relation between the observed duration, isotropic equivalent energy (or luminosity), and photon peak energy, which is apparently incompatible with observations.  This relation is derived in \\S~\\ref{sec:ES} and compared to observations in \\S~\\ref{sec:obs}. The implications are discussed in \\S~\\ref{sec:dis} along with possible caveats. ",
        "conclusions": "\\label{sec:dis}  It has been shown that the simplest version of the external shock model implies a relation between different observed quantities of the GRB (Eq.~\\ref{psi}), which can conveniently be expressed in the form $\\xi_e \\lesssim \\xi_{e,{\\rm max}} = \\Psi^{-1/2}$, where $\\xi_e$ is the number of accelerated electrons per proton. Naively, for the standard assumption that $\\xi_e \\approx 1$, one would expect $\\Psi \\sim 10^{-2}-10^{-1}$ for typical values of $\\epsilon_e \\sim 0.1$ and $10^{-3} \\lesssim \\epsilon_B \\lesssim 0.1$. It is conceivable, however, that only a small fraction of the electrons participate in the acceleration process (i.e. $\\xi_e \\ll 1$).\\footnote{It should be pointed out that in principle even $\\xi_e > 1$ is possible, especially near $R_{\\rm dec}$, due to pair enrichment of the ambient medium from pair production by gamma-ray photons that are scattered on the external medium \\citep{rml02,mrrr01,KP04,B02,TM00}.}  A comparison with observations shows, however, that $\\Psi \\gg 1$ (and $\\xi_{e,{\\rm max}} \\ll 1$) for the vast majority of GRBs (Figs.~\\ref{fig1} and \\ref{fig2}).  This implies that GRBs could arise from synchrotron emission in the external shock only if the efficiency of electron acceleration in relativistic collisionless shocks is very low ($\\xi_e \\ll 1$). An external shock origin might still be possible if the radiation process responsible for the gamma rays is other than than synchrotron radiation \\citep[e.g.,][]{Wang06}. Alternatively, the prompt gamma-ray emission might arise from completely different mechanism, such as internal shocks \\citep[e.g.][]{rl02}.  It should be noted that afterglow observations already provide interesting constraints on the efficiency of electron acceleration \\citep{EW05}.  Current observations imply that the characteristic energy of accelerated electrons is comparable to the proton post-shock temperature. They also imply that the efficiency $\\xi_e$ is similar for highly relativistic and sub-relativistic shocks and plausibly suggest that $\\xi_e \\sim 1$. However, even values of $\\xi_e$ as low as $\\sim m_e/m_p$ cannot be ruled out, since currently testable afterglow predictions remain unchanged for $(E_{\\rm iso},n_{\\rm ext}) \\to (E_{\\rm iso},n_{\\rm ext})/\\xi_e$ and $(\\epsilon_e,\\epsilon_B) \\to \\xi_e(\\epsilon_e,\\epsilon_B)$ for any $\\xi_e$ in the range $m_e/m_p \\leq \\xi_e \\leq 1$ \\citep{EW05}.  Estimates of the energy in the afterglow shock from late time radio observations when the flow is only mildly relativistic and starts to approach spherical symmetry \\citep[often called ``radio calorimetry'';][]{FWK00,BKF04,onp04,Frail05,grl05} typically yield $E_k \\sim 10^{51.5}\\;$erg assuming $\\xi_e = 1$. However, as noted by \\citet{EW05}, afterglow observations actually constrain $\\xi_e E_k$ rather than $E_k$. The true kinetic energies at late times are thus given by $E_k \\sim 10^{52.5}(\\xi_e/0.1)^{-1}\\;$erg. The initial energy content of the outflow could be even larger due to early radiative losses (i.e., during the prompt GRB and early afterglow stages).  It is difficult to accurately account for the magnitude of such losses, as they depend on poorly known questions about postshock energy exchange between protons and electrons. Nevertheless, a lower limit on the radiated energy is given by $E_\\gamma = E_{\\rm\\gamma,iso}(1-\\cos\\theta_0) \\approx E_{\\rm\\gamma,iso}\\theta_0^2/2$ (additional energy may be radiated outside the observed photon energy range, or during the early afterglow). Other possible channels of energy loss are the escape of accelerated non-thermal protons from the blast wave (high energy cosmic rays) or the production of high energy neutrinos via pion decay. Any added losses would inevitably lead to a further increase in the energy requirements.  Therefore, very small values of $\\xi_e$ would imply very large energy contents.  Another test of the simple external shock model is provided by a comparison of the correlation found by \\citet{Firmani06} between the isotropic equivalent luminosity, the burst duration, and the peak energy, with that predicted by Eq.~\\ref{psi0}.  In the cosmological frame of the GRB this correlation reads $L_{\\rm iso} \\propto E_{\\rm p}^{1.62\\pm0.08}T_{0.45}^{-0.49\\pm0.07}$, where $T_{0.45}$ is defined by \\citet{Firmani06} to be the time during which 45\\% of the counts above background are measured (which is expected to scale linearly with $T_{\\rm GRB}=T_{90}$). This is in disagreement with Eq.~\\ref{psi0}, which, for a reasonably small scatter in $\\epsilon_B^{1/2}(\\epsilon_e/\\xi_e)^2$, gives $L_{\\rm iso} \\propto E_{\\rm p}^2 T_{\\rm GRB}^2$.  It is natural to hope that the values of $\\epsilon_B$, $\\epsilon_e$, $p$ and $\\xi_e$ are universal, since they are determined by the microphysics of the collisionless shock.  However, the wide distribution of $\\Psi$ values seen in Fig.~\\ref{fig2} suggests otherwise. That is, in the simplest version of the external shock model, a large scatter in $\\epsilon_B^{1/2}(\\epsilon_e/\\xi_e)^2$ is required. The presence of a significant number of non shock-accelerated electrons in the external shock ($\\xi_e \\ll 1$) appears to be more prominent for the sub-sample of bursts found to be spectroscopically associated with a supernova (Fig.~\\ref{fig2}), most of which have a smooth temporal profile \\citep{yuki}. The low values of $\\xi_e$ do not, however, increase the total energy requirements to unreasonable values for these events as they have rather low values of $E_{\\rm iso}$. Under this interpretation, a wide range of shock microphysical parameters may be the rule, rather than the exception.  In conclusion, observations of the prompt emission in GRBs with known redshifts, which are becoming far more accessible in the {\\it Swift} era, can provide an important diagnostic of the external shock model. Current observational constraints do not allow for efficient electron acceleration in the external shock, if its synchrotron emission produces the observed prompt gamma-ray emission. Although there is no a priori reason to suspect that $\\xi_e$ should be large, $\\xi_e \\ll 1$ would dramatically increase the total kinetic energy budget."
    },
    "0608/physics0608034_arXiv.txt": {
        "abstract": "Scintillation light from gamma ray irradiation in liquid xenon is detected by two Hamamatsu R9288 photomultiplier tubes (PMTs) immersed in the liquid. UV light reflector material, PTFE, is used to optimize the light collection efficiency. The detector gives a high light yield of 6 photoelectron per keV (pe/keV), which allows efficient detection of the 122 keV $\\gamma$-ray line from $^{57}$Co, with a measured energy resolution of $(8.8\\pm0.6)$\\% (\\(\\sigma\\)). The best achievable energy resolution, by removing the instrumental fluctuations, from liquid xenon scintillation light is estimated to be around 6-8\\% ($\\sigma$) for $\\gamma$-ray with energy between 662 keV and 122 keV. ",
        "introduction": "Liquid xenon (LXe) has been used since many years ago for its good ionization properties \\cite{Carugno:NIM96,Aprile:LXeGRIT}, while applications based on its scintillation light are not fully developed until recently mostly due to inefficient detection of the UV scintillation light at 175 nm. Some of the recent applications, based on the LXe scintillation light, are using photomultiplier tubes (PMTs) in the liquid (e.g. MEG experiment \\cite{MEG} and the ZEPLIN-III detector \\cite{ZEPLINIII}) or optical reflectors (e.g. XMASS experiment \\cite{Yamashita:04} and the ZEPLIN-II detector \\cite{ZEPLINII}) to achieve high light collection efficiency. With the improved light collection, better energy resolution and lower energy threshold can be achieved. Our group started to develop a new system with PMTs immersed in the liquid and with PTFE UV light reflectors to increase the light collection efficiency \\cite{Aprile:IEEE03}, as part of the R\\&D program for the XENON dark matter search \\cite{XENON}. In this paper, we report the studies of light yield and energy resolution from LXe scintillation light for different $\\gamma$-ray energies, from a small detector with optimized light collection efficiency. ",
        "conclusions": "In this paper we presented the results related to scintillation spectroscopy obtained from a liquid xenon chamber with a good light collection efficiency by immersing PMTs in the liquid and using UV light reflectors. The relative scintillation yield from gamma rays to relativistic heavy ions in liquid xenon is measured. The LET dependence of relative scintillation yield from $\\gamma$ rays is discussed. The high light collection efficiency yields a good energy resolution of 8.8\\%($\\sigma$) for 122 keV $\\gamma$ rays from $^{57}$Co at zero field. At zero field, the best energy resolution (i.e. without instrumental fluctuations) that can be achieved, by using liquid xenon scintillation light alone, is estimated to be around $6-8\\%$($\\sigma$) for $\\gamma$ rays between 662 keV and 122 keV."
    },
    "0608/astro-ph0608715_arXiv.txt": {
        "abstract": "Free of any atmospheric contamination, the \\textit{Hubble Space Telescope} provides the best available spectrophotometry from the far-UV to the near-IR for stars as faint as $V\\sim16$. The \\textit{HST} CALSPEC standard star network is based on three standard candles: the hot, pure hydrogen white dwarf (WD) stars G191B2B, GD153, and GD71, which have Hubeny NLTE model flux calculations that require the atomic physics for only one atom. These model flux distributions are normalized to the absolute flux for Vega of $3.46\\times10^{-9}$~erg cm$^{-2}$ s$^{-1}$~\\AA$^{-1}$ at 5556~\\AA\\ using precise Landolt $V$~band photometry and the $V$~bandpass function corrected for atmospheric transmission by M.~Cohen. The three primary WD standards provide absolute flux calibrations for FOS, STIS, and NICMOS spectrophotometry from these instruments on the \\textit{HST}. About 32~stellar spectral energy distributions (SEDs) have been constructed with a primary pedigree from the STIS data, which extends from 1150~\\AA\\ for the hot stars to a long wavelength limit of 1~$\\mu$m. NICMOS grism spectrophotometry provides an extension to 1.9~$\\mu$m in the IR for 17 of the \\textit{HST} standards and longward to 2.5~$\\mu$m for a few of the brighter stars. Included among these \\textit{HST} standards are Vega, the Sloan standard BD+17$^{\\circ}$4708, three bright solar analog candidates, three cool stars of type M or later, and five hot WDs. In addition, four K~giants and four main sequence A-stars have NICMOS spectrophotometry from 0.8--2.5~$\\mu$m. The WD fluxes are compared to their modeled SEDs and demonstrate an internal precision of 1--2\\%, while the A-stars agree with the Cohen IR fluxes to $\\sim$2\\%. Three solar analog candidate stars differ from the solar spectrum by up to 10\\% in the region of heavy line blanketing from 3000--4000~\\AA\\ and show differences in shape of $\\sim$5\\% in the IR around 1.8~$\\mu$m. ",
        "introduction": "To establish a set of stellar flux standards, one or more stars must be established with known absolute spectral energy distributions (SEDs). One method for establishing these so-called ``standard candles'' in the sky is to determine the composition, effective temperature, and surface gravity and then compute a model atmosphere SED. The simplest case with the best known atomic physics is for a hot pure hydrogen white dwarf (WD) star, where the temperature and gravity are determined from fits of model line profiles to the observed Balmer lines (e.g., Finley, Koester, \\& Basri 1997, Barstow et~al.\\ 2003). These relative SEDs become absolute when normalized to the measured flux for Vega using precise $V$~band photometry (e.g., Landolt 1992). Bohlin (2003) switched from LTE to NLTE for calculating the model flux distributions with the Hubeny TLUSTY code (Hubeny \\& Lanz 1995).  \\begin{figure} \\centering \\includegraphics[width=\\textwidth]{bohlin_fig1.eps} \\caption{STIS or FOS fluxes below 1~$\\mu$m and NICMOS above 0.8~$\\mu$m. The noise level is evident in the fainter stars near the STIS long wavelength limit of 1~$\\mu$m and in the NICMOS spectrum of the faint SNAP-1 beyond 1.9~$\\mu$m.} \\end{figure}  Following the establishment of the absolute SEDs for the primary stars, a flux distribution for another star is established by the response relative to the primaries when observed by a linear spectrometer of constant sensitivity. In practice, no such instrumentation exists, especially for ground-based observations through an atmosphere with a transparency that can vary on short time scales. Even in space, focus variations, sensitivity degradation with time, and non-linear detectors limit the precision of the transfer of the calibration from the primary standards to bright secondary standard stars. For fainter stars photon, statistics also contributes to the uncertainties.  From its installation in 1997 until its death in 2004, the low dispersion ($R\\sim1000$) modes of STIS were the premier instrumentation for establishing flux standards from 1150--10200~\\AA\\  (Bohlin 2000; Bohlin, Dickinson, \\& Calzetti 2001). Considerable effort has been expended in tracking the STIS changes in sensitivity with time (Stys, Bohlin, \\& Goudfrooij 2004 and unpublished updates) and in correcting the CCD observations for loss of charge transfer efficiency (CTE; Goudfooij \\& Bohlin 2006). In the 0.8--2.5~$\\mu$m range, the NICMOS $R\\sim200$ grism spectrophotometry has been calibrated and the count rate dependent non-linearity has been characterized by Bohlin, Lindler, \\& Riess (2005) and Bohlin, Riess, \\& de~Jong (2006).  \\begin{figure} \\centering \\includegraphics[angle=90,width=.89\\textwidth]{bohlin_fig2.eps} \\caption{CALSPEC stellar flux distributions as in Figure~1. The lower resolution of the NICMOS observations (red) is evident where the cool star spectra overlap with STIS (green). The solar analog SF1615+001A is dotted for clarity.} \\end{figure}  Figures 1--2 include most of the SEDs of the stars with NICMOS observations. Figure~1 shows the hotter WD stars, while Figure~2 illustrates the solar analog and cooler stars. The wavelength coverage, the spectral types, and the dynamic range are typical of the standard observational capabilities of the \\textit{HST} spectrographs. Because of the highly over subscribed demand for \\textit{HST} time, sky coverage of the \\textit{HST} CALSPEC standard star network is minimal.\\footnote{The absolute spectral energy distributions discussed in this paper are available in digital form at http://www.stsci.edu/hst/observatory/cdbs/calspec.html.} ",
        "conclusions": ""
    },
    "0608/astro-ph0608523_arXiv.txt": {
        "abstract": "We show the evolution of non-spherically symmetric balls of a self-gravitating scalar field in the Newtonian regime or equivalently an ideal self-gravitating condensed Bose gas. In order to do so, we use a finite differencing approximation of the Shcr\\\"odinger-Poisson (SP) system of equations with axial symmetry in cylindrical coordinates. Our results indicate: 1) that spherically symmetric ground state equilibrium configurations are stable against non-spherical perturbations and 2) that such configurations of the SP system are late-time attractors for non-spherically symmetric initial profiles of the scalar field, which is a generalization of such behavior for spherically symmetric initial profiles. Our system and the boundary conditions used, work as a model of scalar field dark matter collapse after the turnaround point. In such case, we have found that the scalar field overdensities tolerate non-spherical contributions to the profile of the initial fluctuation. ",
        "introduction": "\\label{sec:introduction}  Recently, scalar fields have played different roles in several scenarios related to astrophysical phenomena. The reason is that such fields are quite common in theoretical physics, specially branches related to theories beyond the standard model of particles, high dimensional theories of gravity an tensor-scalar theories alternative to General Relativity. In the present research we deal with the scalar field dark matter model (SFDM), which assumes the dark matter to be a classical minimally coupled real scalar field determined by a cosh-like potential. Such potential behaves exponentially at early stages of the universe and as a free field (quadratic potential) at late times, which provides the field with the necessary properties to mimic the behavior and successes of cold dark matter at cosmic scales. In fact in \\cite{Sahni2000,MatosUrena2000,MatosUrena2001} it was shown that the mass parameter of the scalar field gets fixed by a desired cut-off of the power spectrum, which has two effects: i) the theory gets fixed and ii) there is no overabundance of substructure, which standard cold dark matter cannot achieve. One important consequence is that the boson has to be ultralight with masses around $m \\sim 10^{-21,-23}$eV. This is a substantially important bound, because in the standard dark matter models there are no such ultralight dark matter candidates. The benefit obtained however, is two fold: the scalar field can represent a Bose Condensate of such ultralight particles and the Compton wavelength forbids the scalar field to form cuspy structures. In fact, in \\cite{FerrerGrifols2001} it was shown that in order for the scalar field interaction to become long range, the system only needs to be condensed.  After the fluctuation analysis about this candidate and its corresponding concordance with observations -the model mimics the properties of the $\\Lambda$CDM at cosmic scales-, the next step has to be in the direction of the study of structure formation and the explanation of local phenomena, like rotation curves in galaxies. Fortunately there have been important advances in such direction \\cite{Arbey2001,Mbelek2004}. About the gravitational collapse, in \\cite{AlcubierreDark2002} it was shown that relativistic self-gravitating scalar field configurations can be formed when they have galactic masses provided the mass of the boson is ultralight. Nevertheless, because the gravitational field in galaxies is weak, the race turned into the newtonian limit of the system of equations, which was developed in \\cite{GuzmanUrena2003}. The price to be paid is that it is not possible to apply the approach at very early stages of the evolution of the universe, and the profit is that the scalar field in the non-relativistic regime provides a clear interpretation within the Bose Condensate formalism and classical Quantum Mechanics. In both cases, the strong gravity and the Newtonian regimes, a wide range of arbitrary spherically symmetric initial configurations collapse and form gravitationally bounded and virialized objects with a smooth density everywhere called oscillatons (except those that are related to unstable initial configurations that collapse into black holes in the strong field regime) \\cite{GuzmanUrena2004,AlcubierreDark2003}. This property seems to be fundamental in order to form galactic halos, because several high resolution observations are consistent with regular galactic dark matter profiles in the center of the galaxies \\cite{flatcore}, which implies that this type of condensate could be an alternative to solve the cuspy density profiles of dark halos.  Two pieces of the model that are in progress are the high energy fully relativistic case, including exponential-like scalar field potentials and the free field weak energy newtonian, both are complementary and necessary to explore the SFDM hypothesis. Inspired in a late-time astrophysical scenario, at stages after the turnaround point where weak field applies and the field is free, the question is whether dark matter halos are gravitationally bounded objects of scalar field  which have been formed through a gravitational collapse of initial scalar field overdensities. The newtonian version of the Einstein-Klein-Gordon system of equations is provided by the Schr\\\"odinger-Poisson equations (SP), which are the ones that lead the gravitational collapse of the system. This approximation should work for the evolution of an initial density profile after the epoch when the overdensity fluctuation starts to evolve independently of the cosmic expansion.  In the recent past, it has been found that in spherical symmetry the SP system has equilibrium solutions of two types: stable, for which the wave function is nodeless (called sometimes ground state configurations) and others, for which the wave function has nodes (called sometimes excited configurations) that decay into ground state solutions. From these solutions only the ground state ones are stable \\cite{GuzmanUrena2004,Harrison2002}; even further, it has been found that such configurations behave as late-time attractors for initially quite arbitrary spherically symmetric density profiles \\cite{GuzmanUrena2003,GuzmanUrena2006}.  In \\cite{GuzmanUrena2003} it was shown that free scalar field overdensities after the turnaround virialize and tend to form ground state configurations. Due to a very general scale invariance of the Schr\\\"odinger-Poisson system of equations such results were also valid for a structure of arbitrary mass \\cite{GuzmanUrena2004}. In \\cite{GuzmanUrena2006} was shown that the spherical collapse of SFDM tolerates the introduction of a self-interaction term, which on the other hand is associated to the self-interaction term of a self-gravitating Bose-condensate \\cite{Wang2001}, and helps at allowing diverse sizes and masses of the final configurations.  In this paper we go a step forward and study the collapsing process of non-spherical initial configurations, in particular those involving at most quadrupolar terms. In order to achieve this goals an axisymmetric code that solves the SP system is needed. We choose to deal with cylindrical coordinates denoted by $(x,z)$, where $x$ is the radial coordinate and $z$ the axial coordinate. The SP system for the free field case in these coordinates reads  \\begin{eqnarray} i \\frac{\\partial \\psi}{\\partial t} &=& -\\frac{1}{2} \\left( \\frac{\\partial^{2} \\psi}{\\partial x^2} + \\frac{1}{x}\\frac{\\partial \\psi}{\\partial x} + \\frac{\\partial^{2} \\psi}{\\partial z^2} \\right) + U \\psi \\label{eq:schroedinger}\\\\ \\frac{\\partial^{2} U}{\\partial x^2} &+& \\frac{1}{x}\\frac{\\partial U}{\\partial x} + \\frac{\\partial^{2} U}{\\partial z^2} = \\psi^{\\ast}\\psi.\\label{eq:poisson} \\end{eqnarray}  \\noindent where $\\hbar = c = 1$ and we are using the rescaled variables $x\\rightarrow mx$, $z\\rightarrow mz$, $t \\rightarrow mt$ and the wave function $\\psi \\rightarrow \\sqrt{4\\pi G}\\psi$. This set of coupled partial differential equations in two spatial dimensions plus time is the core of the present manuscript. As mentioned before, these equations have stable solutions in spherical symmetry (see \\cite{GuzmanUrena2006} for solutions also including a non-linear term in the Schr\\\"odinger equation); such solutions have shown to be not only stable, but also late time attractors for quite arbitrary initial density profiles \\cite{GuzmanUrena2004,GuzmanUrena2006}. Therefore, our main task in the present manuscript will be to show that such solutions are still attractors even for non spherically symmetric initial density profiles.  In the next section we present the code we constructed for the present purpose. In section \\ref{sec:nonsphericalperturbations} we show how a ground state configuration reacts under non-spherical perturbations. In section \\ref{sec:nonsphericalcollapse} we show the evolution and fate of non-spherical initial density profiles. Finally in section \\ref{sec:conclusions} we draw some conclusions. ",
        "conclusions": "\\label{sec:conclusions}  We have presented a new code designed to solve the SP system with axial symmetry in cylindrical coordinates. We have shown it passes the necessary testbeds of stability and consistency with expected results related to spherically symmetric ground state equilibrium configurations.  We perturbed ground state equilibrium configurations with rather general axially symmetric shells of particles, which were proportional to the first axisymmetric spherical harmonics. Moreover, these shells were endowed with an initial speed toward the equilibrium configuration. The main result is that ground state configurations are stable against this type of perturbations.  Finally, we have evolved axisymmetric initial profiles, and showed that spherically symmetric ground state equilibrium configurations play the role of late-time attractors. This is a generalization of the same result when the initial profiles are spherically symmetric \\cite{GuzmanUrena2006}. We have thus shown that the final state of a non-spherical self-gravitating scalar field fluctuation is a ground state spherically symmetric solution. Thus we showed that such configurations are late-time attractors not only for initial spherically symmetric configurations but also for quite general initial axisymmetric shapes.  In the context of the scalar field dark matter model we have quite a new result: the collapse of overdensities tolerates an initial non-spherical contribution to the initial profiles, and moreover, initially axisymmetric profiles tend toward a spherical ground state.  About the process of virialization, with the results at hand we are unable to establish whether the gravitational cooling (powered by the emission of scalar field) is the only responsible for the relaxation of the system. The cases involving a non-zero quadrupolar contribution to the gravitational potential might allow one to formulate the study of another channel, that of the emission of gravitational waves, which would be the subject of a further investigation.  Definitely a major application of the tools shown in this paper is the demonstration of the stability properties of rotating solutions to the SP system of equations constructed in the past \\cite{Harrison2002,Schupp1996}. It remains unclear whether excited state solutions representing spinning configurations are stable, and whether such solutions could contribute to the SFDM possibilities in terms of quick virialization of collapsed structures and other dynamical properties, for instance, the non-cupsy density profile might demand a considerably important spherical contribution in the case the wave function has odd parity."
    },
    "0608/astro-ph0608009_arXiv.txt": {
        "abstract": "s{Constellation-X (Con-X) will carry out two powerful and independent sets of tests of dark energy based on X-ray observations of galaxy clusters, providing comparable accuracy to other leading dark energy probes. The first group of tests will measure the absolute distances to clusters, primarily using measurements of the X-ray gas mass fraction in the largest, dynamically relaxed clusters, but with additional constraining power provided by follow-up observations of the Sunyaev-Zel'dovich (SZ) effect. As with supernovae studies, such data determine the transformation between redshift and true distance, $d(z)$, allowing cosmic acceleration to be measured directly. The second, independent group of tests will use the exquisite spectroscopic capabilities of Con-X to determine scaling relations between X-ray observables and mass.  Together with forthcoming X-ray and SZ cluster surveys, these data will help to constrain the growth of structure, which is also a strong function of cosmological parameters.} ",
        "introduction": "The late-time accelerated expansion of the Universe is now a well measured fact \\cite{Riess:04,Allen:04,Spergel:06}. However, the underlying cause of this cosmic acceleration remains unknown. Assuming general relativity, a new energy component of the Universe, so-called dark energy, is required. In order to pin down its nature, a number of powerful future experiments are planned.  Con-X~\\footnote{\\sf http://constellation.gsfc.nasa.gov/} data will constrain dark energy with comparable accuracy and in a beautifully complementary manner to the best other techniques available circa 2018. Using a modest $\\sim10-15\\%$ (10-15Ms) investment of the available observing time over the first 5 years of the Con-X mission, we will be able to measure the X-ray gas mass fraction (or predict the Compton $y$-parameter) to $5\\%$ or $3.5\\%$ accuracy, corresponding to $3.3\\%$ or $2.3\\%$ in distance, for 500 or 250 clusters, respectively, with a median redshift $z\\sim1$.  When combined with CMB data or suitable priors, the predicted dark energy constraints from Con-X X-ray data are comparable to those projected by e.g. future supernovae, weak lensing and baryon oscillation experiments. Only by combining such independent and complementary methods can a rigorous and precise understanding of the nature of dark energy be achieved. ",
        "conclusions": ""
    },
    "0608/hep-th0608186_arXiv.txt": {
        "abstract": "String gas cosmology is rewritten in the Einstein frame. In an effective theory in which a gas of closed strings is coupled to a dilaton gravity background without any potential for the dilaton, the Hagedorn phase which is quasi-static in the string frame corresponds to an expanding, non-accelerating phase from the point of view of the Einstein frame. The Einstein frame curvature singularity which appears in this toy model is related to the blowing up of the dilaton in the string frame. However, for large values of the dilaton, the toy model clearly is inapplicable. Thus, there must be a new string phase which is likely to be static with frozen dilaton. With such a phase, the horizon problem can be successfully addressed in string gas cosmology. The generation of cosmological perturbations in the Hagedorn phase seeded by a gas of long strings in thermal equilibrium is reconsidered, both from the point of view of the string frame (in which it is easier to understand the generation of fluctuations) and the Einstein frame (in which the evolution equations are well known). It is shown that fixing the dilaton at some early stage is important in order to obtain a scale-invariant spectrum of cosmological fluctuations in string gas cosmology. ",
        "introduction": "String gas cosmology is a model of superstring cosmology which is based on coupling to a classical dilaton gravity background a gas of classical strings with a mass spectrum corresponding to one of the consistent perturbative superstring theories \\cite{BV,TV} (see also \\cite{Perlt} for early work, and \\cite{RHBrev1,RHBrev2,BattWat} for reviews). String gas cosmology has been developed in some detail in recent years \\cite{ABE,BEK,others}. In particular, it was shown \\cite{Watson,Patil1,Patil2,Edna,Watson2,Sugumi,other2} that string modes which become massless at enhanced symmetry points lead to a stabilization of the volume and shape moduli of the six extra spatial dimensions (see \\cite{BV,Mairi} for arguments in the context of string gas cosmology on how to naturally obtain the separation between three large and six string-scale dimensions \\footnote{See, however, \\cite{Col,Danos} for some caveats.}) .  String gas cosmology is usually formulated in the string frame, the frame in which stringy matter couples canonically to the background dilaton space-time. The existence of a maximal temperature \\cite{Hagedorn} of a gas of weakly interacting strings in thermal equilibrium has crucial consequences for string cosmology. As discussed in \\cite{BV}, as we follow our universe back in time through the radiation phase of standard cosmology, then when the temperature approaches its limiting value, the energy shifts from the radiative modes to the string oscillatory and winding modes. Thus, the pressure approaches zero. In the string frame, and for zero pressure, as follows from the equations derived in \\cite{TV,Ven}, the universe is quasi-static. There is an attractive fixed point of the dynamics in which the scale factor of our large three dimensions is constant, but the dilaton is dynamical (we consider the branch of solutions in which the dilaton is a decreasing function of time). We call this phase the quasi-static Hagedorn phase. Since the string frame Hubble radius is extremely large in the Hagedorn phase (infinite in the limiting case that the scale factor is exactly constant), but decreases dramatically during the transition to the radiation phase of standard cosmology, all comoving scales of interest in cosmology today are sub-Hubble initially, propagate on super-Hubble scales for a long time after the transition to the radiation phase before re-entering the Hubble radius at late times. Thus, it appears in principle possible to imagine a structure formation mechanism driven by local physics.  As was recently suggested \\cite{ABV}, string thermodynamic fluctuations in the Hagedorn phase may lead to a nearly scale-invariant spectrum of cosmological perturbations. There would be a slight red tilt for the spectrum of scalar metric perturbations. A key signature of this scenario would be a slight blue tilt for the spectrum of gravitational waves \\cite{BNPV} (see also \\cite{BNPV2,Ali} for more detailed treatments). Since this result is surprising from the point of view of particle cosmology, and since the analyses of \\cite{ABV,BNPV} contained approximations, it is an interesting challenge to analyze the cosmology from the point of view of the Einstein frame, the frame in which cosmologists have a better physical intuition.   In this paper, we analyze both the background dynamics of string gas cosmology and the generation and evolution of cosmological perturbations in the Einstein frame. A naive extrapolation of the background solutions of \\cite{TV} into the past would yield a cosmological singularity. Such an extrapolation is clearly not justified once the dilaton reaches values for which we enter the strong coupling regime of string theory. Rather, at early times one must have a new phase of the theory in which the dynamics    is consistent with the qualitative picture which emerges from string thermodynamics. This phase is meta-stable and will decay into a phase with rolling dilaton. The modified background evolution can solve the horizon and singularity problems in the context of string gas cosmology.  An improved analysis of generation of fluctuations in the string frame shows that the conclusions of \\cite{ABV,BNPV} concerning the spectra of cosmological perturbations and gravitational waves (we are considering the case where our three large dimensions are toroidal) are only obtained if the dilaton velocity can be neglected. ",
        "conclusions": "In this paper we have recast string gas cosmology in the Einstein frame rather than in the string frame in which the analysis usually takes place. Our analysis sheds new light on several important cosmological issues.  At the level of the background evolution, it becomes clear that solutions of the dilaton gravity equations (\\ref{E1} - \\ref{E3}) with decreasing dilaton contain an initial singularity. From the point of view of the Einstein frame, there is an initial curvature singularity which follows from the presence of a singularity in the dilaton field in the string frame. Obviously, however, these solutions are not applicable at very early times since they correspond to times when the string theory is strongly coupled. Hence, there must be, prior to the phase of rolling dilaton, a strong coupling Hagedorn phase in which both the size of space and the dilaton must be quasi-static. Provided that this phase lasts sufficiently long, thermal equilibrium over all scales relevant to current observations can be established.  Note that during the period when the solutions of (\\ref{E1} - \\ref{E3}) have a rolling dilaton, then in the Einstein frame the expansion of space never accelerates. The evolution corresponds to that of a radiation-dominated universe. The Einstein frame Hubble radius increases linearly in time throughout. However, the presence of the strong coupling Hagedorn phase can solve the horizon problem in the sense of making the comoving horizon larger than the comoving scale corresponding to our currently observed universe.  How to model the strong coupling Hagedorn phase now becomes a crucial question for string gas cosmology \\footnote{See also \\cite{Riotto} for an interesting discussion of these issues.}. Since the singularity of string gas cosmology is (from the point of view of the string frame) associated with the dilaton becoming large, and thus with string theory entering a strongly coupled phase, it is interesting to conjecture that a process like tachyon condensation \\cite{tachyon} will occur       and  resolve the singularity (like in the work of \\cite{Eva}). If this phase lasts for a long time, it will produce a large space in thermal equilibrium.  There are other possible scenarios in which there is a precursor phase of the rolling dilaton period which establishes thermal equilibrium on large scales. One such possibility was discussed in \\cite{Natalia} and makes use of a pre-Hagedorn phase in which the extra spatial dimensions initially expand, driven by a gas of bulk branes. The resulting increase in the energy stored in the branes leads to the increase in size and entropy which solves both the horizon and entropy problems. Once the extra spatial dimensions have contracted again to the string scale, the size of our three spatial dimensions can be macroscopic while the temperature of matter is of string scale.  Another possibility to obtain a solution to the horizon problem and to justify thermal equilibrium over large scales is to invoke a bouncing cosmology such as obtained in the context of higher derivative gravity models in \\cite{Biswas} (see also \\cite{MBS} for an earlier construction). The phase of contraction could produce the high densities required to form a string gas with the necessary requirements.  We have also studied the mechanism for the generation of fluctuations proposed in \\cite{ABV,BNPV2} in more detail, both from the point of view of the string frame and the Einstein frame. We have shown that a small value of the dilaton velocity in the Hagedorn phase is required in order that the string thermodynamic fluctuations are able to generate a scale-invariant spectrum of cosmological fluctuations. If dilaton velocity terms are important, then a Poisson spectrum is produced. If the dilaton velocity is negligible, and if thermal equilibrium on the scales of interest can be justified, then a scale-invariant spectrum of metric fluctuations is induced.  \\centerline{\\bf Acknowledgements}  Two of us (R.B. and A.N.) are grateful to Lev Kofman, Andrei Linde and Slava Mukhanov for detailed discussions, and for emphasizing the need to analyze the dynamics of string gas cosmology completely in the Einstein frame. R.B. thanks Misao Sasaki and the Yukawa Institute for Theoretical Physics for hospitality during the time when this project was initiated. We wish to thank Misao Sasaki, Scott Watson and in particular Cumrun Vafa for fruitful conversations. J.S. would like to thank the Department of Physics, McGill University for kind hospitality during the course of this work. The research of R.B. and S.P. is supported by an NSERC Discovery Grant, by the Canada Research Chairs program, and by an FQRNT Team Grant. The work of D.E. is supported in part by PPARC. S.K. is supported by a JSPS Postdoctoral Fellowship for Research Abroad. The research of J.K. at Perimeter Institute is supported in part by the Government of Canada through NSERC and by the Province of Ontario through MEDT. The work of A.N. is supported in part by NSF grant PHY-0244821 and DMS-0244464. The work of J.S. is supported by the Japan-U.K. Research Cooperative Program, the Japan-France Research Cooperative Program, and the Grant-in-Aid for  Scientific Research Fund of the Ministry of Education, Science and Culture of Japan No.18540262 and No.17340075."
    },
    "0608/astro-ph0608572.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "Our analysis showed that the proper motions of UCAC2 stars are distorted by a complex magnitude equation. Thus, for example, no noticeable effect of the magnitude equation was found in the proper motions of stars brighter than $\\approx11^m$, but this effect is significant in the magnitude range $12-15^m$. We showed that the magnitude equation has the greatest effect in $\\mu_\\alpha\\cos\\delta$ of northern-sky stars. This is because different sources for the southern and northern skies were used as the first epoch in deriving the proper motions of UCAC2 stars. We determined the coefficient of the magnitude equation for northern sky stars fainter than $13^m$ by comparing the proper motions of UCAC2 and PUL3SE stars, $-0.60\\pm0.05$ mas yr$^{-1}$ mag$^{-1}$. We associate this magnitude equation with the effect of the NPM catalog, Yellow Sky plates. The derived magnitude equation mainly affects the determination of the rotation tensor components. Its elimination leads to a $25\\%$ change in the rotation around the z axis (the Oort constant B) and to a $100\\%$ change in the rotation around the $y$ axis for $13^m$ stars.  The method of statistical parallaxes was used to estimate the distances to stars with random errors no larger than $14\\%$. The linear solar velocity relative to the local standard of rest, which is well determined for the local centroid ($d\\approx150$ pc), was used as a reference. This could affect the distance estimates for distant stars.  The method with fields was shown to reduce the correlations between the unknowns ($M_{\\scriptscriptstyle13}^{\\scriptscriptstyle-}$)--($M_{\\scriptscriptstyle13}^{\\scriptscriptstyle+}$) and ($M_{\\scriptscriptstyle32}^{\\scriptscriptstyle-}$)--($M_{\\scriptscriptstyle23}^{\\scriptscriptstyle+}$), being determined by an order of magnitude. When Eqs. (1) and (2) are solved simultaneously, the correlations between these unknowns are lower (than those for the solution of only Eq. (2)); therefore, the simultaneous solution is preferred.  The mean Oort constants estimated from Tycho-2 stars ($8^m.5-11^m.25$) are $A = 13.88\\pm0.98$ km s$^{-1}$ kpc$^{-1}$ and $B =-12.13\\pm0.63$ km s$^{-1}$ kpc$^{-1}$. The analogous means that we estimated from bright ($8^m-11^m.17$) UCAC2 stars have no statistically significant differences. The mean Oort constants estimated from UCAC2 stars fainter than $16^m.34$ are $A= 12.10\\pm0.58$ km s$^{-1}$ kpc$^{-1}$ and $B=-10.63\\pm0.46$ km s$^{-1}$ kpc$^{-1}$.  We established that the two components of the Ogorodnikov-Milne model that describe the rotation around the Galactic y axis and the deformation in the xz plane depend on heliocentric distance. For the nearest stars ($d\\approx200$~pc), these parameters are $-1.0\\pm0.1$~mas yr$^{-1}$ and decrease, on average, by a factor of 2 at distances of $\\sim$600 \u00ef\u00ea. We associate this with kinematic peculiarities of the Local system of stars. For more distant stars, there are, on average, no statistically significant deformation in the xz plane.  For distant Tycho-2/UCAC2 stars at heliocentric distances of $\\approx$900~pc, the mean rotation around the Galactic y axis is $-0.37\\pm0.04$~mas yr$^{-1}$, which we interpret as a residual rotation of the ICRS/Tycho-2 system relative to the inertial reference frame."
    },
    "0608/gr-qc0608019.txt": {
        "abstract": "We study cosmological self-reproduction in models of inflation driven by a scalar field $\\phi$ with a noncanonical kinetic term ($k$-inflation). We develop a general criterion for the existence of attractors and establish conditions selecting a class of $k$-inflation models that admit a unique attractor solution. We then consider quantum fluctuations on the attractor background. We show that the correlation length of the fluctuations is of order $c_{s}H^{-1}$, where $c_{s}$ is the speed of sound. By computing the magnitude of field fluctuations, we determine the coefficients of Fokker-Planck equations describing the probability distribution of the spatially averaged field $\\phi$. The field fluctuations are generally large in the inflationary attractor regime; hence, eternal self-reproduction is a generic feature of $k$-inflation. This is established more formally by demonstrating the existence of stationary solutions of the relevant FP equations. We also show that there exists a (model-dependent) range $\\phi_{R}<\\phi<\\phi_{\\max}$ within which large fluctuations are likely to drive the field towards the upper boundary $\\phi=\\phi_{\\max}$, where the semiclassical consideration breaks down. An exit from inflation into reheating without reaching $\\phi_{\\max}$ will occur almost surely (with probability 1) only if the initial value of $\\phi$ is below $\\phi_{R}$. In this way, strong self-reproduction effects constrain models of $k$-inflation. ",
        "introduction": "The paradigm of $k$-inflation~\\cite{Armendariz-Picon:1999rj} assumes that the energy density in the early universe is dominated by a scalar field $\\phi$ with an effective Lagrangian $p(X,\\phi)$ consisting solely of a (noncanonical) kinetic term, \\begin{equation} p(X,\\phi)=K(\\phi)L(X),\\quad X\\equiv\\frac{1}{2}\\partial_{\\mu}\\phi\\partial^{\\mu}\\phi,\\label{eq:Kessence L}\\end{equation} where $K(\\phi)$ and $L(X)$ are functions determined by the underlying fundamental theory. The effective Lagrangian~(\\ref{eq:Kessence L}) is assumed to be valid in a sufficiently wide range of values of $\\phi$ and $\\partial_{\\mu}\\phi$, as long as the energy density of the field $\\phi$ is below a suitable scale (e.g.~the Planck scale). Unlike the {}``traditional'' models of inflation, where the dominant energy density is due to a potential $V(\\phi)$, an accelerated expansion in $k$-inflation is driven by the kinetic energy of the scalar field.% \\footnote{An effective scalar field $\\phi$ with noncanonical kinetic terms is also used in models of $k$-\\emph{essence}~(\\cite{Armendariz-Picon:2000dh,Armendariz-Picon:2000ah}, and more recently e.g.~\\cite{Silverstein:2003hf,Alishahiha:2004eh}, where the Lagrangian $p(X,\\phi)$ is arranged to track the equation of state of other dominant matter components and to become dominant only at late times. In the present paper, we do not consider such models of $k$-essence but assume that the energy density of the field $\\phi$ is dominant during inflation. % } The form~(\\ref{eq:Kessence L}) of the Lagrangian is sufficiently general to cover many interesting cases. For instance, Lagrangians of the form\\begin{equation} p(X,\\phi)=K_{1}(\\phi)X^{n_{1}}+K_{2}(\\phi)X^{n_{2}},\\end{equation} where $n_{1},n_{2}$ are fixed numbers, are reduced to the Lagrangian~(\\ref{eq:Kessence L}) after a suitable redefinition of the field $\\phi$.  The assumed initial conditions for $k$-inflation consist of the field $\\phi$ having small spatial gradients within a sufficiently large initial region, and values of $\\phi$ and the time derivative $\\dot{\\phi}$ within suitable (and wide) ranges. These initial conditions guarantee that the energy density due to the field $\\phi$ is positive and dominates other forms of energy in the universe. The evolving homogeneous field $\\phi(t)$ rapidly approaches an attractor solution, $\\phi_{*}(t)$, which drives inflation. The presence of the attractor makes a model largely insensitive to initial conditions, since the attractor solution is unique and is approached exponentially quickly. In the present paper we develop a general definition of attractor solutions and demonstrate their existence in a wide range of models of $k$-inflation.  The back-reaction of quantum fluctuations modifies this classical picture of cosmological evolution, adding random {}``jumps'' to the classical {}``drift'' along the attractor line. In generic models of potential-driven inflation, quantum fluctuations give rise to the phenomenon of eternal self-reproduction~\\cite{Vilenkin:1983xq,Starobinsky:1986fx,Linde:1986fd,Goncharov:1987ir,Linde:1993xx}. As soon as suitable initial conditions occur in one domain, ultimately an infinite volume of space will undergo a long period of inflation followed by reheating. In this way, the eternal character of the self-reproduction process significantly alleviates the problem of initial conditions for inflation.  The main focus of this paper is to investigate self-reproduction in generic models of $k$-inflation with Lagrangians of the form~(\\ref{eq:Kessence L}). We find that quantum {}``jumps'' generically dominate over the {}``drift,'' and that self-reproduction is generically present. Self-reproduction in inflation can be studied using the stochastic ({}``diffusion'') approach based on a Fokker-Planck (FP) equation (see Ref.~\\cite{Linde:1993xx} for a review). Although the FP approach depends on the choice of the spacelike foliation of spacetime ({}``gauge''), one can use the FP equations to compute a number of gauge-independent characteristics, such as the fractal dimension of the inflating domain~\\cite{Aryal:1987vn,Winitzki:2001np} and the probability of reaching the end-of-inflation point $\\phi=\\phi_{E}$ of the configuration space. In particular, the presence of eternal self-reproduction is a gauge-invariant statement~\\cite{Winitzki:2001np,Winitzki:2005ya}. We show that the fractal dimension of the inflating domain is generically close to 3, confirming a strongly self-reproducing behavior in a wide class of models of $k$-inflation.  We also investigate the {}``stationarity'' of the distribution of $\\phi$, in the sense of Refs.~\\cite{Linde:1993xx,Garcia-Bellido:1993wn,Garcia-Bellido:1994ci}. There exists a (model-dependent) boundary value $\\phi_{\\max}$ such that the stochastic formalism breaks down for $\\phi>\\phi_{\\max}$. Imposing a boundary condition at $\\phi=\\phi_{\\max}$, we find that the stationary volume-weighted distribution of the field $\\phi$ is concentrated near $\\phi_{\\max}$. In the terminology of Ref.~\\cite{Garcia-Bellido:1993wn}, this constitutes a {}``runaway diffusion'' behavior, meaning that the volume-averaged value of the field is driven to the imposed upper limit. Finally, we follow the evolution of the field $\\phi$ along a single comoving worldline, starting at a generic value $\\phi_{0}\\lesssim\\phi_{\\max}$. We find that the field will either reach the boundary $\\phi=\\phi_{\\max}$ or, with roughly equal probability, exit the inflationary regime through the end-of-inflation (reheating) point $\\phi=\\phi_{E}$. An exit into reheating becomes overwhelmingly probable only if $\\phi_{0}<\\phi_{R}$, where $\\phi_{R}\\ll\\phi_{\\max}$ is a model-dependent threshold value. Therefore, the semiclassical description of the inflationary evolution along a single comoving worldline will remain valid only if the initial value $\\phi_{0}$ is chosen within the interval $\\phi_{E}<\\phi_{0}<\\phi_{R}$. It is interesting to note that a similar conclusion was reached in a study of stochastic back-reaction effects in models of quintessence~\\cite{Martin:2004ba}, where the initial value of the inflaton and the quintessence was found to be constrained to certain intervals. We compare this situation with the behavior of models of potential-driven inflation, where such a threshold $\\phi_{R}\\ll\\phi_{\\max}$ is absent and the exit into reheating has probability $1$ for all $\\phi_{0}\\lesssim\\phi_{\\max}$, where $\\phi_{\\max}$ is the Planck boundary.  Thus, the present investigation confirms that generic models of $k$-inflation exhibit standard features of eternal inflation, and shows that strong effects of self-reproduction constrain the choice of initial conditions. ",
        "conclusions": ""
    },
    "0608/astro-ph0608052_arXiv.txt": {
        "abstract": "We present new Near-Infrared ($J$,$K$) magnitudes for 114~RR~Lyrae stars in the globular cluster \\omc (NGC~5139) which we combine with data from the literature to construct a sample of 180 RR Lyrae stars with $J$ and $K$ mean magnitudes on a common photometric system. This is presently the largest such sample in any stellar system.  We also present updated predictions for $J$,$K$-band Period-Luminosity relations for both fundamental and first-overtone RR Lyrae stars, based on synthetic horizontal branch models with metal abundance ranging from $Z$=0.0001 to $Z$=0.004. By adopting for the \\omc variables with measured metal abundances an $\\alpha$-element enhancement of a factor of 3 ($\\approx 0.5$ dex) with respect to iron we find a true distance modulus $\\mu_0 = 13.70\\pm0.06\\pm0.06$ (random and systematic errors respectively), corresponding to a distance $d=5.5\\pm0.03\\pm0.03$~Kpc. Our estimate is in excellent agreement with the distance inferred for the eclipsing binary OGLEGC-17, but differ significantly from the recent distance estimates based on cluster dynamics and on high amplitude $\\delta$ Scuti stars. ",
        "introduction": "\\label{introduction}  The Period-Luminosity ($PL$) relation for Classical Cepheids and, in the near-infrared, for RR Lyrae stars is one of the most important methods to estimate distances and to calibrate secondary distance indicators (Bono et al. 2001; Freedman et al. 2001; Saha et al. 2001; Cassisi et al. 2004, hereinafter C04; Catelan, Pritzl, \\& Smith 2004, hereinafter CPS04; Gieren et al. 2005). Current trigonometric parallaxes from Hipparcos (Feast \\& Catchpole 1997), HST (Benedict et al. 2002), and from ground based telescopes (Lane et al. 2000; Kervella et al. 2004) have provided a unique possibility to calibrate the zero-point of these relations. Unfortunately, the number of Galactic calibrators with accurate trigonometric parallaxes is very limited, and for RR Lyrae stars it is restricted to a single object, namely RR Lyr itself. Therefore, we still lack a sound knowledge of the slope of the different PL relations based on purely geometrical methods. To overcome this problem, the astronomical community devoted a paramount observational effort to measure the pulsation properties of these robust distance indicators. The ideal laboratories for these studies are the Globular Clusters (GCs) and the dwarf spheroidals (dSphs) in the Local Group as the RR Lyrae stars in a given system are located at largely the same distance.  It has been demonstrated by Longmore et al. (1990) that cluster RR Lyrae follow a well-defined $K$-band $PL$ relation but we still lack a comprehensive empirical investigation of its properties.  For this reason, during the last few years we undertook a detailed pulsation and evolutionary investigation of RR Lyrae stars (Bono et al. 2003 and references therein) and, at the same time, we also started a large observational project aimed at collecting new Near-Infrared (NIR) data for RR Lyrae stars in  Galactic (M92: Del Principe et al. 2005) and Large Magellanic Cloud clusters (Reticulum: Dall'Ora et al. 2004).  Although the theoretical and empirical scenarios have recently seen new predictions based on Synthetic Horizontal Branch (SHB) models (Cassisi et al. 2004); Catelan, Pritzl, \\& Smith 2004) and new data (Butler 2003; Borissova et al. 2004; Storm 2004), we still lack a detailed NIR photometric analysis of a stellar system which includes a large sample of fundamental (F, \\rrab) and first-overtone (FO, \\rrc) RR Lyrae stars (only a dozen Galactic GCs host more than 50 RR Lyrae, see, e.g., Clement et al. 2001) and which spans a broad range in metal abundances. At present, the largest NIR data set includes 74 RR Lyrae stars whose metal abundance, based on 23 objects, ranges from [Fe/H]$\\approx-2$ to [Fe/H]$\\approx-0.7$, as  observed by Borissova et al. (2004) in the inner regions of the Large Magellanic Cloud.  In this context, the GC \\omc is a perfect target, since it  includes more than 186 RR Lyrae (86 \\rrab, 100 \\rrc; Kaluzny et al. 2004) with a metal content ranging from [Fe/H]$\\sim-2$ to [Fe/H]$ \\sim-1$ (131 objects, Rey et al 2000; 74 objects Sollima et al. 2006). Moreover, the \\omc distance has been estimated using different distance indicators such as the First Overtone Blue Edge (FOBE, Caputo et al. 2002), the eclipsing binary OGLEGC-17 (Thompson et al.  2001; Kaluzny et al. 2002), the $V$-band metallicity relation of RR Lyrae stars (Rey et al. 2000; Catelan 2005), and a dynamical analysis by van de Ven et al. (2006) based on proper motion and radial velocities of a large sample of individual stars.  In this paper, we present new NIR photometry for a substantial number of the RR Lyrae stars in \\omc and we provide new distance estimates which we compare with determinations already present in the literature. ",
        "conclusions": "\\label{Distance_01}  In Table \\ref{tab.distances} we have listed recent empirical distance estimates to \\omc where we have corrected all the moduli for the same reddening and employed the same reddening law as we have used for the RR Lyrae stars. We note that our new near-IR results are in excellent agreement with the largely geometric distance to the eclipsing binary OGLEGC-17 from Thompson et al. (2001) and Kaluzny et al. (2002), and from the First Overtone Blue Edge method based on pulsation predictions from Caputo et al. (2002). Current distance determinations based on $J$ and $K-band$ measurements and on NIR PL($\\tHB$) relations are also in very good agreement with similar estimates based on NIR relations derived by CPS04. In particular, for $\\tau_{HB}=0.94$ and $f=3$ we found $13.70\\pm0.06$ vs $13.72\\pm0.06$ ($K-band$) and $13.71\\pm0.10$ vs $13.76\\pm0.10$ ($J-band$). Note that distance estimates based on CPS04 relations were determined using the same solar metal abundance and we did not include a few very metal-poor RR Lyrae, since CPS04 prediction range from Z=0.0005 to Z=0.006. This finding is also supported by most recent $M_V-$[Fe/H] relations. By adopting the mean V-magnitudes for RR Lyrae stars provided by Kaluzny et al. (2004) and by Butler et al. (1978) together with the calibration given by Bono et al. (2003) we obtain for RR Lyrae with individual metal abundances a distance value of $(m-M)_0 = 13.72\\pm0.11$. We obtain a similar distance modulus ($13.62\\pm 0.11$) using the same magnitudes and the recent calibration of the $M_V-$[Fe/H] relation provided by Catelan (2005). The slope of such a relation was estimated by Cacciari \\& Clementini (2003), while the zero-point is based on the trigonometric parallax of RR Lyr itself (Benedict et al. 2002; Bono et al. 2002). Note that the systematic error for distance estimates based on visual magnitudes is $\\sigma=0.08$, due to the significant contribution from the absorption correction.  On the other hand, we deviate by more than three $\\sigma$ from the distance evaluation based on High Amplitude $\\delta$ Scuti (HADS) stars ($\\mu_0=14.05\\pm0.02$, McNamara 2000) and from the modulus of $\\mu_0 = 13.41\\pm0.13$ from van de Ven et al. (2006) based on cluster star dynamics.  It is worth noting that current uncertainties on metal abundance and on NIR photometry can not account for such a discrepancy.  The NIR $PL$ relations for RR Lyrae stars provide an accurate approach for determining absolute distances to globular clusters and other stellar systems. We find for \\omc a distance modulus of $13.70\\pm0.06\\pm0.06$ where the error estimates are random and systematic errors, respectively. We note that for populous systems like globulars and dwarf galaxies the dominating source of error are the systematic sources. The accuracy of NIR photometry in crowded regions significantly improved during the last few years and the sample sizes will certainly benefit from the new large field of view NIR cameras available at 4m class telescopes (van der Bliek et al. 2004). New $J-$band light curve templates for both fundamental and first overtone RR Lyrae are also urgently required to improve the accuracy of their mean magnitudes in this band."
    },
    "0608/astro-ph0608578_arXiv.txt": {
        "abstract": "We have proposed in previous papers that the high-frequency pair QPOs observed in black-hole binaries with frequency ratio 3:2 are inertial-acoustic oscillations (nearly horizontal oscillations with no node in the vertical direction) or g-mode oscillations, which are resonantly excited on warped relativistic disks. The resonance occurs through horizontal motions. In this model the dimensionless spin parameter $a_*$ of the central sources can be estimated when their masses are known from other observations. This estimate is done for three sources (GRO J1665-40, XTE J1550-564, GRS 1915+105). For all of them we have $a_*\\leq 0.45$. ",
        "introduction": "High-frequency quasi-periodic oscillations (HF QPOs), whose frequencies are in the range of 100 to 450 Hz, have been observed in some black-hole binaries and black-hole candidates. One of characteristics of these  HF QPOs is that they often appear in a pair and their frequencies change little with time \\footnote { In the kHz QPOs of neutron-star X-ray binaries, the frequencies (and their ratio) of the pair oscillations change with time. This is a difinite difference between HF QPOs in black holes binaries and kHz QPOs in neutron star binaries. In our warp models, kHz QPOs of neutron stars are interpreted as disk oscillations with {\\it vertical resonance} (see Kato 2005b), or as the case in which warp has precession (Kato 2005a). }, keeping the frequency ratio close to 3:2. These sources are GRO J1665-40 (300, 450 Hz), XTE J1550-564 (92, 184, 276 Hz) and GRS 1915+105 (41, 67, 113, 168 Hz) (e.g., a review by McClintock and Remillard 2006). Importance of commensurability of pair QPO frequencies on understanding the disk structure in the innermost region was emphasized by Abramowicz and Klu{\\' z}niak (2001) and Klu{\\' z}niak and Abramowicz (2001).  It is known that  oscillations can be excited in a deformed disk by resonant processes. One well-known example is superhumps in tidally-deformed dwarf-novae disks (Whitehurst 1988; Hirose, Osaki 1990; Lubow 1991). Another example is a spiral pattern on ram-pressure deformed galactic disks (Tosa 1994; Kato and Tosa 1994). In black-hole X-ray binaries, similar types of resonant oscillations should occur when the disks are deformed. We think that one of the most probable deformations of disks in the innermost region is a warp. Based on this idea, we examined excitation of disk oscillations on warped disks (Kato 2003b, 2004b), and proposed a resonant excitation model of QPOs (Kato 2004a,b, 2005a,b; Klu{\\' z}niak et al. 2004).  In this warped-disk model, the high-frequency QPOs in black-hole binaries are g-mode oscillations or inertial-acoustic oscillations\\footnote { In this paper inertial-acoustic oscillations represent the fundamental p-mode oscillations, in which oscillations are nearly horizontal and horizontal velocity has no node in the vertical direction. In some recent papers by Kato, however, inertial-acoustic oscillations are treated together with g-mode oscillations, since in mathematical analyses of the present resonance problem both of them can be treated together in a pack without making distinction (see Kato 2004b). Hence, when we used the term of g-mode oscillations, inertial-acoustic oscillations were implicitly included there. This was misleading. Hence, in this paper we explicitly mention inertial-acoustic oscillations without including them in g-mode oscillations. Here, it is noted that we use the following terminology for disk oscillations. The modes in which horizontal velocity has no node in the vertical direction ($n=0$) is called inertial-acoustic oscillations (p-modes). The modes with $n=1$ are g-modes and corrugation modes (c-modes). The modes with $n\\geq 2$ are g-modes and vertical p-modes} or their combination, excited by a horizontal resonance. If this resonance model is correct, it gives a way to estimate the spin of the central source from observed pair frequencies of QPOs, if the mass of the source is observationally known.  Recently, Shafee et al. (2006) evaluated the spins of two black-hole sources (GRO J1655-40 and 4U 1543-47) whose masses are observationally known, by fitting their spectra with model spectra derived from current disk models. Since one of the sources (i.e., GRO J1655-40) which they adopted has 3:2 pair frequencies, we can independently estimate the spin of the source. The purpose of this paper is to present the frequency-spin relation based on the warped-disk model and to estimate spins of some black-hole X-ray binaries, including GRO J1655-40. ",
        "conclusions": ""
    },
    "0608/astro-ph0608264_arXiv.txt": {
        "abstract": "{We constructed hydrodynamical model atmospheres for mid M-type main-, as well as pre-main-sequence (PMS) objects. Despite the complex chemistry encountered in these cool atmospheres a reasonably accurate representation of the radiative transfer is possible, even in the context of time-dependent and three-dimensional models. The models provide detailed information about the morphology of M-type granulation and statistical properties of the convective surface flows.  In particular, we determined the efficiency of the convective energy transport, and the efficiency of mixing by convective overshoot. The convective transport efficiency was expressed in terms of an equivalent mixing-length parameter \\mlp\\ in the formulation of mixing-length theory (MLT) given by \\citet{Mihalas78}. \\mlp\\ amounts to values around $\\approx 2$ for matching the entropy of the deep, adiabatically stratified regions of the convective envelope, and lies between 2.5 and 3.0 for matching the thermal structure of the deep photosphere.  For current spectral analysis of PMS objects this implies that MLT models based on $\\mlp=2.0$ overestimate the effective temperature by 100\\pun{K} and surface gravities by 0.25\\pun{dex}. The average thermal structure of the formally convectively stable layers is little affected by convective overshoot and wave heating, i.e., stays close to radiative equilibrium conditions. Our models suggest that the rate of mixing by convective overshoot declines exponentially with geometrical distance to the Schwarzschild stability boundary. \\changed{It increases at given effective temperature with decreasing gravitational acceleration.} ",
        "introduction": "The increasing number of stars, brown dwarfs, and extrasolar planets of spectral class M or later discovered by infrared surveys and radial velocity searches has spawned a great deal of interest in the atmospheric physics of these objects.  Their atmospheres are substantially cooler than the solar atmosphere, allowing the formation of molecules, or even liquid and solid condensates. Convection is a ubiquitous phenomenon in these atmospheres shaping their thermal structure and the distribution of chemical species. Hydrodynamical simulations of solar and stellar granulation including a realistic description of radiative transfer have become an increasingly powerful and handy instrument for studying the influence of convective flows on the structure of late-type stellar atmospheres as well as on the formation of their spectra \\citep[e.g.,][]{Nordlund82,Steffen+al89,Chan+Sofia89,Nordlund+Dravins90,CBTMH91,Ludwig+al94,Gadun+Pikalov96, Steiner+al98,Stein+Nordlund98,Asplund+al99b,Voegler+Schuessler03,Robinson+al04}. Here we report on efforts to construct hydrodynamical model atmospheres for mid M-type objects. The spectral type just borders the temperature where the formation of condensates becomes important.  The motivation of this investigation was twofold: first, pre-main-sequence (PMS) evolutionary models of M-type stars and brown dwarfs based on mixing-length theory \\citep[MLT,][]{BoehmVitense58} to describe the convective energy transport depend sensitively on the poorly constrained mixing-length parameter\\footnote{The ratio between the mixing-length and local pressure scale height.}~\\mlp\\ \\citep{Baraffe+al02}. Our hydrodynamical models represent convection essentially from first principles, and are free of the uncertainties of MLT allowing to put the stellar models on a firmer footing. Second, the distribution of dust clouds in cool brown dwarfs depends on the efficiency of mixing of their atmospheres by convective overshoot \\citep{Ackerman+Marley01,Allard+al03,Helling+al04}. For other work on the modeling of dust cloud formation in very low mass stars and brown dwarfs, see e.g. \\citet{Cooper+al03} and \\citet{Tsuji05}, and references therein. While by construction MLT cannot describe convective overshoot it is naturally represented in our three-dimensional hydrodynamical models.  The present investigation is extending a previous study of an M-dwarf atmosphere by \\citet[][hereafter LAH]{Ludwig+al02} to PMS objects at lower surface gravity.  A preliminary account of the results was given in \\citet{Ludwig03}. We start in Sect.~\\ref{s:modelover} with an overview of the model construction, in particular related to approximations we adopted in the radiative transfer.  In Sect.~\\ref{s:morphology} we discuss the general morphology of the convective flows in the M-type atmospheres and present some statistical properties. In Sect.~\\ref{s:convection} we provide estimates of the efficiency of the convective energy transport in terms of an effective mixing-length parameter, and discuss consequences for the analysis of PMS M-type objects in the framework of present standard model atmospheres.  We continue in Sect.~\\ref{s:mixing} by characterizing the properties of the atmospheric mixing found in the hydrodynamical models, and conclude with final remarks in Sect.~\\ref{s:remarks}. In our investigation we take repeatedly recourse to the solar atmosphere as standard benchmark. ",
        "conclusions": ""
    },
    "0608/astro-ph0608444_arXiv.txt": {
        "abstract": "Using {\\it Spitzer} IRAC mid-infrared imaging from the Great Observatories Origins Deep Survey, we study $z_{850}$-dropout sources in the Hubble Ultra Deep Field. After carefully removing contaminating flux from foreground sources, we clearly detect two $z_{850}$-dropouts at $3.6\\mu$m and $4.5\\mu$m, while two others are marginally detected. The mid-infrared fluxes strongly support their interpretation as galaxies at $z\\approx7$, seen when the Universe was only 750~Myr old. The IRAC observations allow us for the first time  to constrain the rest-frame optical colors, stellar masses, and ages of the highest redshift galaxies. Fitting stellar population models to the spectral energy distributions, we find photometric redshifts in the range $6.7-7.4$, rest-frame colors $U-V=0.2-0.4$, $V$-band luminosities $L_V=0.6-3 \\times 10^{10}~L_\\sun$, stellar masses $1-10 \\times 10^{9}$~\\msun, stellar ages $50-200$~Myr, star formation rates up to $\\sim25$~\\msun~yr$^{-1}$, and low reddening $A_V<0.4$. Overall, the $z=7$ galaxies appear substantially less massive and evolved than Lyman break galaxies or Distant Red Galaxies at $z=2-3$, but fairly similar to recently identified systems at $z=5-6$. The stellar mass density inferred from our $z=7$ sample is $\\rho_* = 1.6^{+1.6}_{-0.8}\\times 10^{6}$~\\msun Mpc$^{-3}$ (to 0.3$L^*_{z=3}$), in apparent agreement with recent cosmological hydrodynamic simulations, but we note that incompleteness and sample variance may introduce larger uncertainties. The ages of the two most massive galaxies suggest they formed at $z\\gtrsim8$, during the era of cosmic reionization, but the star formation rate density derived from their stellar masses and ages is not nearly sufficient to reionize the universe. The simplest explanation for this deficiency is that lower-mass galaxies beyond our detection limit reionized the universe. ",
        "introduction": "Observations of massive galaxies at high redshift with the {\\it Hubble Space Telescope} and the {\\it Spitzer Space Telescope} are revolutionizing our knowledge of the early formation history of stars and galaxies. Blue star forming galaxies at $z=2-3$ with stellar masses $10^{10}-10^{11}$\\msun\\ are routinely identified from optical imaging (Steidel et al. 1996a,b,2004), and have been studied in detail \\citep[e.g.,][Shapley et al. 2001,2005]{PDF01}, while near-infrared surveys at $z=2-3$ have uncovered substantial numbers of redder, more evolved galaxies with larger stellar masses $>10^{11}$\\msun\\  \\citep{Fr03,Ya04,Da05}. Some of these red galaxies appear to have stellar ages $> 1.5$ Gyr, implying that they formed most of their stars before $z\\sim5$ \\citep{La05}, and suggesting that massive galaxies should exist well beyond these redshifts. Direct detection of such galaxies would place strong constraints on galaxy formation models (e.g., \\citealt{So01,Na05}).  Tantalizingly, the most recent surveys with the Advanced Camera for Surveys (ACS) and the Near-Infrared Camera and Multiobject Spectrograph (NICMOS) on the {\\it Hubble Space Telescope} have identified sources at $z=5-6.5$ with fairly evolved stars and stellar masses of $1-4 \\times 10^{10}$\\msun\\ \\citep[e.g.,][]{Ya05,Ey05,Do05} or perhaps more \\citep{Mo05}. Critical for these results was access to the rest-frame wavelengths longward of the Balmer/4000 \\AA\\ break offered by the InfraRed Array Camera (IRAC; \\citealt{Fa04}) on {\\it Spitzer}. Without mid-infrared photometry to very faint magnitudes the stellar ages and masses of $z\\gtrsim5$  galaxies are  poorly constrained.  In this Letter, we extend mass estimates to $z=7-8$ by analyzing the mid-infrared fluxes of 6 $z_{850}$-dropout candidates found in the Hubble Ultra Deep Field (UDF) by Bouwens et al. (2004, hereafter B04). These candidates were selected from exceptionally deep optical ACS (Beckwith et al. in prep) and near-infrared NICMOS imaging \\citep{Th05}, and when combined with the ultradeep IRAC data available there, offer us an ideal opportunity to verify their reality and to study their stellar populations. The stellar masses and ages of $z_{850}$-dropout galaxies would provide us with the first direct look at galaxy formation at $z\\gtrsim7$, building on the comprehensive $z\\sim6$ study (Bouwens et al. 2006b). Where necessary, we assume an $\\Omega_M=0.3, \\Omega_\\Lambda=0.7,$ cosmology with $H_0=70$~km~s$^{-1}$Mpc$^{-1}$. Magnitudes are expressed in the AB photometric system.  \\ifemulate\\figa\\fi ",
        "conclusions": "Using the GOODS dataset (Dickinson et al. in prep), we have estimated and analyzed the {\\it Spitzer}/IRAC mid-infrared fluxes of 4 $z_{850}$-dropouts candidates, which were identified in the UDF by B04 and remeasured more accurately in B06. The sources are rare, with a surface density of 0.7~arcmin$^{-2}$, and very faint, with observed magnitudes of $H_{160}=26-27$ and $3.6\\mu=25-27$, placing them well beyond the spectroscopic capabilities of current telescopes, but in reach of future facilities such as JWST and ALMA. IRAC directly confirms the reality of two sources, while two others are marginally detected. Modeling of the broadband fluxes strongly supports their interpretation as $z\\approx7$ galaxies with substantial stellar masses $1-10 \\times 10^9$ and ages $50-200$~Myr.   Using the redshift selection function for the $z_{850}$-dropout sample (B04), we obtain an effective volume of 9000 Mpc$^{3}$ to $0.3L^*_{z=3}$ and we infer a stellar mass density of $\\rho_* = 1.6^{+1.6}_{-0.8}\\times 10^{6}$~\\msun Mpc$^{-3}$. Comparing to the stellar mass density at lower redshifts, computed using similar techniques to similar luminosities, we find a continuing decrease where the density at $z\\approx7$ is 95\\%, 22\\% of that at $z=6,5$  \\citep{Ya06,Sta06}. Recent Smooth Particle Hydrodynamics (SPH) simulations in a $\\Lambda$CDM universe predict stellar mass- and number densities for massive ($>1.8 \\times 10^{9}$M$_\\sun$) galaxies of $0.9 \\times 10^{6}$~\\msun~Mpc$^{-3}$ and $2.5 \\times 10^{-4}$~Mpc$^{-3}$ (SPH G6 run, \\citealt{Na05}), remarkably close to the minimum estimates for our sample ($0.7 \\times 10^{6}$~\\msun Mpc$^{-3}$ and $2.2 \\times 10^{-4}$~Mpc$^{-3}$ to the same mass limit). However, we cannot exclude additional stellar mass residing in massive non-starforming or dust-enshrouded galaxies, which the $z_{850}$-dropout criteria would have missed.  The ages of the two most massive (IRAC-detected) $z_{850}$-dropouts suggest that the bulk of their stellar mass formed at even higher redshifts $z\\gtrsim8$, during the epoch of cosmic reionization \\citep{Sp06}. We can place a simple upper limit on the contribution of high-mass galaxies to reionization, by calculating the maximum SFR densities implied by the observed stellar masses and ages for these objects. Taking the maximum masses and dividing it by the minimum ages ($\\sim60$Myr), we infer a substantial SFR density 0.04~M$_\\sun$yr$^{-1}$~Mpc$^{-3}$. This is higher than calculated directly from the rest-frame UV luminosities at $z=6-7$ to the same limits \\citep{Bo06b}, but  still more than $3$ times too small to reionize the universe at the lowest probable redshift $z_{reion}=8.6$ \\citep{Sp06} for canonical assumptions \\citep[][see also Yan et al. 2006]{Ma99,Bo06b}.  The simplest explanation for the lack of ionizing photons provided by high-mass galaxies at $z=7$ is that low-mass galaxies beyond our current detection limits were primarily responsible for reionization \\citep{LB03,YW04,Bo06b}. Other possibilities are that massive galaxies are missing from current surveys and are unaccounted for in the models, or that the observed galaxies had a top heavy IMF, which would increase the ionizing efficiency per unit stellar mass. Obviously, incompleteness, sample variance, and large scale structure dominate the uncertainties in our results. Larger, very deep near-infrared surveys from the ground and from space will address this issue in more detail."
    },
    "0608/astro-ph0608358_arXiv.txt": {
        "abstract": "{\\emph{Context}. We investigate mid-infrared and X-ray properties of the dusty torus in unification scenarios for active galactic nuclei.   \\emph{Aims}. We use the relation between mid IR and hard X-ray luminosities to constrain AGN unification scenarios.  \\emph{Methods}. With VISIR at the VLT, we have obtained the currently highest angular resolution ($0 \\farcs 35$ FWHM) narrow-band mid infrared images of the nuclei of 8 nearby Seyfert galaxies. Combining these observations with X-ray data from the literature we study the correlation between their mid IR and hard X-ray luminosities.  \\emph{Results}. We find that the rest frame $12.3 \\, \\mu$m ($L_{\\mathrm{MIR}}$) and 2-10 keV ($L_{\\mathrm{X}}$) luminosities are correlated at a highly significant level. The best fit power-law to our data is $\\log{L_{\\mathrm{MIR}}} \\propto \\left(1.60 \\pm 0.22 \\right) \\log{L_{\\mathrm{X}}}$, showing a much smaller dispersion than earlier studies.  \\emph{Conclusions}. The similarity in the $\\log{L_{\\mathrm{MIR}}}$ / $\\log{L_{\\mathrm{X}}}$ ratio between Sy1s and Sy2s even using high angular resolution MIR data implies that the similarity is intrinsic to AGN and not caused by contamination from extra-nuclear emission. This supports clumpy torus models. The exponent of the correlation constrains the inner geometry of the torus. ",
        "introduction": "The unification model for active galactic nuclei (AGN) interprets the different appearance of Seyfert 1 (Sy1) and Seyfert 2 (Sy2) galaxies uniquely as the result of an orientation effect \\citep{anton93,barthel94,urry96}. The central engine is considered to be surrounded by an optically and geometrically thick molecular torus. Associated with this torus are large masses of dust that supposedly reprocess the X-ray and UV emission from the accretion disk and re-emit it in the mid infrared (MIR) regime \\citep{pier93}.  It is thus very attractive to search for correlations between IR continuum and hard X-ray emission in order to test the unification scenarios for AGN. A tight correlation between the $10.5 \\, \\mu$m continuum and the absorption-corrected 2-10 keV luminosities for 8 nearby Seyfert galaxies was reported by \\citet{krabbe01} using $1\\farcs 2$ resolution MIR data. More recently \\citet{lutz04} found a correlation between the rest frame $6 \\, \\mu$m luminosity and the absorption-corrected hard X-ray luminosity for a sample of $71$ AGN. This sample was comprised of objects for which $24 \\arcsec$ angular resolution ISOPHOT spectra and hard X-ray observations were available; in particular it does not contain Compton-thick objects. However, the authors reported two problems that the unification model faces when compared  with their observations. (I) The scatter of the relation is about an order of magnitude larger than expected from the results of \\citet{krabbe01}. (II) There is no significant difference between early type (Sy types 1 to 1.5) and late type (Sy types 1.8 to 2) objects in the average ratio of mid-infrared to hard X-ray emission, as would be expected from an optically and geometrically thick torus dominating the mid-IR AGN continuum.  Here we study this relation through rest wavelength $12.3 \\, \\mu$m photometry of AGN, improving upon \\citet{krabbe01} and \\citet{lutz04} by a factor of 3 and 80, respectively, in terms of angular resolution. The wavelength was chosen as, among the bands accessible from the ground, it best traces hot and warm dust.  Throughout this paper we assume $H_{0} = 72\\,$km/s, $\\Omega_{\\Lambda} = 0.73$ and $\\Omega_{\\mathrm{m}} = 0.27$ \\citep{spergel03}. ",
        "conclusions": "\\label{discussion}  We have presented new high-resolution MIR observations of 8 nearby Seyferts, obtained with VISIR at the VLT. We found that their rest-frame $12.3 \\,\\mu$m and absorption-corrected 2-10 keV luminosities are strongly correlated following $L_{\\mathrm{MIR}} \\propto L_{\\mathrm{X}}^{1.6}$. Early- and late-type Seyfert galaxies follow the same relation. At the low-luminosity end the correlation is not well defined as one of the two object's MIR radiation is very probably not dominated by a dusty torus.  Both \\citet{krabbe01} and \\citet{lutz04} found a correlation close to $L_{\\mathrm{MIR}} \\propto L_{\\mathrm{X}}^1$. The exponent different to the one we find is most probably due to the difference in angular resolution ($24 \\arcsec$ for ISOPHOT, $1 \\farcs 2$ for Krabbe et al.'s sample and $ \\sim 0 \\farcs 35$ for VISIR) and to the choice of the rest frame wavelength as the intervening medium is expected to be more optically thin at $12.3 \\, \\mu$m than at shorter wavelengths. In particular, using the corresponding ISOPHOT $6 \\; \\mu$m data for our sample also yields an exponent close to 1. In addition we find a dispersion similar to the one in \\citet{krabbe01}; the larger one reported by \\citet{lutz04} can be attributed to their relatively poor angular resolution and the fact that the emission at $6 \\, \\mu$m is slightly more dependent on inclination effects \\citep[e.g. see][]{hoenig06}. On the other hand we confirm the absence of a significant difference in the $L_{\\mathrm{MIR}}$ / $L_{\\mathrm{X}}$ ratio between early and late-type objects found by both other studies. At our high angular resolution this is clearly in disagreement with smooth torus models, in which a flux difference of one order of magnitude is expected for type 1 and type 2 AGN of the same bolometric luminosity \\citep[e.g.][]{pier93}. This cannot be accounted for by dust clumps situated in the ionisation cone \\citep{efstat95} because their contribution to the nuclear MIR flux should not exceed 50\\%  \\citep[see][]{galliano05}. Clumpy torus models \\citep[e.g.][]{nenkova02} do not  predict such an offset and seem to be compatible with our results.  \\citet{hoenig06} recently used 3D radiative transfer modelling to demonstrate that clumpy tori can appear as optically thin in the MIR with most of the radiation originating in the innermost part of the torus. In their model individual clouds are optically thick but their filling factor is very small. Consequently lines-of-sight toward the innermost clouds are usually not blocked by other clouds. This implies that $L_{\\mathrm{MIR}}$ is proportional to the surface area $A$ of the inner edge of the torus and the covering factor of clouds projected on $A$. Our finding that Sy1s and Sy2s follow the same $L_{\\mathrm{X}}$ -- $L_{\\mathrm{MIR}}$ relation is a strong indication that the torus appears optically thin at $12.3 \\, \\mu$m.  The correlation we find then requires a dependency of $A$ or the covering factor on $L_{\\mathrm{X}}$ and thus the bolometric luminosity $L_{\\mathrm{Bol}}$. The simplest explanation for this observed effect is flaring: If the torus' inner edge is defined by the dust sublimation radius $r_0$, we obtain $A = 2 \\pi r_0 \\cdot H(r_0)$ where $H(r)$ is the scale height of the cloud distribution. For a flaring disk $H(r)$ can be written as $H(r) \\propto r^{\\alpha}$ with $\\alpha$ being constant.  This assumption yields $A \\propto r_0^{1+\\alpha}$. The dust sublimation radius $r_0$ follows $r_0 \\propto L_{\\mathrm{Bol}}^{0.5} \\propto L_{\\mathrm{X}}^{0.5}$. This leads to $L_{\\mathrm{MIR}} \\propto L_{\\mathrm{X}}^{0.5 + 0.5\\alpha}$. Our observations yield $L_{MIR} \\propto L_{X}^{1.6}$ and thus $\\alpha \\approx 2$. Including \\object{NGC 4579} results in $\\alpha \\approx 1.5$.  Another possibility is that its scale height depends on the radial accretion rate $\\dot{M}$. The torus model presented by \\citet{beckert04} predicts $H(r)/r \\propto \\dot{M}^{\\, 0.5}$. For AGN with comparable radiative efficiencies $\\dot{M}$ is proportional to $L_{\\mathrm{X}}$. This yields $L_{\\mathrm{MIR}} \\propto A \\propto L_{\\mathrm{X}}^{1.5}$ which is very close to the correlation we find.  The high angular resolution of our MIR observations indicates that the similarity of the $L_{\\mathrm{MIR}}$ / $L_{\\mathrm{X}}$ ratio for Sy1s and Sy2s is intrinsic rather than caused by extra-nuclear emission. Although contributions of dust clumps situated in the ionisation cone cannot be ruled out they should not dominate the nuclear MIR flux. Assuming our sample is representative of the general AGN population our results thus support recent models of clumpy dust tori \\citep[e.g.][]{nenkova02,hoenig06} and, furthermore, indicate that the filling factor of clouds is small. The geometry of the torus can be constrained by the slope of the $L_{\\mathrm{X}}$ -- $L_{\\mathrm{MIR}}$ correlation. Clumpy tori with low filling factors might also account for the similarity of \\emph{Spitzer} spectra of type 1 and type 2 QSOs \\citep{sturm06}."
    },
    "0608/astro-ph0608672_arXiv.txt": {
        "abstract": "An ensemble cluster has been formed from a data set comprising a complete magnitude-limited sample of 680 giant galaxies ($M_{B}^{0}$ \\raisebox{-1ex}{$\\stackrel{\\textstyle <}{\\sim}$} $-19$) in 8 low-redshift clusters, normalised by the velocity dispersions and virial radii for the early-type cluster populations.  Distinct galaxy populations have been identified, including an infall population.  A majority (50--70\\% or greater) of the infall population are found to be in interacting or merging systems characterised by slow gravitational encounters. The observed enhancement of galaxy--galaxy encounters in the infall population compared to the field can be explained by gravitational shocking.  It is shown that disc galaxy mergers in the infall population integrated over the estimated lifetime of the cluster ($\\sim$ 10 Gyr) can readily account for the present cluster S0 population. ",
        "introduction": "Although the transformation of luminous cluster disc galaxy populations from mainly spirals to predominantly S0s between intermediate redshifts ($z \\sim 0.5$) and the present is well established (e.g. Dressler 1980; Dressler et al. 1997), there is as yet no agreed mechanism for this transformation. One of the earliest explanations proposed was galaxy--galaxy interactions and mergers (e.g. Lavery \\& Henry 1988; Lavery, Pierce \\& McClure 1992). This explanation is attractive both on account of the observed abundance of interacting/merging systems in intermediate redshift clusters (e.g. Dressler et al. 1994a, 1994b) and because the Hubble sequence is readily explainable as a sequence of decreasing merger damage (Schweizer 2000). Nevertheless this explanation was resisted because it was considered that slow galaxy--galaxy interactions and mergers are likely to be rare in virialised clusters (e.g. Ostriker 1980; Makino \\& Hut 1997; Ghigna et al. 1998).  Although many other explanations have been attempted, including ram-pressure stripping of the cold interstellar gas of spirals by the hot ionised intracluster medium (e.g. Quilis, Moore \\& Bower 2000), galaxy `harassment', the impulsive heating of a galaxy disc by high-speed encounters (e.g. Moore et al. 1996, 1999), and `strangulation', the stripping of an (hypothesised) hot halo gas from spirals (e.g.  Larson et al. 1980; Bekki, Couch \\& Shioya 2002), these explanations are not without difficulties (for a recent discussion, see Mihos 2004).  Alternatively, attempts have been made to remove the problem entirely from clusters by proposing preprocessing of disc galaxies to earlier type systems in groups in the field (e.g. Balogh et al. 2004).  However, as remarked by Mihos (2004), neglect of slow encounters has been premature. Numerical simulations which track galaxy halos in a rich cluster in an $\\Omega_{m} = 1$ cosmology show that between $z = 0.5$ and $z=0$, although no mergers occur in the central virialised region of the cluster, in the outskirts the merger rate is 5--9\\% (see Ghigna et al. 1998). Gnedin (2003) has similarly studied the interaction and merger rates in clusters under different cosmologies. Over the lifetime of a cluster, the distribution of encounter velocities shows a large tail, but a significant fraction of encounters, largely those in the cluster periphery or those occurring at higher redshift before the cluster has fully collapsed, have relatively low velocity.  For example, for an $\\Omega_{m} = 0.4$ $\\Omega_{\\lambda} = 0.6$ cosmology, some 42\\% of encounters have a velocity $\\le$ 550 km ${\\rm s}^{-1}$.  As Mihos (2004) has noted, galaxy clusters form not by accreting individual galaxies randomly from the field environment, but rather through the infall of less massive groups falling in along the filaments that make up the `cosmic web'. Such infalling groups provide sites with much lower velocity dispersions than that of the cluster, thus permitting strong, slow encounters more normally associated with the field. For low-redshift clusters, such infall continues at the present (e.g. Moss \\& Dickens 1977; Tully \\& Shaya 1984; Sodre et al. 1989; Colless \\& Dunn 1996; Biviano et al. 1997; Rines et al. 2003).  Besides the theoretical considerations outlined above, there is also accumulating observational evidence of the potential importance of galaxy mergers for morphological transformation of cluster disc galaxies, both at low and intermediate-redshift (e.g. Dressler et al. 1997; Moss, Whittle \\& Pesce 1998; Moss \\& Whittle 2000, 2005; Koopmann \\& Kenney 2004; Sato \\& Martin 2006).  In the present paper the infall population of nearby clusters is studied using data from an extensive morphological and star formation survey of 8 low-redshift clusters (Moss \\& Whittle 2000, 2005). These data are used to construct an ensemble cluster from which general properties of all the clusters can be deduced.  It is shown that a very high fraction ($\\sim$ 50--70\\%) of the infall population are either interacting or possible mergers.  This interaction/merger rate is too high to be accounted for simply on the basis of infalling groups typical of the field, but implies that the accretion process of the cluster enhances the galaxy--galaxy interaction/merger rate.  It is suggested that galaxy shocking (Struck 2005) can provide a mechanism to enhance the interaction/merger rate, and that galaxy--galaxy interactions and mergers can provide an explanation for the transformation of most luminous cluster spirals to S0s over the past $\\sim$ 10 Gyr.  The paper is organised as follows.  In section 2, the survey data are summarised and construction of an ensemble cluster is discussed.  In section 3, various sub-populations of the ensemble cluster are identified, including that of the infall population. Member galaxies of the infall population are considered in detail and an estimate of the interaction/merger rate for this population is given. Discussion is given in section 4. It is shown that there is an enhancement of galaxy mergers in the infall population in comparison to the field; an explanation (gravitational shocking) is proposed for this enhancement; and finally it is shown that it is entirely plausible that the majority of the current S0 cluster population have resulted from galaxy mergers in the past $\\sim$ 10 Gyr.  Conclusions are given in section 5. ",
        "conclusions": "Analysis of an ensemble cluster comprising a complete magnitude-limited sample of giant galaxies ($M_{B}^{0}$ \\raisebox{-1ex}{$\\stackrel{\\textstyle <}{\\sim}$} $-19$) in 8 low-redshift clusters has shown that disturbed galaxies in these clusters form a infall population.  It has further been shown that the disturbance of the stellar populations of these galaxies can readily be explained by slow galaxy--galaxy encounters in at least 50--70\\% of cases. The resulting enhancement of slow encounters in the cluster infall population can be attributed to gravitational shocking (Struck 2005) of infalling galaxy groups.  A simple estimate of the galaxy merger rate demonstrates that the cluster giant S0 population can be accounted for as the outcome of minor mergers over the past $\\sim$ 10 Gyr.  Notwithstanding this result, obviously not all cluster S0 galaxies need be products of merging galaxies in the infall population; some early-type cluster galaxies are expected to result from galaxy interactions in groups or the field (e.g. Zabludoff \\& Mulchaey 1998) and this may help explain the correlation between star formation and radial distance to several virial radii from the cluster centre (e.g. Lewis et al. 2002; Gomez et al. 2003; although for an alternative view, see Poggianti 2004). However, contrary to Balogh et al. (2004), it is not necessary to assume that the majority of cluster S0s have formed by such preprocessing in the field.  Indeed the fact that most galaxies in the infall population are spirals (see section \\ref{dgin} above), rules out such an explanation. Rather, the bulk of transformation of spirals to S0s takes place during, and as an inherent part of the process of virialisation of the infall population.  Moreover, at least for the cluster galaxies, other processes may play a significant role in the formation of cluster giant S0 galaxies.  For the Virgo cluster, the majority of spiral galaxies have their H$\\alpha$ disks truncated most likely due to intracluster medium-interstellar medium (ICM-ISM) stripping as a cause of the reduced star formation (Koopmann \\& Kenney 2004). A similar effect has been found in other low-redshift clusters (Thomas et al. 2006). Such ram-pressure stripping of gas is likely to contribute to the transformation of cluster spirals to S0s.  Nevertheless, as the present work has shown, it is plausible to suppose that for most cluster giant S0 galaxies their formation has involved slow galaxy--galaxy interactions and ensuing minor mergers."
    },
    "0608/astro-ph0608391_arXiv.txt": {
        "abstract": "Given the number of recently discovered galaxy-galaxy lens systems, we anticipate that a gravitationally lensed supernova will be observed within the next few years.  We explore the possibility that stars in the lens galaxy will produce observable microlensing fluctuations in lensed supernova light curves.  For typical parameters, we predict that $\\sim$70\\% of lensed SNe will show microlensing fluctuations $> 0.5$ mag, while $\\sim$25\\% will have fluctuations $> 1$ mag. Thus microlensing of lensed supernova will be both ubiquitous and observable.  Additionally, we show that microlensing fluctuations will complicate measurements of time delays from multiply imaged supernovae: time delays accurate to better than a few days will be difficult to obtain.  We also consider prospects for extracting the lens galaxy's stellar mass fraction and mass function from microlensing fluctuations via a new statistical measure, the time-weighted light curve derivative. ",
        "introduction": "Within the past year there has been an explosion in the number of known galaxy-galaxy strong lens systems due to the success of the Sloan Lens ACS (SLACS) survey \\citep{bolton04,SLACS1,SLACS2,SLACS3}. The lens galaxies consist of a subset of the Sloan Digital Sky Survey (SDSS) Luminous Red Galaxy (LRG) catalog \\citep{LRG} and were initially selected based on the presence of several ``rogue'' oxygen emission lines in the otherwise well characterized LRG spectra \\citep{bolton04}.  The hypothesis that these emission features originated from an unresolved background galaxy within the 3\\arcsec\\ SDSS fiber suggested a substantial probability of lensing.  Follow-up HST snapshots with the ACS camera revealed that many of the source galaxies are in fact lensed \\citep{bolton04,SLACS1}.  Oxygen emission lines trace star formation in both the local and distant universe \\citep{kenn,hogg}, so it seems likely that the SLACS sources are undergoing intense star formation.  Coupled with the sheer number of lenses discovered with SLACS ($\\sim$30 at present count), this opens up the exciting possibility of observing lensed supernovae (SNe) in the near future.  Lensed SNe offer a variety of interesting applications.  For example, \\citet{og1} showed that SN magnification distributions can be used to break the degeneracy between the density profile of the lens galaxy and the measured value of the Hubble constant. \\citet{cooraySN} investigated the possibility of using weakly lensed SNe to map out large-scale density fluctuations in the universe.  \\citet{kolbar} proposed the use of SN magnification distributions to eliminate the mass-sheet degeneracy in cluster mass reconstruction.  In this paper, we point out that strongly lensed SNe will also be \\emph{microlensed} by the stars in the lens galaxy, and demonstrate that the microlensing fluctuations will contain valuable information about the stellar population of the lens galaxy.  The first lensed SN to be discovered will probably be Type II, as opposed to the Type Ia SNe used to measure the expansion history of the universe \\citep{riess98,perl,tonry03,knop,riess04}. While SNe II appear to be somewhat less common than SNe Ia when averaged over galaxies of all Hubble types, SNe II have a higher rate among spiral and irregular galaxies \\citep{kenn,mannucci}.  The presence of [\\ion{O}{2}] emission lines in the SLACS sources suggests that they are predominantly gas-rich spirals and irregulars.  Since star forming galaxies typically host one SN II every 50--100 years \\citep{mannucci}, we expect a SN II to be observed in one of the $\\sim$30 SLACS source galaxies within the next several years.  Although there is scatter among the peak brightnesses of SN II's, the mean B-band luminosity is approximately $M_B \\sim -17.5$ \\citep{youngSN}.   This implies an apparent magnitude of a SN in the SLACS source population (which has a mean redshift of $z_s \\approx 0.53$) of $m \\sim 24.5$. Since there is an additional amplification $\\mu_0 \\sim 10$ due to the main lens galaxy, the apparent magnitude for multiply imaged SNe in the SLACS sample should be $m \\sim 22$, a value that is certainly within reach of existing telescopes with reasonable exposure times.  In the SLACS sample, discoveries of lensed SNe will come from monitoring known lens systems.  The complementary approach is to mine a catalog of observed SNe to find those that are lensed. This approach could be used with the Supernova Acceleration Probe \\citep[SNAP;][also http://snap.lbl.gov]{SNAP}, which is intended to observe $\\sim$2000--3000 SNe Ia per year down to 26--27th magnitude. \\citet{og2} forecast that SNAP will discover 2--20 lensed SNe Ia per year, with the uncertainties due mainly to the star formation rate at high redshift.  Although SNAP will focus on SNe Ia, the deep imaging will naturally reveal many SNe II as well, raising the total number of lensed SNe by a factor of $\\sim$2.  While there are many categories and subcategories of observed SNe \\citep[see][for a thorough review]{filippenko}, the physics behind SN explosions has been studied extensively and both light curves and spectra are reasonably well reproduced in simulations \\citep{eastman2,eastman}.  The progenitors of SNe II are thought to be massive \\citep[$\\ga 8$--$10\\,M_\\odot$;][]{leonard} stars that undergo core collapse when the amount of energy released in fusion reactions is no longer sufficient to counteract self-gravitational forces.  The resulting explosion creates an expanding envelope that can be modeled with the Expanding Photosphere Method \\citep{kirshner}.  Light curves of some SN II have been observed with time sampling approaching 1 day \\citep[e.g., SN1999em;][]{leonard}.  Light that reaches us from a distant SN will of course be affected by all intervening matter, such as dust and gas, leading to photometric and spectroscopic differences between the intrinsic and observed light.  Lensed SNe are unique in that they may also exhibit gravitationally induced variability due to intervening stars lying within the lens galaxy.  This phenomenon known as microlensing is familiar from quasar lensing, with the best known example being Q2237+030 \\citep{wozniak}.  The brightnesses of the lensed images fluctuate on time scales of months to years due to relative motion of the lens galaxy and source quasar.  For a detailed example of microlensing light curve analysis see \\citet{csk}, and for detailed treatments of quasar microlensing theory see \\citet{WT,SW,SWL}.  In a microlensed SN, by contrast, what changes with time is the \\emph{size} of the source.\\footnote{The situation is similar to nanolensing of gamma ray bursts as explored by \\citet{walker03}.} The effects of source size in microlensing have been studied \\citep[e.g.,][]{mortonson,dobler2}, with the general conclusion that the amplitude of microlensing fluctuations decreases with increasing source size.  Thus, we na\\\"ively expect that microlensing fluctuations may start out large and decrease in amplitude as the SN expands.  Another qualitative result is that large sources are not very sensitive to microlensing by small stars.  Therefore we might hope to probe the mass function of stars in the lens galaxy by comparing the microlensing at different times or source sizes.  The notion of using multiple source sizes to ``map out'' lens mass distributions on various scales has been explored for quasar lensing by \\citet{met} and \\citet{dobler}.  In particular, \\citeauthor{dobler} plotted curves of magnification versus source size and pointed out that multiwavelength observations of quasar lenses could measure these curves.  However, with quasar lenses it is possible to obtain only a few discrete points on the curves, say for X-ray, optical, and radio wavelengths, and there is no clear mapping between the observable quantity (wavelength) and source size.  With SNe, by contrast, there is a simple mapping between the observable quantity (time) and the source size via the photosphere expansion rate.  Furthermore, with good time sampling it would be possible to measure the entire magnification versus source size curve.  Of course, seeing microlensing in SN light curves is contingent upon having a SN expansion rate that is neither too large (which would ``wash out'' microlensing effects at early times) nor too small (which would mean that the SN remains effectively a point source for its observable lifetime).  In \\S 2 we examine the typical SN expansion rate in units that are meaningful for microlensing, and estimate the optical depth for microlensing by stars in both the lens and host galaxies.  In \\S 3 we describe our microlensing simulations, and in \\S 4 we show sample microlensed SNe light curves.  In \\S 5 we interpret the light curves and consider whether microlensing fluctutations could be used to extract information about the stellar population: the relative abundances of stars and smooth dark matter at the image positions, and the mass function of stars.  We also consider whether microlensing will affect time delay measurements in lensed SNe, which have been proposed as a tool for precise measurements of $H_0$ \\citep{og1}.  Finally, in \\S 6 we summarize our findings and discuss possibilities for future theoretical and observational advances. ",
        "conclusions": "In anticipation of the discovery of lensed SNe in the near future with SLACS and SNAP, we have undertaken a pilot study of the fluctuations in the observed light curves due to microlensing by stars in the lens galaxy.  Our main findings can be summarized as follows. \\begin{itemize}  \\item The microlensing optical depth and the angular size of the source SN virtually ensure that microlensing fluctuations will be present in lensed SN light curves.  \\item Microlensing can produce an overall brightness offset of $\\sim$1 mag, short time scale brightness fluctuations of $\\sim$0.1--0.2 mag, and longer time scale fluctuations of $\\sim$0.5--1.0 mag.  \\item We estimate that $\\sim$70\\% of lensed SN light curves will have microlensing fluctuations larger than 0.5 mag, and $\\sim$25\\% will have fluctuations larger than 1 mag.  The probability that microlensing effects will \\emph{not} be present in any given light curve is $\\la 5\\%$.  \\item Microlensing fluctuations can generate dramatic and easily observable changes to the intrinsic SN light curve.  These features will complicate efforts to use time delays to measure the Hubble constant to high precision.  Without detailed knowledge of the intrinsic structure of the light curve, time delays accurate to better than a few days will be difficult to obtain.  \\end{itemize}  We should say that if microlensing fluctuations are \\emph{not} observed in lensed SN light curves, there are two possible interpretations: either the ratio of stellar mass density to smooth dark matter mass density (at the image positions) is so low that the microlensing optical depth is near zero; or the SN photosphere undergoes an initial period of very rapid expansion ($v \\gg 15,000$ km/s at early times).  The microlensing fluctuations depend on the stellar mass fraction and mass function in the lens galaxy, although there are such large difference between different random star configurations that it may be challenging to extract the stellar population parameters.  We have introduced the time-weighted derivative (TWD) of the light curve as a quantity that may be useful for this analysis.  The TWD is relevant because of the correlation between fluctuation time scales, stellar Einstein radii, and the photosphere expansion velocity.  We find that TWD statistics may be most useful for determining the relative abundances of stars and dark matter at the image positions.  More generally, the amount of structure in the light curves is a sign that more effort should be exerted to figure out how to analyze lensed SN light curves to determine properties of the stellar population.  Once a lensed SN is typed via spectroscopy or the late-time shape of the photometric light curve, it may be possible to model the intrinsic light curve.  The efficacy of EPM implies that the source geometry at different epochs in the observations will be fairly well known, and so a light curve fitting method \\citep[such as that described by][]{csk} may ultimately be the best analysis.  The fitting method must be extended to efficiently incorporate \\emph{expanding} sources as opposed to \\emph{moving} sources; however the basic idea of finding a family of possible stellar configurations and using a Bayesian likelihood analysis to extract mass function parameters will remain the same."
    },
    "0608/astro-ph0608028_arXiv.txt": {
        "abstract": "We present the results of 3-dimensional simulations of the direct collapse to a black hole of a rotating, 60\\,M$_\\odot$, zero metallicity, population III star.  Because the structure of this star (angular momentum, density and temperature profiles) is similar to many collapsar gamma-ray burst progenitors, these calculations have implications beyond the fate of population III stars.  These simulations provide a first 3-dimensional look at a realistic collapsar progenitor, and the results are very different from any previous 2-dimensional calculations.  If the angular momentum of the progenitor is high, non-axisymmetric instabilities in the collapsing core cause spiral structures to form, and these structures shape later outflows.  These outflows are driven by the imbalance between viscous heating and inefficient neutrino cooling and ultimately develop into a $10^{52}$\\,erg explosion.  Without magnetic fields, this collapse will not produce relativistic jets, but the explosion is indeed a hypernova.  We conclude with a discussion of the implications of such calculations on the explosions, nucleosynthesis, neutrino flux and gravitational wave emission from the collapse of massive stars. ",
        "introduction": "In the last two decades, observations have brought us a wealth of information about gamma-ray bursts (GRBs), transforming these phenomena from enigmatic curiosities to well-constrained astrophysical objects.  We now know much about the immediate surroundings, host galaxies, redshift distribution, energetics, and beaming of these bursts \\citep[e.g.][]{Sol05}.  These observations provide a number of strong constraints on any theoretical proposal for the engine behind these explosions.  The extreme GRB energies ($>10^{51}$\\,erg) and rapid observed variability ($\\sim$\\,ms) led theorists to propose models that invoked accretion onto compact objects \\citep*[see][]{Nar92}, and very quickly a range of formation scenarios were proposed to produce these accretion disks \\citep*{Pop99,Fry99}.  One such model invoked the collapse of the core of a rotating massive star down to a black hole \\citep{Woo93}.  As the rest of the star falls onto the core, its angular momentum causes it to hang up in a disk, producing a rapidly-accreting ($\\gtrsim 0.1 {\\rm M_\\odot\\,s^{-1}}$) black hole.  It is the accretion of this stellar material that is believed to power the GRB either through neutrino/anti-neutrino annihilation or winding of magnetic fields. With the discovery of supernovae-like light-curves concurrent and cospatial with long-duration GRBs, this ``collapsar'' model became one of the leading engines behind this class of GRBs.  Most of the work studying the collapsar model has focused on achieving the conditions (collapse to black hole, sufficient angular momentum to form a disk, etc.) required to produce an accreting black hole engine. Very little work has focused on studying the collapse and explosion process itself.  \\citet{Mac99} made the first such calculations, using a 2-dimensional, axially-symmetric variant of Prometheus \\citep*{Arn89,Fry91}.  They used simplified physics---an equation of state for electron-degenerate matter \\citep*{Blin96} with a functional term to approximate nuclear statistical equilibrium and neutrino cooling mimicked by a single cooling term.  They immediately collapsed the entire iron core to a black hole and placed an inner absorbing boundary at a radius of 50\\,km to mimic that black hole.  They began with the helium cores of several model stars and then artifically set the angular momentum profile to ensure that the centrifugal force at collapse would be $\\leq 0.02$ times the gravitational force.  \\citet{Pro03} used the same equation of state and cooling routine in the magnetohydrodynamic code Zeus-2D to model this collapse in axial symmetry all the way down to $1.5 R_{\\rm S}$ (where $R_{\\rm S} = 2GM/c^2$ is the Schwarschild radius), including the effects of magnetic fields.  They used the same 25\\,M$_\\odot$ progenitor with its ad-hoc angular momentum profile from \\citet{Mac99}.  And, like \\citet{Mac99}, at the onset, they collapsed the iron core down to a black hole.  In this paper, we present the first 3-dimensional simulations of the collapsar, following the accretion of material onto the black hole down to the innermost stable circular orbit.  Our progenitor is a zero-metallicity, rotating 60\\,M$_\\odot$ star and the rotation profile in our collapse is taken directly from the pre-collapse progenitor. We discover a number of features that the previous axially-symmetric models of collapsars did not, and could not, produce.  Most notably, we find that the initial accretion torus develops spiral wave instabilities, leading to enhanced angular momentum transport and viscous heating and the strong production of gravitational waves.  In \\S 2, we discuss our initial conditions and computational technique and the general features evident in our results.  We describe the transport of angular momentum through our simulated collapsar in \\S 3 and discuss the outflows from the central accreting torus in \\S 4.  We conclude in \\S 5 with a discussion of the implications of these 3-dimensional accretion instabilities, including analysis of the gravitational waves produced in these models. ",
        "conclusions": ""
    },
    "0608/astro-ph0608502_arXiv.txt": {
        "abstract": "We present an analysis of the observed broad iron line feature and putative warm absorber in the long 2001 {\\it XMM-Newton} observation of the Seyfert-1.2 galaxy MCG--6-30-15.  The new {\\tt kerrdisk} model we have designed for simulating line emission from accretion disk systems allows black hole spin to be a free parameter in the fit, enabling the user to formally constrain the angular momentum of a black hole, among other physical parameters of the system.  In an important extension of previous work, we derive constraints on the black hole spin in MCG--6-30-15 using a self-consistent model for X-ray reflection from the surface of the accretion disk while simultaneously accounting for absorption by dusty photoionized material along the line of sight (the warm absorber).  Even including these complications, the {\\it XMM-Newton}/EPIC-pn data require extreme relativistic broadening of the X-ray reflection spectrum; assuming no emission from within the radius of marginal stability, we derive a formal constraint on the dimensionless black hole spin parameter of $a=0.989^{+0.009}_{-0.002}$ at $90\\%$ confidence.  The principal unmodeled effect that can significantly reduce the inferred black hole spin is powerful emission from within the radius of marginal stability.  Although significant theoretical developments are required to fully understand this region, we argue that the need for a rapidly spinning black hole is robust to physically plausible levels of emission from within the radius of marginal stability.  In particular, we show that a non-rotating black hole is strongly ruled out. ",
        "introduction": "\\label{sec:intro}  Accreting black holes are the driving force behind some of the most powerful processes in the universe.  The central regions of active galactic nuclei (AGN), in particular, are prodigious sources across the electromagnetic spectrum including the X-ray band. The X-ray continuum (which is accurately approximated by a power-law up to $\\sim 100 \\keV$) is thought to be produced by Comptonization processes in the corona (or base of a jet) surrounding the inner part of the accretion disk.  A portion of the X-ray photons produced in the corona irradiate the underlying optically-thick disk, producing the so-called ``X-ray reflection'' signatures in the observed spectrum (Guilbert \\& Rees 1988; Lightman \\& White 1988).  These X-ray reflection signatures consist of fluorescent and recombination emission lines sitting on a continuum due to Compton scattering and the summed radiative recombination continua of the excited ions in the photoionized disk surface (George \\& Fabian 1991; Ross \\& Fabian 2005).  The Fe-K$\\alpha$ line is the most prominent of these features due to its energy (at $6.4 \\keV$ it is visible above the direct power law continuum), and the high astrophysical abundance and fluorescent yield of iron.  Additionally, this line is often significantly broadened by both the standard Doppler effect and relativistic processes, the effects of which increase the closer the line is emitted to the event horizon (Fabian \\etal 1989; Laor 1991).  These relativistic processes include light bending, beaming and gravitational redshifting, all of which result in a greatly elongated and skewed line profile; in particular, the line profile can display an extended low-energy tail primarily resulting from gravitational redshift.  The line emission region is sufficiently close to the black hole that frame-dragging effects associated with the black hole spin can be important in determining the line profiles.  The broad iron line is therefore a powerful probe of the relativistic effects on the spacetime immediately surrounding the black hole.  The first broad iron line robustly detected and resolved in an AGN was found by the {\\it Advanced Satellite for Cosmology and Astrophysics (ASCA)} in the Seyfert-1.2 galaxy MCG--6-30-15 (Tanaka \\etal 1995; Iwasawa \\etal 1996), and since then has been extensively studied with {\\it BeppoSAX} (Guainazzi \\etal 1999), {\\it RXTE} (Lee \\etal 1999, 2000), {\\it Chandra} (Lee \\etal 2002; Young \\etal 2005) and {\\it XMM-Newton} (Wilms \\etal 2001; Fabian \\etal 2002; Reynolds \\etal 2004).  All of these results show that the broad iron line feature is consistent with a highly redshifted line from the inner parts of an accretion disk; no alternative hypothesis has yet explained the spectrum of MCG--6-30-15 satisfactorily (Fabian \\etal 1995; Reynolds \\& Wilms 2000; Vaughan \\& Fabian 2004; Young \\etal 2005).  Subsequent {\\it ASCA}, {\\it Chandra} and {\\it XMM-Newton} studies have discovered similarly broad iron line profiles in several other Seyferts, such as MCG--5-23-16 (Dewangen, Griffiths \\& Schurch 2003), NGC 3516 (Turner \\etal 2002), Mrk335 (Gondoin \\etal 2002), and Mrk766 (Pounds \\etal 2003a).  In addition to having a sample of objects that have been observed with robust broad iron lines, it is equally important to have a precise model to use in fitting the data.  The two line profiles currently included as standard in the current versions of the commonly used spectral fitting package {\\sc xspec} (Arnaud 1996) are useful as a starting point, but ultimately quite limited in terms of their ability to accurately parameterize the line.  The {\\tt diskline} model (Fabian \\etal 1989) describes the line profile from a disk around a non-rotating (Schwarzschild) black hole, and, due to the approximations employed, does not include relativistic light bending. Similarly, the {\\tt laor} model (Laor 1991) has important constraints as well: this is a fully-relativistic model, but the dimensionless spin parameter of the black hole is hard-wired at $a=0.998$, the equilibrium spin of a black hole accreting from a Novikov \\& Thorne (1974) accretion disk. Furthermore, due to the computational realities of the early 1990s, the relativistic transfer functions underlying the {\\tt laor} model are pre-calculated and tabulated, yielding noise (or even gross inaccuracies) in the line profiles produced, especially at very high disk inclination angles.  Given these limitations, as well as the high quality of AGN spectra currently being obtained with {\\it Chandra} and {\\it XMM-Newton} (and hopefully in the future with the remaining operable instruments aboard {\\it Suzaku}), it is imperative that X-ray astronomers have access to effective models that are fully relativistic, accurate, and that allow black hole spin to be fit as a free parameter.  This latter point is crucial if we are to hope to constrain the spin of astrophysical black holes using broad iron lines and other X-ray reflection signatures.  Three new relativistic line models have recently been developed for this purpose and implemented in a form that can be readily used by X-ray astronomers: the {\\tt ky} suite (Dov\\v{c}iak \\etal 2004) and similar codes by Beckwith \\& Done (2004) and \\v{C}ade\\v{z} \\& Calvani (2005).  These models achieve comparable results for the morphologies of the line profiles, and all offer significant improvements over the {\\tt diskline} and {\\tt laor} results in terms of accuracy and precision over a wider range of physical parameters.  Most importantly, these models leave the spin of the black hole as a free parameter, include emission from within the innermost stable circular orbit in the disk, and compute fully relativistic photon transfer functions.  In practice, the {\\sc xspec} modules that implement both the {\\tt ky} and Beckwith \\& Done models use relativistic transfer functions stored in very large (multi-gigabyte) pre-calculated tables. Using these models, both groups conclude that fitting broad iron lines cannot truly constrain the spin of the black hole.  Since their models allow emission from {\\it any} radius of the accretion disk outside of the event horizon, they can produce iron line profiles with arbitrarily redshifted wings even if the underlying black hole spacetime has no spin.  One must consider the physicality of this assumption, however: it is not possible to get line emission from radii deep within the plunging region for several reasons.  Firstly, the closer one gets to the event horizon, the smaller the geometrical area of the disk available for emission becomes, while the gravitational redshift affecting any emission from such regions increases correspondingly.  Both of these effects would tend to minimize the contribution to the line profile from most of the emission inside the radius of marginal stability. Further, the ionization of the disk material within $\\sim 4 \\,r_{\\rm g}$ of a Schwarzchild black hole is inevitably too high to produce significant line emission in the first place, and the optical depth of this material also decreases rapidly within the radius of marginal stability (Reynolds \\& Begelman 1997; Young \\etal 1998). For all of these reasons, one simply cannot produce significant line emission from any arbitrary radius outside the event horizon.  This must be taken into account when modeling such line emission.  In this paper, we undertake a rigorous re-analysis of the broad iron line seen in the $\\sim 350 \\ks$ July-2001 {\\it XMM-Newton} observation (Fabian \\etal 2002).  We assess the constraints on the iron line profile and black hole spin including the complications introduced by other spectral components displayed by this system, especially the substantial column of absorbing photoionized gas seen along the line of sight to the central disk of this AGN.  To facilitate this investigation, we present a new variable-spin accretion disk line profile model, {\\tt kerrdisk}, that we are currently developing for public use in the {\\sc xspec} package.  The principal difference between {\\tt kerrdisk} and the previously mentioned variable-spin line profile models is that our model does not rely on gigabyte-sized, precomputed tables of the photon transfer function over many thousands of radii within the disk.  Rather, as we will discuss below, we use the approach of Cunningham (1975) to decompose the line profile integral into a simple analytic component (describing the effects of relativistic beaming) and an additional transfer function (describing the effects of light bending) which is a slowly varying function of black hole spin and disk inclination angle and hence only needs to be coarsely tabulated. This technique allows us to obtain high-quality line profiles with very modest sized (40\\,MB) tables, making it substantially more portable and convenient than some other models. This accuracy, portability and comparative speed of computation make {\\tt kerrdisk} a valuable addition to the X-ray astronomer's software arsenal.  Additionally, a trivial change of the driver code allows the user to implement any desired limb-darkening law or compute the line profiles to arbitrary resolution and accuracy (if, for example, the user is interested in the narrow cusps of the line profile).  Neither of these modifications is possible if the limb-darkening and resolution are hard-wired into the tabulated relativistic transfer functions.  In \\S2, we motivate our study by giving a brief recap of the history of the broad iron line in MCG--6-30-15.  \\S3 describes our new iron line model, including a detailed comparison of {\\tt kerrdisk} to the other line profile models already discussed. \\S4 presents our analysis of MCG--6-30-15 using these variable-spin line models.  This work is performed under the ``standard'' assumption that there is no observable X-ray reflection from inside the radius of marginal stability of the accretion disk.  A future paper will extend the analysis to the more general case where physically reasonable levels of emission from this region are allowed.  Under this assumption we arrive at the conclusion that the black hole in MCG--6-30-15 is a rapid-rotator.  \\S5 and \\S6 discuss our results and draw conclusions, respectively. ",
        "conclusions": "\\label{sec:conclusions}  Both Dov\\v{c}iak \\etal (2004) and Beckwith \\& Done (2004) make valid points about the difficulty of accurately calculating the spin parameter of a black hole based on broad line profiles from the disk alone.  Given the number of parameters governing the {\\tt kerrdisk} line profile (nine, to be exact), some degeneracy between them is certainly to be expected, and it is possible to generate statistically indistinguishable line profiles using different combinations of parameters.  This does not render the use of broad lines as spin diagnostics obsolete, however: provided that one uses care in determining whether parameter values are physically reasonable, it is possible to at least constrain the spin parameter to be within a certain range of values.  In the example of MCG--6-30-15 above, we have clearly shown that the data rule out a non-spinning black hole if we also demand that the values for the disk emissivity indices are physically realistic.  The plots of $\\Delta \\chi^2$ vs. $a$ clearly show the improvement in fit achieved when one frees the spin parameter.  In this case the fit strongly tends toward maximal spin.  The combination of such a high spin and such a high iron abundance in the disk may understandably give one pause when considering the precision of our best-fitting Model~5.  Both $a$ and $Z_{\\rm Fe}$ push the upper limits of their respective parameter spaces in order to try to account for the extreme breadth and strength of the Fe-K$\\alpha$ feature in the spectrum.  As we have mentioned, however, in this work we are seeking to establish a robust method for constraining black hole spin via X-ray spectroscopy and are not claiming that our best-fit parameter values should be interpreted as exact.  The spin of the black hole may indeed reach $a=0.989$, but there could also be some contribution to the line from highly redshifted radiation originating from within the ISCO that we have not yet accounted for, which could reduce the need for such a high spin value.  What we can say with confidence is that fits to the data with a Schwarzchild black hole are not physically sound.  Likewise, the disk may indeed have nearly ten times the solar value of iron, but underabundances in some of the lighter elements may also simulate an effectively iron-rich environment which could contribute to the observed width of the Fe-K$\\alpha$ line (Reynolds, Fabian \\& Inoue 1995).  {\\it Suzaku} observations of this source will be invaluable for untangling the parameters of the ionized reflection model.  Much work remains to be done on the subjects of modeling accretion disks and isolating black hole spins, and improvements in the former will undoubtedly help us place more accurate constraints on the latter.  If MCG--6-30-15 is indeed a rapidly spinning black hole, as seems likely given our spectral modeling, it is astrophysically interesting for several reasons.  As mentioned in \\S\\ref{sec:intro}, rapidly spinning holes can, in principle, experience a magnetic torque by the fields threading the accretion disk at the radius of marginal stability.  This torque can theoretically extract rotational energy from the hole itself, significantly enhancing the amount of dissipation in the inner accretion disk.  The steepest dissipation profiles would be obtained if the magnetic torque is applied completely at the radius of marginal stability (Agol \\& Krolik 2000). Therefore, only for rapidly spinning black holes would one expect to observe such a steep dissipation profile: in this scenario $r_{\\rm ms}$ is dragged inward very close to the event horizon, so the torque is strongest here.  Such an effect would manifest itself via strongly redshifted reflection features in the spectrum, since the strong dissipation very close to the event horizon would mean that most of the emission would originate from this region.  Based on its own best-fit emission profile, MCG--6-30-15 may in fact be giving us a glimpse of this phenomenon at work (Wilms \\etal 2001).  It should be noted, however, that interpreting the detected reflection features in this way demands that little of the observed emission originate from the plunging region within $r_{\\rm ms}$.  Taken in this context, the relatively steep emissivity index we found for the ionized disk in our best-fitting Model~5 ($a=0.989$, $\\alpha_1=6.06$) is not unexpected, and may be indicative of this type of magnetic torquing."
    },
    "0608/astro-ph0608516_arXiv.txt": {
        "abstract": "A model for gas outflows is proposed which simultaneously explains the correlations between the (i) equivalent width of low ionization and Ly$\\alpha$ lines, (ii) outflow velocity, and (iii) star formation rate observed in Lyman Break Galaxies (LBGs). Our interpretation implies that LBGs host short-lived ($30 \\pm 5$ Myr) starburst episodes observed at different evolutionary phases. Initially, the starburst powers a hot wind bound by a denser cold shell, which after $\\approx 5$ Myr becomes dynamically unstable and fragments; afterwards the fragment evolution is approximately ballistic while the hot bubble continues to expand. As the fragments are gravitationally decelerated, their screening ability of the starlight decreases as the UV starburst luminosity progressively dims. LBG observations sample all these evolutionary phases. Finally, the fragments fall back onto the galaxy after $\\approx 60$ Myr. This phase cannot be easily probed as it occurs when the starburst UV luminosity has already largely faded; however, galaxies dimmer in the UV than LBGs should show infalling gas. ",
        "introduction": "Recent observations (Adelberger \\etal 2003, Shapley \\etal 2003 (S03), Steidel \\etal 2003, Adelberger \\etal 2005) based on a combination of quasar absorption-line and faint-galaxy techniques applied to the same cosmic volumes have revealed that UV-selected, star forming galaxies at redshift $z\\approx 3$ (in brief,  Lyman Break Galaxies, LBGs) show spectral absorption lines due to heavy elements. These absorption lines, which can be very strong ($N_\\nCIV \\gg 10^{14}$~cm$^{-2}$),  are often called ``interstellar\" because they resemble the absorption lines produced by interstellar material in local galaxies. Nevertheless, these lines appear (without exception) to be  {\\it blueshifted} with respect to the galaxy systemic velocity deduced from nebular lines. The typical velocity difference is $\\Delta v \\approx 250$~km~s$^{-1}$, but a non-negligible fraction of LBGs show differences in excess of 300 ~km~s$^{-1}$ (Pettini \\etal 2002). These absorbing systems are often interpreted as arising in cool and/or neutral structures (maybe a shell) part of an outflow powered by supernovae associated with the detected galaxy star formation activity.  An additional element in support of this hypothesis comes from the detection of {\\it redshifted} Ly$\\alpha$ photons backscattered from gas behind the stars. Quite interestingly, this emission could trace the receding part of the outflowing shell. Finally, strong Ly$\\alpha$ absorption is produced by the intergalactic gas within $1 h^{-1}$ comoving Mpc of most LBGs, but for about 1/3 the absorption is weak or absent. The transparency of the medium around some LBGs could arise from the gas collisional ionization/heating by the wind (Croft \\etal 2002, Adelberger \\etal 2003,  Kollmeier \\etal 2003, Bruscoli \\etal 2003, Bertone 2005, Kollmeier \\etal 2006), from a galaxy proximity (photoionization) effect, as pointed out by Maselli \\etal (2004), or simply reflect cosmic variance and artifacts of the data reduction (Desjacques \\etal 2004, Desjacques, Haehnelt \\& Nusser 2005).  In general,  one would expect to observe also infalling gas in the surroundings of LBGs arising either from gas accretion along the cosmic web filaments or because at least part of the outflowing gas could rain back onto the galaxy if its velocity does not exceed the escape speed of the system (a process known as galactic fountain). In fact, the infall ram pressure might be of primary importance in confining and reverting  the wind expansion (Fujita \\etal 2004). This situation is clearly appreciated from the inspection of the results of recent, high resolution simulation of LBG outflows (Ferrara, Scannapieco \\& Bergeron; see Fig. 2 of that paper). However, observations so far have not  found any experimental evidence of such physically-based expectation.  The central aim of this paper is to understand why outflows are routinely observed in LBGs whereas gas infall seem to escape detection.  This puzzle is possibly related to the following scenario. During the first phase of a starburst the galactic wind produces a cavity filled with hot gas, enclosed by a cold and dense shell. Afterwards the shell fragments and the velocity of the cold debris starts decreasing while the hot gas escapes from the galaxy.  Whether or not they survive for some time, the UV light from the galaxy should become progressively less absorbed by the debris since their covering factor decreases with time as they travel to larger distances from the galaxy.  Observations may show a time evolution of the equivalent width (EW) of metal absorption lines and dust reddening produced by the decrease of the screening effect of the cold debris on the UV light from the galaxy. If the debris are not disrupted or ejected from the galaxy we should also observe the gas return (infall) process. The relevant question to pose then seems: Can we observe galaxies during all these different evolutionary stages?  A large statistical galactic sample (more than 800 galaxies at intermediate redshifts) that can be used to test this model has been compiled by S03. From the analysis of this data set a very interesting number of relationships emerge (some of which already noticed by S03 and others obtained by correlating their data in a different way), which are listed in the following: \\begin{itemize} \\item[(1)] The wind velocity, $v_w$, is proportional to the star formation rate (SFR): $v_w  \\propto \\dot M_\\star$ \\item[(2)] The Ly$\\alpha$ emission line EW is  inversely proportional to the wind velocity: ${\\rm EW}_{Ly\\alpha}\\propto 1/v_w$ \\item[(3)] The EW of metal absorption lines is proportional to the wind velocity: ${\\rm EW}_{met}\\propto v_w$ \\item[(4)] The reddening is proportional to the wind velocity: $E(B-V)\\propto v_w$ \\end{itemize}  Should we interpret these properties as produced by intrinsic variations of the physical properties ({\\it eg.}, SFR, mass, size) of LBGs?  Or do they rather reflect the different ages of a single (broadly speaking) galactic population with uniform properties? The relationships (3) and (4) naturally descend from the properties of a time evolving uniform population: as the cold debris move outward they are slowed down by gravity/energy losses, their covering factor decreases, and the EW and reddening also decrease. Property (2) can easily be understood in the same framework.  Contrary to the metal absorption lines, the EW of the Ly$\\alpha$ emission line increases with decreasing wind velocity ({\\it ie.}, increasing time from the starburst) since the scattering by the outflowing neutral gas is reduced and a larger fraction of the blue wing of the Ly$\\alpha$ line is transmitted.  Naively, relationship (1) can be interpreted as due to a variation of the energy input of the wind for galaxies with different star formation rate. However, a closer look at that relationship (\\fig~\\ref{fig:v_sfr}) shows that the dependence of $v_w$ on the SFR is not as the one expected under this assumption.  Assuming, for simplicity, a constant density profile for the gas in the galaxy core, the velocity is a very weak function of the SFR: $v \\propto \\dot M_\\star^q$ with $q=1/5$ (shown by the solid line).  Even assuming a more general gas density profile of the galaxy core, $\\rho_g \\propto r^{-\\beta}$, would produce $q =1/(5-\\beta)$ that is too flat to fit the observations for any realistic choice of $\\beta$.  A very steep slope $\\beta=3.5$ (shown by the dashed line) fits the observations, but would produce unrealistically large outflow velocities incompatible with the data. Moreover, we will show later that the cold shell becomes unstable when the slope of the density profile becomes steeper than $\\beta \\simeq 2$: at that point  fragments decouple from the hot wind. Thus, the steepest physical dependence of $v_w$ on the SFR is $q \\approx 1/3$  when $\\beta=2$, while the best fit to the data requires $q \\approx 2/3$.  Guided by the aforementioned arguments, in this paper we propose that the different observed wind velocities correspond to different starburst evolutionary times, and are not directly related to variations of their mean SFR or other internal properties. The proportionality of the wind velocity and SFR is incidental. It reflects the fact that both quantities decrease as the time elapsed from the starburst increases.  We have organized the paper as follows. In \\S~\\ref{sec:sca} we describe our model analytically assuming a power law density profile for the gas. In \\S~\\ref{sec:res} we assume a more realistic density profile of the gas and we compare the numerical results to the observations. In \\S~\\ref{sec:dis} we discuss the implications of our model and provide an explanation for the lack of evidences for infalling gas in LBGs.   \\begin{figure} \\centerline{\\psfig{figure=fig_wind.eps,height=9cm}} \\caption{Wind velocity as a function of the LBG star formation rate. The points with errorbars show the data from S03; the dashed line is the best fit to the data with a power law $v_{w}   \\propto \\dot M_\\star^{2/3}$.  A pure  luminosity effect would produce a much flatter power law:  $v_w  \\propto \\dot M_\\star^{1/5}$ (solid line).} \\label{fig:v_sfr} \\end{figure} ",
        "conclusions": ""
    },
    "0608/astro-ph0608385_arXiv.txt": {
        "abstract": "The radiation background produced by the 21 cm spin-flip transition of neutral hydrogen at high redshifts can be a pristine probe of fundamental physics and cosmology.  At $z \\sim 30$--$300$, the intergalactic medium (IGM) is visible in 21 cm absorption against the cosmic microwave background (CMB), with a strength that depends on the thermal (and ionization) history of the IGM.  Here we examine the constraints this background can place on dark matter decay and annihilation, which could heat and ionize the IGM through the production of high-energy particles.  Using a simple model for dark matter decay, we show that, if the decay energy is immediately injected into the IGM, the 21 cm background can detect energy injection rates $\\ga 10^{-24}$ eV cm$^{-3}$ sec$^{-1}$.  If all the dark matter is subject to decay, this allows us to constrain dark matter lifetimes $\\la 10^{27} \\sec$. Such energy injection rates  are much smaller than those typically probed by the CMB power spectra.  The expected brightness temperature fluctuations at $z\\sim 50$ are a fraction of a mK and can vary from the standard calculation by up to an order of magnitude, although the difference can be significantly smaller if some of the decay products free stream to lower redshifts.  For self-annihilating dark matter, the fluctuation amplitude can differ by a factor $\\la 2$ from the standard calculation at $z \\sim 50$.  Note also that, in contrast to the CMB, the 21 cm probe is sensitive to both the ionization fraction and the IGM temperature, in principle allowing better constraints on the decay process and heating history.  We also show that strong IGM heating and ionization can lead to an enhanced H$_2$ abundance, which may affect the earliest generations of stars and galaxies. ",
        "introduction": "Introduction}  The cosmic ``dark ages,\" stretching from the last scattering surface of the cosmic microwave background (CMB) at $z \\sim 1100$ to the formation of the first luminous sources at $z \\sim 30$, are one of the last frontiers for observational cosmology.  The problem is the difficulty of finding any probes:  because the intergalactic medium (IGM) has mostly decoupled from the CMB, there is no background radiation field against which we can study it, and (by definition) there are no local light sources.  This is unfortunate because, at least in principle, the physics of the dark ages is sufficiently simple that we can hope to understand it in detail.  The only factors that enter are the CMB, the expanding Universe, recombinations, Compton scattering (which couples the CMB to the IGM), and gravitational growth -- almost entirely in the linear regime.  Thus, the IGM properties during the dark ages (and especially its fluctuations) constitute a strong test of fundamental cosmological parameters in an entirely analogous way to the CMB.  Conversely, if we take the physics as well-understood, probes of the dark ages would offer stringent tests of exotic physics.  Two examples are dark matter decay and annihilation.  If either of these processes happen (even for only a fraction of the dark matter), they could inject high-energy photons (or other particles) into the IGM.  These would then scatter and deposit some or all of their energy as ionizations and heat.  An altered reionization history may affect the total optical depth and, as a consequence, the CMB power spectrum. Decaying dark matter particles were in fact  initially advocated  as a possible explanation for the high optical depth  detected in the first year of \\emph{WMAP} data \\cite{chen04-decay, hansen04, kasuya04, avelino04, pierpaoli04,Bean03}.  Such data allowed an upper limit on the energy injection  rate and consequently on the decay time ($t \\ge H_0^{-1}$).  (These constraints, and others mentioned throughout this paper, were obtained with the first-year \\emph{WMAP} data, but we do not expect the newer data \\cite{spergel06} to improve them substantially given the many uncertainties on the various reionization processes affecting the total optical depth.)  However, the CMB is only affected if a substantial fraction of the CMB photons interact with the IGM -- or in other words if the optical depth to electron scattering is substantially altered by the decay products \\cite{chen04-decay, pierpaoli04, mapelli06}.  Compared to the IGM temperatures of $\\sim 10$--$10^3 \\kel$ at $z \\sim 30$--$300$, ionization requires a substantial energy input, and decays with long timescales (or low energy injection rates) cannot be ruled out with the CMB.  Nevertheless, because the IGM temperature is so small, these models can still significantly affect the thermal history of the IGM \\cite {chen04-decay, mapelli06} and hence the history of star formation. A number of recent particle physics models motivate such scenarios (see the discussion below in \\S \\ref{particles}).  Dark matter annihilation has also been considered recently, primarily as an explanation for gamma-ray sources in the local Universe (e.g., \\cite{profumo06}).  But any such scenario will also predict an annihilation background from encounters between IGM particles; because the physical density increases like $(1+z)^3$, such events could be relatively common in the early Universe and so affect the CMB \\cite{padmanabhan05, mapelli06}.  Of course, these scenarios will continue to affect the IGM throughout the dark ages as well.  The 21 cm background can offer  powerful constraints on dark matter decay and annihilation (or indeed any exotic process that injects energy into the IGM during the dark ages).  The hyperfine level populations of the ground state of neutral hydrogen are determined through competition between absorption of (and emission stimulated by) CMB photons (and possibly UV photons) and collisions.  Since the IGM is dense at high-redshifts, collisions dominate, driving the level populations into equilibrium with the kinetic temperature of the gas.  Because the latter is colder than the CMB, the IGM is a net absorber of CMB photons at $z \\sim 30$--$300$.  Thus, the 21 cm transition can be used to map fluctuations in the IGM during the dark ages, offering a pristine probe of cosmology \\cite{loeb04}.  Previous work has emphasized the possibility of constraining the matter power spectrum with this tool \\cite{loeb04, bharadwaj04-vel, barkana05-infall, naoz05}.  Here we point out that the fundamental properties of the fluctuations -- on all scales -- depend sensitively on the thermal history (see also Ref.~\\onlinecite{shchekinov06}).  Thus the 21 cm background can be used to constrain  dark matter decay and annihilation.  It is much more sensitive than the CMB to dark matter decay because (i) it depends on the thermal history, not just the ionized fraction, and (ii) it is directly sensitive to the late time behavior (when most of the energy is injected).  It is less useful for constraining annihilation scenarios because in that case a large fraction of the energy is injected at early times.  The excess heating and ionization induced by dark matter decay and annihilation can also affect the chemistry of the IGM.  In the simple chemical environment of the primordial IGM, their most important effect will be on the H$_2$ abundance, which is significant because that molecule is an important coolant for low-temperature gas \\cite{saslaw67, galli98} and is thought to play a key role in the formation of the first stars \\cite{haiman96, tegmark97, abel02, bromm02}.  Unfortunately, the implications -- and even magnitude -- of any possible boost to early structure formation are controversial \\cite{padmanabhan05,biermann06,ripamonti06}.  Here, we provide a more detailed look at ${\\rm H_{2}}$ formation and incorporate some hitherto neglected effects in the calculation.  The rest of this paper is organized as follows.  We briefly review the physics of the 21 cm background in \\S \\ref{21cm}.  We describe our simple model for dark matter decay and annihilation, and review some particle physics motivation, in \\S \\ref{dm}.  We then present our results for the 21 cm background in \\S \\ref{results} and for the H$_2$ abundance in \\S \\ref{htwo}.  Finally, we discuss their implications in \\S \\ref{disc}.  In our numerical calculations, we assume a cosmology with $\\Omega_m=0.26$, $\\Omega_\\Lambda=0.74$, $\\Omega_b=0.044$, $H=100 h \\hunits$ (with $h=0.74$), $n=0.95$, and $\\sigma_8=0.8$, consistent with the most recent measurements \\citep{spergel06}, although we have increased $\\sigma_8$ from the best-fit \\emph{WMAP} value to improve agreement with weak lensing.  We quote all distances in comoving units, unless otherwise specified. ",
        "conclusions": ""
    },
    "0608/astro-ph0608666_arXiv.txt": {
        "abstract": "Theoretical models of very metal-poor intermediate-mass Asymptotic Giant Branch (AGB) stars predict a large overabundance of primary nitrogen. The very metal-poor, carbon-enhanced, s-process-rich stars, which are thought to be the polluted companions of now-extinct AGB stars, provide direct tests of the predictions of these models. Recent studies of the carbon and nitrogen abundances in metal-poor stars have focused on the most carbon-rich stars, leading to a potential selection bias against stars that have been polluted by AGB stars that produced large amounts of nitrogen, and hence have small [C/N] ratios. We call these stars Nitrogen-Enhanced Metal-Poor (NEMP) stars, and define them as having [N/Fe] $> +0.5$ and [C/N] $< -0.5$. In this paper, we report on the [C/N] abundances of a sample of 21 carbon-enhanced stars, all but three of which have ${\\rm [C/Fe]} < +2.0$. If NEMP stars were made as easily as Carbon-Enhanced Metal-Poor (CEMP) stars, then we expected to find between two and seven NEMP stars. Instead, we found no NEMP stars in our sample. Therefore, this observational bias is not an important contributor to the apparent dearth of N-rich stars. Our [C/N] values are in the same range as values reported previously in the literature ($-0.5$ to $+2.0$), and all stars are in disagreement with the predicted [C/N] ratios for both low-mass and high-mass AGB stars. We suggest that the decrease in [C/N] from the low-mass AGB models is due to enhanced extra-mixing, while the lack of NEMP stars may be caused by unfavorable mass ratios in binaries or the difficulty of mass transfer in binary systems with large mass ratios. ",
        "introduction": "\\label{Sect:Intro}  Very Metal-Poor (VMP) stars ($\\mem{[Fe/H]}\\le -2.0$)\\footnote{We adopt the usual spectroscopic notation that [A/B] $\\equiv$ {\\rm log}$_{\\rm 10}$(N$_{\\rm A}$/N$_{\\rm B}$)$_{\\star}$ -- {\\rm log}$_{\\rm 10}$(N$_{\\rm A}$/N$_{\\rm B}$) $_{\\odot}$, and log~$\\epsilon$(A) $\\equiv$ {\\rm log}$_{\\rm 10}$(N$_{\\rm A}$/N$_{\\rm H}$) + 12.0, for elements A and B.} provide essential tools for the study of element production in stars and galactic chemical processing during the early stages of the evolution of our Galaxy. The number of VMP stars with measured elemental abundance ratios has been increasing rapidly in the past decade. The large modern surveys for metal-poor stars, most importantly the HK survey of Beers and colleagues \\citep{beers:92,beers:99b} and the Hamburg/ESO Survey (HES) of Christlieb and collaborators \\citep{christlieb:03}, have produced medium-resolution confirmation spectra of many thousands of metal-poor candidates; the most interesting of these have been (and are being) followed up with high-resolution spectroscopic studies on large-aperture telescopes \\citep[e.g.  ][]{cayrel:04,honda:04b,johnson:04,barklem:05}.  One of the most important discoveries of these new surveys is that at least 20\\,\\% of all VMP stars exhibit conspicuous enrichments of the CNO elements, most notably C \\citep{beers:05,lucatello:06}. These Carbon-Enhanced Metal-Poor (CEMP) stars are defined by Beers \\& Christlieb as metal-poor stars with $\\mem{[C/Fe]} > +1.0$. Many of the CEMP stars exhibit overabundances of the elements associated with s-process nucleosynthesis. Aoki et al. (2003) estimate this fraction to be between 70\\,\\% and 80\\,\\%; Beers \\& Christlieb refer to these as the CEMP-s stars. The observed abundance patterns for CEMP-s stars suggest nucleosynthetic origin in low- or intermediate-mass stars that have evolved through the thermally pulsing asymptotic giant branch (AGB) phase, and later transferred this processed material to a surviving low(er)-mass companion. In this sense the CEMP-s stars are the metal-poor analogs of the classical CH stars \\citep{keenan:42}. The AGB star that originally enriched the presently observed companion VMP star is now a white dwarf; its presence is often revealed by tell-tale radial-velocity variations observed for the companion star. Based on the still-limited numbers of CEMP-s stars for which sufficient multi-epoch spectroscopic data has been obtained, a $100$\\,\\% binary fraction is possible \\citep{lucatello:05}. It is thus expected that the large overabundances of a number of elements observed in CEMP-s stars should reflect the AGB nucleosynthetic yields, perhaps with some modification due to giant-branch evolution of the companion.  In all AGB stars C is produced by the triple-$\\alpha$ reaction in He-shell flashes and convectively dredged up into the stellar envelope. For stars of low initial mass (2--3\\,\\msun), this eventually leads to a C-rich composition, with ${\\rm C/O} > 1$. In these models N is not enhanced because it is burned during the He-shell flash. During the interpulse phase the convective envelope has no mixing connection with the H-shell, according to standard models, and no alteration of N is expected. For larger masses ($>3.5\\,\\msun$, depending on metallicity) the efficiently dredged up carbon is transformed into N by the hot-bottom burning (HBB) process \\citep{lattanzio:92b,boothroyd:93,forestini:97,herwig:04a,ventura:05}. As a result, the models predict small C/N ratios in the stellar envelope ($\\mem{[C/N]} \\approx -1$).  Simulations of low- and intermediate-mass stellar evolution confirm that the well-established trends of C and N production in solar-metallicity AGB stars extend to very low metallicity \\citep{herwig:04b}. The amount of N produced is independent of the initial metallicity of the star, because it is based on the primary production of C in the He-burning shell. Therefore, HBB in intermediate-mass VMP AGB stars provides a primary source of nitrogen in the early Universe. Figure~\\ref{Fig:AGBabund} shows the evolution of [C/N] ratios on the surface of metal-poor AGB stars from the models of \\citet{herwig:04b}. Two effects are immediately clear: (1) a lower predicted [C/N] ratio in more massive AGB stars, and (2) a lower predicted [C/Fe] on the surfaces of these same stars.   \\begin{figure} \\includegraphics[width=8cm]{f1.eps} \\caption{\\label{Fig:AGBabund}Surface abundances of AGB stars of different masses for metallicity [Fe/H] = $-2.3$ over several dredge-up episodes. As more dredge-up occurs, the surface abundance of C increases. For the higher-mass cases, dredged-up C is processed into N via HBB, leading to lower C/N ratios.  The final total yields of the elements produced by the AGB stars and ejected into the ISM by mass loss are given in \\citet{herwig:04b}.} \\end{figure}  The predicted time-averaged yields for AGB stars are shown in Figure~\\ref{Fig:CFeCNold}. In this Figure we also plot [C/N] ratios for CEMP stars that were available in the literature when we began this project. We would expect that at least one-third of the stars should have been polluted by intermediate-mass AGB stars, assuming (1) stars between 3.5\\,{\\solmass} and 7.5\\,{\\solmass} undergo HBB, (2) stars between 1.0\\,{\\solmass} and 7.5\\,{\\solmass} produce enough C to enhance their companions during mass transfer, and (3) the binary mass ratios from \\citet{pinn:06} are correct for low-metallicity stars. It is evident that the observed [C/N] ratios in CEMP stars do not fit the predictions of either the low-mass or the intermediate-mass AGB models, but instead fall in an intermediate regime that is not covered by the models.  We are left with two questions. First, why do the observations of stars with large [C/Fe] show larger N (smaller [C/N]) than predicted by models for the evolution of 2--3\\,\\msun\\ stars? This problem is reminiscent of the high N abundances observed in the moderately metal-poor CH stars, and may be related to mixing processes that are not accounted for in models of AGB evolution \\citep{vanture:92b}. This issue will be discussed later in this paper. Second, where are the VMP stars that were, in the context of the mass-transfer scenario, polluted by the N-rich 3.5--7.5\\,\\msun\\ stars? The solution for this problem is less obvious. The HBB responsible for the very efficient production of N in intermediate-mass stars is a robust prediction of stellar-evolution models. If the binary mass-transfer scenario is invoked for the CEMP-s stars with ${\\rm [C/N]} \\approx 1$, then one may ask why there are no NEMP stars with ${\\rm [C/N]} \\approx -1$, as might be expected to arise from systems in which the donor is a more massive HBB AGB star. Before considering possible solutions to this second problem we need to test one obvious possibility, that the absence of NEMP stars among previously analysed VMP stars may be simply the result of a selection bias against their detection.  Because of the requirement for large-aperture telescopes in order to obtain high-resolution, high-$S/N$ spectra of the CEMP stars in the HK and HES, only a limited number of these objects have been observed to date. Special attention has been focused on the subsample of CEMP stars that are likely to be the most iron-poor and/or carbon-rich \\citep[e.g.][]{norris:97b,aoki:02b}. This may have led to an observational bias against the discovery of NEMP stars. The CH features at 4305\\,{\\AA} are routinely covered in the medium-resolution spectroscopic follow-up of metal-poor candidates from the HK survey and the HES; thus, CEMP stars are readily identified and placed on target lists for examination at high resolution. However, N-rich stars are not so easily recognized. The medium-resolution confirmation spectra generally extend no bluer than about 3600\\,{\\AA}; the only N-sensitive feature included in this range is the CN band at 3850\\,{\\AA}, which requires large enhancements of {\\it both} C and N to be strong.  As a result, the majority of stars in the current literature with detailed studies of their elemental abundance patterns are those with ${\\rm [C/N]} > 0$, i.e., they are more C-rich than they are N-rich. Thus, the lack of stars with ${\\rm [C/N]} < 0$ could well be an observational selection effect, especially since the lower C-enhancements predicted in intermediate-mass AGB stars would make their companions potentially less C-rich, and hence less likely to be followed up at high spectral-resolution.  Previous low-resolution surveys of metal-poor (${\\rm [Fe/H]} <-1.0$) dwarf stars that included the blue NH band have discovered some stars with stronger than average nitrogen abundances, including HD~74000 and HD~160617 \\citep{bessell:82}, HD~97916 and HD~166913 \\citep{laird:85} and HD~25329 \\citep{carbon:87}. Higher-resolution follow-up studies \\citep[e.g.,][]{beveridge:94,mashonkina:03} have shown that these stars are somewhat s-process rich (${\\rm [Ba/Fe]}\\sim +0.5$). They do not appear to be C-enhanced, however, and their binary status is not confirmed. These stars are probably related to the phenomenon discussed here, but we confine ourselves in the rest of the paper to the discussion of CEMP stars found in recent surveys.  We have undertaken a medium-resolution observing campaign to address the question of whether observational selection biases might be responsible for the lack of known NEMP stars. We examine a sample of VMP stars with moderate carbon enhancements, $+0.5 \\le \\rm{[C/Fe]} \\le +1.0$ (based on their medium-resolution confirmation spectra), in order to better constrain the range of the N enhancements in these stars, and to see if the N abundances obtained agree better with expectations based on AGB models. This is accomplished using near-UV medium-resolution spectroscopy that covers the region of the NH band at 3360--3370\\,{\\AA} for a sample of 21 moderately carbon-enhanced metal-poor stars, along with a number of similar stars from the literature with available high-resolution results. In \\S\\,\\ref{Sect:Observations} we describe our sample selection criteria, observations, and data reduction procedures. Details of our abundance analysis for this sample are provided in \\S\\,\\ref{Sect:AbundanceAnalysis}. In \\S\\,\\ref{Sect:Results} we summarize our results. A discussion of the theoretical expectations and a comparison with our present results is provided in \\S\\ref{Sect:Discussion}. ",
        "conclusions": "\\tablewidth{0pt} \\tablecaption{Summary of Errors} \\tablehead{\\colhead{Star} & \\colhead{{\\teff} range} & \\colhead{$\\sigma$(log$\\epsilon$(C))} & \\colhead{$\\sigma$(log$\\epsilon$(N))} & \\colhead{$\\sigma$(log$\\epsilon$(Fe))} & \\colhead{$\\sigma$([C/Fe])} & \\colhead{$\\sigma$([N/Fe])} & \\colhead{$\\sigma$([C/N])} } \\startdata turnoff & \\teff $> 6400$K & 0.30 & 0.36 & 0.24 & 0.31 & 0.35 & 0.31 \\\\ subgiant & $5400$K$ <$\\teff$<6400$K & 0.30 & 0.31 & 0.24 & 0.30 & 0.32 & 0.30 \\\\ giant & \\teff $<5400$K & 0.25 & 0.25 & 0.21 & 0.29 & 0.25 & 0.41 \\enddata \\end{deluxetable*}  \\label{Sect:Discussion}  We expected to find that between 12\\,\\% and 35\\,\\% of our sample would turn out to be NEMP stars (see \\S\\,\\ref{Sect:IMAGBpollution}), depending on our assumption concerning the mass ratios in the binary progenitors. Instead, we found no candidates in our sample of 21 stars, and only two in the recent literature. The presence of these two ${\\rm [C/N]}\\sim -1.00$, s-process-rich stars suggest that NEMP stars do in fact exist, and their abundance patterns are as expected from the models, just not in the predicted numbers that should be found. However, there remain several outstanding questions concerning the nature of our sample which we should examine before asking if models of primary nitrogen production by intermediate-mass AGB stars are incorrect.   \\subsection{Are we observing stars that have been polluted by AGB stars?}\\label{Sect:AGBpollution}  AGB stars are by far the most likely culprits for the pollution of our sample. The majority of CEMP stars are CEMP-s stars, and much observational evidence, reviewed in the Introduction, points to their being the result of AGB mass transfer. However, it is possible that stars with lower [C/Fe] values have a smaller percentage that are s-process-rich than those with ${\\rm [C/Fe]} > 1.0$. We have information on the s-process abundances for four stars in the sample: CS~22947-187 and CS~31085-090 are s-process-rich, while CS~30314-067 and CS~22958-042 are not and there appears to be no correlation with C-richness.  We are currently obtaining follow-up high-resolution spectra for many of stars in our sample to check for radial-velocity variations and s-process enhancements.  \\subsection{Are we observing stars that have been polluted by intermediate-mass AGB stars?}\\label{Sect:IMAGBpollution}  Although we are confident that most of our stars have been polluted by AGB stars, we have considerably less information on the distribution of masses of those AGB stars. Our initial estimate was based on the mass ratio ($q$) distribution deduced by \\citet{pinn:06} from a variety of studies, particularly the data on massive binaries in the Small Magellanic Cloud from \\citet{harries:03} and \\citet{hilditch:05}. Their most likely distribution of $q$ fractions is that 45\\,\\% of binary systems are ``twins'', with the primary and secondary masses the same to within 5\\,\\%, while the $q$ of the other 55\\,\\% of systems are matched by a flat mass distribution. We assumed that the observed stars, being main-sequence turnoff, subgiant, giant or horizontal branch stars, all have main-sequence masses close to 0.9\\,{\\solmass}. Therefore, their ``twin'' AGB stars would be $< 1$\\,{\\solmass} and would have created high [C/N] ratios. Therefore, the NEMP stars must come from the flat mass distribution. We adopt 3.5\\,{\\solmass} as the dividing line between those AGB stars that undergo HBB and those that do not. Finally, the maximum mass of an AGB star at ${\\rm [Fe/H]}=-2.3$ is 7.5\\,{\\solmass} (Poelarends et al., in prep). Combining all of this information, we find that there should be roughly 33\\,\\% NEMP stars and 67\\,\\% CEMP-s stars contained in our sample.  There may be difficulties with our predictions.  For example, while \\citet{pinn:06} consider the mass ratios using studies that cover a wide range of masses, none of the studies is able to cover a wide range of $q$. The SMC binary studies only reach a mass ratio of $\\sim 0.4$, which is larger than minimum $q$ we are interested in, 0.12, and it is conceivable that the distribution will change as the mass difference between the two components becomes more extreme. Such studies are difficult because of the different luminosities and timescales of the two kinds of stars, but we do have some information on the $q$ distribution of binaries with solar-type primaries \\citep{duquennoy:91}. For the longer period ($P > 3000$ days) binaries, the data suggested a secondary star drawn from an IMF-weighted distribution. However, for the shorter period binaries, \\citet{mazeh:92} found a distribution that was much flatter, with perhaps a rise toward $q=1$. For lack of other information, if we make the assumption that the mass of the companion of the observed stars is drawn from a \\citet{salpeter:55} mass function between 1\\solmass{} and 7.5 \\solmass, we find that there should be 1 NEMP for every 6.9 CEMP-s stars. This is a much smaller ratio, but still not small enough to be supported by the observational data. All of these numbers were obtained assuming that there was not a bias towards finding NEMP stars in the C-mild sample; if such a bias existed, the disagreement with theory would only be strengthened.  \\subsection{Could binarity be affecting our results?}  The fact that these AGB stars were in binary systems may have affected their nucleosythesis yields by increasing their mass-loss rates and ending the production of N by HBB. This could be due to tidal synchronization between rotational and orbital motions, as observationally demonstrated for RGB stars by \\citet{demedeiros:02}. Even after HBB is terminated, dredge-up can continue if the minimum envelope mass for dredge-up is less than that for HBB, and the resulting [C/N] ratio may be altered. \\citet{frost:98} argued that the presence of high- luminosity C stars in the Large Magellanic Cloud could be explained by continued dredge-up after HBB has ceased.  It is not clear if this same reasoning would apply to our present sample, which has much lower metallicity. However, if tidal synchronization leads to a spin-up mixing, N production should be enhanced, possibly balancing the effect of enhanced mass loss. Without detailed models we can only speculate, but it is not immediately clear that binarity could inhibit N production at the bottom of more massive AGB stars.  \\subsection{Could mass ratios affect mass transfer?}  \\citet{jorissen:92b} summarize the situation regarding mass transfer in both Roche-lobe overflow and wind models. For large mass ratios, only wind models were stable, because the large mass-transfer rates in overflow scenarios puff up the secondary star and lead to a contact binary situation \\citep{tout:91}.  \\citet{vanture:92b} argued that the mass transfer from a wind would not be sufficient to create CH stars, and therefore only in binary systems with Roche-lobe overflow would the nucleosynthesis of the AGB star be preserved in the still-shining secondary. Much work remains to be done in understanding mass-transfer scenarios, but if NEMP stars can only be formed by wind accretion, then their production could be far less efficient than for CEMP stars.  \\subsection{Could mass transfer rates be affected C- vs. N-richness?}  Our understanding of mass loss in AGB stars is seriously incomplete. However, one variable that may be important is the dust-to-gas ratio, with mass-loss rates increasing for increasing amounts of dust. Carbon is a major contributor to dust in AGB stars, and the dust-to-gas ratio has been parameterized by the C-richness \\citep[e.g.,][]{arndt:97}. Therefore, since HBB stars produce less C-rich AGB envelopes, their mass loss may be lower and therefore the number of binary companions that are enriched enough to be included in our sample may be small. However, a lower mass-loss rate, given enough time, could still conceivably add up to the same amount of enrichment. Additionally, \\citet{arndt:97} found only a small dependency on C-enhancement.  \\subsection{Does hot bottom-burning occur in intermediate mass AGB stars?}  Another explanation for the absence of NEMP stars in our sample could be the inability of more massive AGB stars (maybe only at very low metallicity) to produce N via HBB. The efficiency of envelope convection, parameterized by the mixing length parameter, determines the efficiency of HBB in stellar evolution models. A larger convective efficiency would imply even more efficient HBB, but could decrease the dredge-up efficiency, in particular if significant convective overshoot is present. This would limit the ability of massive AGB stars to produce primary N. A smaller convective efficiency would lead to less efficient HBB, and a smaller mass range for efficient N production, lowering the expected number of NEMP stars. However, we have not performed calculations with a range of convective efficiencies, and therefore we cannot quantify these possible effects. Stellar evolution modeling of HBB in massive AGB stars also predicts the production of Li through the \\citet{cameron:71} mechanism. This property of HBB is observationally confirmed by the discovery that almost all of the luminous AGB stars with ${\\rm C/O} < 1$ in the Magellanic Clouds are lihtium rich \\citep[e.g.][]{smith:95}.  \\subsection{Could extra-mixing be affecting our results?}  The main goal of our study was to establish whether there existed a possible observational selection bias against the discovery of NEMP stars. However, we also noted that previous literature data, as well as our new objects have C/N ratios that are too low compared to standard low-mass AGB models without HBB. Our 2\\,{\\solmass} and 3\\,{\\solmass} model sequences produce large amounts of C, but no N.  The most likely solution to this discrepancy would be an additional mixing process, such as cool-bottom processing, which has been studied in great detail in the context of abundance anomalies in globular cluster stars \\citep{wasserburg:95,denissenkov:03b}. Such an extra-mixing process below the bottom of the convective envelope could also operate in low-mass AGB stars, leading to mild CN cycling that produces some N, but not as complete as in HBB \\citep{nollett:03}. We believe that the present observational data support such a scenario. We plan to test this idea quantitatively in the future with stellar evolution models that include this process.   \\subsection{Final Thoughts}  It is probable that the formation of C-rich stars involves a delicate balancing act in which the N-rich stars are disadvantaged. The most likely possibilities are the binary mass ratios in the early Universe or the process of mass transfer for large mass-ratio systems. With ever larger samples of metal-poor stars with sufficient spectroscopic data becoming available, we expect to continue studies of the relative fractions of CEMP vs. NEMP stars. To confirm that the majority of the CEMP stars in our sample were indeed polluted by AGB progenitors, we and other authors are obtaining high-resolution spectra for many of these stars to look for s-process enhancements, and to check for detectable radial-velocity variations. In addition, other elements, such as Li, Na, F, and the isotopes of Mg, which are affected by nucleosynthesis in AGB stars, will be examined in due course."
    },
    "0608/astro-ph0608499_arXiv.txt": {
        "abstract": "The E/SO sequence of a cluster defines a boundary redward of which a reliable weak lensing signal can be obtained from background galaxies, uncontaminated by cluster members. For bluer colors, both background and cluster galaxies are present, reducing the average distortion signal by the proportion of unlensed cluster members. In deep Subaru and HST/ACS images of A1689, we show that the tangential distortion of galaxies with bluer colors falls rapidly toward the cluster center relative to the background reference level provided by the red background. We use this dilution effect to derive the cluster light profile and luminosity function to large radius, with the advantage that no subtraction of far-field background counts is required. The light profile of A1689 is found to decline steadily to the limit of the data, $r<2~h^{-1}$Mpc, with a constant slope, $d\\log(L)/d\\log(r)=-1.12\\pm0.06$, unlike the lensing mass profile which steepens continuously with radius, so that $M/L$ peaks at an intermediate radius, $\\simeq 100~h^{-1}$ kpc. A flatter behaviour is found for the more physically meaningful ratio of dark-matter to stellar-matter, when account is made of the color-mass relation of cluster members. We derive a cluster luminosity function with a flat faint-end slope of $\\alpha=-1.05\\pm0.07$, nearly independent of radius and with no faint upturn to $M_{i'}<{-12}$. We establish that the very bluest objects are negligibly contaminated by the cluster ($[V-i']_{AB}<0.2$), because their distortion profile rises continuously towards the center following the red background ($[V-i']_{AB}>1.2$), but is offset higher by $\\simeq 20\\%$. This larger amplitude is consistent with the greater estimated depth of the faint blue galaxies $\\langle z\\rangle\\sim2$ compared to $\\langle z\\rangle\\sim0.85$ for the red background, a purely geometric effect. With a larger sample of background galaxies behind several clusters we may use this geometric effect to constrain the cosmological parameters in a model independent way, by comparing the weak lensing strength over a wide range of photometrically derived redshifts. Finally, we improve upon our earlier modeling of the mass profile by combining both the red and blue background populations, and very clearly exclude low concentration profiles predicted for massive CDM dominated haloes. ",
        "introduction": "The influence of ``Dark matter'' is strikingly evident in the centers of massive galaxy clusters, where large velocity dispersions are measured and giant arcs are often visible.  Cluster masses may be estimated by several means, leading to exceptionally high central mass-to-light ratios, $M/L_R \\sim 100-300h(M/L_R)_{\\odot}$ \\citep{2001ApJ...552..427C}, far exceeding both the mass of stars comprising the light of the cluster galaxies and the mass of plasma derived from X-ray emission and the SZ effect.  Reasonable consistency is claimed between dynamical, hydrodynamical and lensing-based estimates of cluster masses, supporting the conventional understanding of gravity.  However, the high mass-to-light ratio requires that an unconventional non-baryonic dark material dominates the mass of clusters, whose origin remains very unclear.  In detail, discrepancies are often reported between masses derived from strong lensing and X-ray measurements, with the claimed X-ray masses often lower in the centers of clusters. High resolution X-ray emission and temperature maps reveal that the majority of local clusters undergo repeated merging with sub-clumps, and obvious shock fronts are seen in some cases \\citep{2002ApJ...567L..27M,2004ApJ...608..179R}. The double cluster 1E~0657--56 (a.k.a., the ``bullet'' cluster, $z=0.296$) is the most extreme example studied, where the associated gas forms a flattened luminous shock-heated structure lying between the two large distinct clusters, clearly indicating these two bodies collided recently at a high relative velocity \\citep{2004ApJ...606..819M}, with the gas remaining in between while the cluster galaxies have passed through each other relatively collisionlessly. Very interestingly, the weak lensing signal follows the double structure of the clusters, rather than the gas in between, cleanly demonstrating that the bulk of the matter is collisionless and dark \\citep{2004ApJ...604..596C}. This system favors the standard cold dark matter (CDM) scenario, places a restrictive limit on the interaction cross-section of any fermionic dark matter, and disfavors a class of alternative gravity theories in which only baryons are present \\citep{2004ApJ...604..596C}. Smaller but significant discrepancies of the same sort are also claimed for other interacting clusters \\citep{2002ApJ...580L..17N,2005ApJ...618...46J}.   Many clusters show no apparent signs of significant ongoing interaction; these have centrally symmetric X-ray emission and little obvious substructure \\citep{1998MNRAS.296..392A}, and for some of these the lensing and X-ray (or dynamical) derived masses are claimed to agree, as would be expected for relaxed systems \\citep[e.g.,][]{2002ApJ...572...66A,2003AJ....126.2152R,2005ApJ...628L..97D}. More generally, since the dyanmics of dark matter and most cluster galaxies is essentially collisionless, we would expect them to have similar radial profiles.  Biasing inherent in hierarchical growth may significantly modify this similarity \\citep{1997MNRAS.286..795K}, and dynamical friction is expected to concentrate the relatively more massive galaxies in the core. These together with tidal effects may help to account for the unique properties of cD galaxies. Hence comparisons of the mass profile with the light profile of the cluster galaxies are expected to provide an additional insight into the formation of clusters and the nature of dark matter.  Recent improvements in the quality of data useful for weak and strong lensing studies now allow the construction of much more definitive mass profiles that are sufficiently precise to test the predictions of popular models, relatively free of major assumptions.  The inner mass profiles of several clusters have been constrained in some detail via lensing, using multiply lensed background galaxies \\citep{1996ApJ...471..643K,1997ApJ...491..477H,2002ApJ...574L.129S, 2003A&A...403...11G,2004ApJ...604...88S,2005ApJ...621...53B, 2005ApJ...629L..73S}. The statistical effects of weak lensing have been used to extend the mass profile to larger radii. The mass profiles derived from these observations have been claimed to show NFW-like behaviour \\citep*{1997ApJ...490..493N}, with a continuous flattening towards the center, but with higher density concentrations than expected. This is particularly evident for A1689, for which we have constructed the highest quality lensing based mass profile to date, combining over 100 multiply-lensed images and weak lensing effects of distortion and magnification from Subaru \\citep{2005ApJ...621...53B,2005ApJ...619L.143B}.  A lesson learned from this earlier work is the importance of carefully selecting a background population to avoid contamination by the lensing cluster. It is not enough to simply exclude a narrow band containing the obvious E/SO sequence, following common practice, because the lensing signal of the remainder is found to fall rapidly towards the cluster center, in contrast to the uncontaminated population of background galaxies lying redward of the cluster sequence \\citep{2005ApJ...619L.143B}. In the well studied case of A1689, there has been a long standing discrepancy between the strong and weak lensing effects, with the weak lensing signal underpredicting the observed Einstein radius by a factor of $\\sim 2.5$ \\citep{2001A&A...379..384C,2005A&A...434..433B}, based only on a minimal rejection of obvious cluster members using one or two-band photometry.  In our recent weak lensing analysis of Subaru images of A1689 the above behaviour was found when we rejected only the cluster sequence in the same way as others. This resulted in a relatively shallow trend of the weak lensing signal with radius and consequently an underprediction of the Einstein radius accounting for the earlier discrepancy \\citep{2005ApJ...619L.143B} . If, however, one selects only objects redder than the cluster sequence for the lensing analysis, then the weak tangential distortion continues to rise all the way to the Einstein radius in very good agreement with the strong lensing strength. This red population is naturally expected to comprise only background galaxies, made redder by relatively large k-corrections and with negligible contamination by cluster members, since the bulk of the reddest cluster members are the early-type galaxies defined by the cluster sequence. However, for galaxies with colors bluer than the cluster sequence, cluster members will be present along with background galaxies since the cluster population extends to bluer colors of the later-type members, overlapping in color with the blue background. The effect of the cluster members is simply to reduce the strength of the weak lensing signal when averaged over a statistical sample, in proportion to the fraction of cluster members whose orientations are randomly distributed, therefore diluting the lensing signal relative to the reference background level derived from the red background population.  We can turn this dilution effect to our advantage and use it to derive properties of the cluster population, in particular the radial light profile, for comparison with the dark matter profile. Deriving a light profile this way has advantages over the usual approach to defining cluster membership.  The inherent fluctuations in the number counts of the background population are a significant source of uncertainty in the usual approach of subtracting the far-field level when defining the cluster population \\citep[e.g.,][]{2001A&A...367...59P,2005MNRAS.360..727A,2005MNRAS.364.1147P}. This uncertainty is often cited as a potential explanation for the substantial variation reported between luminosity functions derived for different clusters, particularly in the outskirts of clusters, where not only is the density of galaxies lower, but their colors are bluer, thus harder to distinguish from the background using photometry alone.  In \\S~\\ref{data} we describe the observations and photometry of the Subaru images of A1689. In \\S~\\ref{dist_Sub} we describe the distortion analysis applied to the Subaru data. The distortion analysis of ACS images of  A1689 is described in \\S~\\ref{dist_ACS}. In \\S~\\ref{Photo-z} we describe the photometric redshift analysis of the Subaru and ACS images with reference to the \\cite{2004AJ....127..180C} sample of deep multi-color Subaru images. Our weak lensing dilution analysis of the Subaru images is described in \\S~\\ref{dilution}. In \\S~\\ref{lum_sec} we go on to derive the cluster luminosity profile and color, and in \\S~\\ref{lum_func} the cluster luminosity function is deduced at several radial positions. In \\S~\\ref{ML_sec} we determine the $M/L$ profile, and in \\S~\\ref{MLconsist} we do a consistency check for the mass derived in this paper with previous estimations. Our conclusions are summarized in \\S~\\ref{discussion}.  The concordance $\\Lambda$CDM cosmology is adopted ($\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$ but $h$ is left in units of $H_0/100$ km s$^{-1}$ Mpc$^{-1}$, for easier comparison with earlier work). ",
        "conclusions": "\\label{discussion}  We have explored a new approach to deriving the luminous properties of cluster galaxies by utilizing lensing distortion measurements, based on the dilution of the lensing distortion signal by unlensed cluster members which we assume are randomly oriented. We have tested this assumption for a restricted sample of bright cluster galaxies which project beyond the faint galaxy background population so the level of background contamination is negligible, confirming that the cluster galaxies are randomly oriented with a negligible net tangential distortion for the purposes of our work.  This dilution approach is applied to A1689 to derive the radial light profile of the cluster, a color profile and radial luminosity functions. The light profile is found to be smoothly declining and fitted with a power-law slope $d\\log(L)/d\\log(r)=-1.12\\pm0.06$. We also see a mild color gradient corresponding to a change in the cluster population from early- to mid-type galaxies in moving from the center out to the limit of our data at 2~$h^{-1}$Mpc. Unlike the light profile the gradient of mass profile is continuously steepening, such that the ratio of $M/L$ peaks at intermediate radius. We find that the cluster luminosity function has a flat faint-end slope of $\\alpha=-1.05\\pm0.07$, nearly independent of radius and with no faint upturn to $M_{i'}<{-12}$.  A major advantage of our approach is that we do not need to define far-field counts for subtracting a background, as in the usual method where there is a limitation imposed by the clustering of the background population that limits the radius to which a reliable subtraction of the background can be made.  We have also established that the bluest galaxies in the field of A1689 lie predominantly in the background, as their radial distortion profile follows closely the red galaxies, but with an offset indicating the blue population lies at a greater mean distance than the red background galaxies, and consistent with the estimated mean redshifts of these two populations. With a larger sample of clusters, this purely geometric effect can potentially be put to use to provide a simple model-independent measure of the cosmological curvature.   The mass profile of A1689 was reexamined using our combined background sample of red and blue galaxies. The distortion profile derived from this sample is consistent with our earlier work, but somewhat more statistically significant, so we have examined the mass profile more carefully out to a larger radius. We have found that the distortion profile is steeper than predicted for CDM haloes appropriate for cluster sized masses $C_{vir}\\sim 5$. This discrepancy is particularly clear at large radius $r>2\\arcmin$, where an acceptable fit is found to an NFW profile but with a concentration $C_{vir}=27^{+3.5}_{-5.7}$. This finding is consistent with our earlier work which showed that although an overall best fit profile of $C_{vir}\\simeq 14$ to the joint strong and weak lensing based data presented in \\cite{2005ApJ...619L.143B}, the curvature of the data is more pronounced than an NFW profile being shallower in the inner region and steeper at larger radius, so that the derived value of the concentration increases with radius depending on the radial limits being examined.  This result is surprising and may require a significant departure from the standard CDM model, either in terms of the mass content, or the epoch at which the bulk of the cluster was assembled. For example, one possibility to achieve earlier formation of clusters is to allow deviation from Gaussianity of the primordial density fluctuation field, as has been considered recently by, e.g., \\cite{2006MNRAS.368.1583S}. A1689 is amongst the most massive known clusters, and projection effects may play a role in boosting somewhat the lensing signal along the line of sight. We therefore aim to test the generality of this result with a careful study of a statistical sample of clusters.  Upcoming spatially resolved SZ measurements will add a significant new ability to determine cluster mass profiles over a large range of radius, and allow for improved consistency checks between the various independent means of estimating masses. The combination of X-ray, lensing and SZ measurements will soon lead to far greater accuracy in understanding the nature of cluster mass profiles.  We plan an improvement to the weak lensing work with deeper multi-color imaging from Subaru for measuring reliable photometric redshifts for a sizable fraction of the background population. This added dimension of depth will enhance the weak lensing signal and reduce the systematic problems of cluster and foreground contamination of the lensing signal. We also aim to extend this work to well studied clusters at lower redshift with archived Subaru imaging and detailed X-ray and upcoming SZ observations as in principle, the lensing signal should be equally strong for lower redshift clusters, given the maximal ratio of lens to source distances, $d_{ls}/d_s\\simeq 1$, for faint background sources."
    },
    "0608/hep-th0608121_arXiv.txt": {
        "abstract": "It has recently been shown that a Hagedorn phase of string gas cosmology may provide a causal mechanism for generating a nearly scale-invariant spectrum of scalar metric fluctuations, without the need for an intervening period of de Sitter expansion. A distinctive signature of this structure formation scenario would be a slight blue tilt of the spectrum of gravitational waves. In this paper we give more details of the computations leading to these results. ",
        "introduction": "String gas cosmology (SGC) \\cite{BV,TV} (see also \\cite{Perlt} for early work, \\cite{rbr1,rbr2} for recent overviews, and \\cite{str} for a critical review) is an approach to string cosmology, based on adding minimal but crucial inputs from string theory, namely new degrees of freedom - string oscillatory and winding modes - and new symmetries - T-duality - to the now standard hypothesis of a hot and small early universe (see \\cite{others} for other work on string gas cosmology, and in particular \\cite{ABE} for a discussion of the role which D-branes play in this scenario). One key aspect of SGC is that the temperature cannot exceed a limiting temperature, the Hagedorn temperature $T_H$ \\cite{Hagedorn}. This immediately provides a qualitative reason which leads us to expect that string theory can resolve cosmological singularities \\cite{BV}. From the equations of motion of string gas cosmology \\cite{TV,Ven} it in fact follows that if we follow the radiation-dominated Friedmann-Robertson-Walker (FRW) phase of standard cosmology into the past, a smooth transition to a quasi-static Hagedorn phase will occur. In this phase, the string metric is quasi-static while the dilaton is time-dependent. Reversing the time direction in this argument, we can set up the following new cosmological scenario \\cite{BV}: the universe starts out in a quasi-static Hagedorn phase during which thermal equilibrium can be established over a large scale (a scale sufficiently large for our current universe to grow out of it following the usual non-inflationary cosmological dynamics). The quasi-static phase, however, is not a stable fixed point of the dynamics, and eventually a smooth transition to a radiation-dominated FRW phase will occur, after which point the universe evolves as in standard cosmology. This transition is connected to the decay of string winding modes into string oscillatory degrees of freedom, as described initially in \\cite{BEK} and studied later in more depth in \\cite{Col,Frey}.  Recently \\cite{Ali1,Ali3} it was discovered that, under certain assumptions which will be detailed below, string thermodynamic fluctuations during the Hagedorn phase lead to scalar metric fluctuations which are adiabatic and nearly scale-invariant at late times, thus providing a simple alternative to slow-roll inflation for establishing such perturbations. It is to be emphasized that this mechanism for the generation of the primordial perturbations is intrinsically stringy - particle thermodynamic fluctuations would lead to a spectrum with a large and phenomenologically unacceptable blue tilt \\cite{Ali3}. As discussed in \\cite{Ali2}, the stringy thermal fluctuations also generate an almost scale-invariant spectrum of gravitational waves (tensor metric fluctuations). Whereas the spectrum of scalar modes has a slight red tilt (like the fluctuations generated in most inflationary models), the gravitational wave spectrum has a slight blue tilt. This is a prediction which distinguishes our string gas cosmology scenario from standard inflationary models. In inflationary models, a slight red tilt of the spectrum is expected since the amplitude of gravitational waves is set by $H$, and $H$ is decreasing in time, leading to lower values of $H$ when scales with larger values of the comoving wavenumber $k$ exit the Hubble radius.  In this paper, we review the analyses of our previous results \\cite{Ali1,Ali3,Ali2} and present more details. The outline of this paper is as follows. We begin with a discussion of the background string gas cosmology, emphasizing the smooth transition between the initial quasi-static Hagedorn phase and the later radiation-dominated period. Next, we describe how the scalar and tensor metric fluctuations are related to correlation functions of the string gas energy-momentum tensor. In Section 4 we then discuss the computation of the string gas correlation functions. Sections 5 and 6 describe the evolution of scalar and tensor metric fluctuations, respectively. In the final section we discuss our results, point out unresolved issues and future directions for research.  In the following, we assume that our three spatial dimensions are already large (sufficiently large such that expansion according to standard cosmology beginning at a temperature of the string scale can produce our observed universe today - if the string scale is of the order $10^{16}$ GeV, then the initial size needs to be of the order of 1mm) during the Hagedorn phase (for a possible mechanism to achieve this see \\cite{Natalia1}), while the extra spatial dimensions are confined to the string scale. For a mechanism to achieve the separation into three large and six small compact dimensions in the context of SGC see \\cite{BV} (see however \\cite{Col,Frey} for caveats), and for a natural dynamical mechanism arising from SGC to stabilize all of the moduli associated with the extra spatial dimensions see \\cite{Watson,Patil1,Patil2,Patil3,Edna,shiggs} (see also \\cite{other2}).  To be specific, we take our three dimensions to be toroidal. The existence of one cycles results in the stability of string winding modes - and this is a key ingredient in our calculations. ",
        "conclusions": "In this paper we have presented more details of the derivation of the power spectra for scalar and tensor metric fluctuations. Here we summarize the main results.  First of all, we have found that our string gas cosmology background leads to a new way of obtaining a roughly scale-invariant spectrum of cosmological perturbations. Our mechanism relies on string theory in two crucial ways. Firstly, the cosmological background is the result of a new symmetry of string theory which is not present in ordinary quantum field theory, namely T-duality. The background equations of motion which are consistent with the T-duality symmetry, namely those of dilaton gravity, predict that the radiation-dominated phase of standard cosmology was preceded by a phase - the Hagedorn phase - during which the scale factor in the string frame is static, the pressure vanishes, and the temperature is close to the limiting temperature, the Hagedorn temperature. As we have shown, string thermodynamical fluctuations in the Hagedorn phase lead to an almost scale-invariant spectrum of cosmological perturbations. This is the second point where string theory plays a crucial role. Point particle fluctuations would have given a completely different spectrum.  We have shown that the spectrum of the scalar metric fluctuations has a slight red tilt, whose value depends on how fast the transition between the Hagedorn phase and the radiation-dominated period occurs.  Our model makes a unique prediction: the spectrum of gravitational waves has a slight blue tilt, unlike the slight red tilt which is predicted in scalar field-driven inflationary models. It would be interesting to investigate how hard it would be to observationally identify this signature. One factor which will help in this regard is the fact that the tensor to scalar ratio $r$ in our scenario is not suppressed by slow-roll parameters. As can be seen immediately by comparing our results for the amplitude of the tensor and scalar power spectra, respectively, the ratio $r$ at a scale $k$ is given by \\be r \\, \\sim \\, (1 - T(t_i(k))/T_H)^2 \\ln^2{\\left[\\frac{1}{l_s^2 k^2}(1 - T(t_i(k))/T_H)\\right]}\\, . \\ee  In principle (if the dynamical evolution from the Hagedorn phase to the radiation-dominated FRW phase were under complete analytical control) this quantity would be calculable. If the string length were known, the factor $(1 - T/T_H)$ could be determined from the normalization of the power spectrum of scalar metric fluctuations. Since the string length is expected to be a couple of orders larger than the Planck length, the above factor does not need to be very small. Thus, generically we seem to predict a ratio $r$ larger than in simple slow-roll inflationary models.  We close by mentioning some limitations of our analysis. Since the dilaton is evolving until the beginning of the radiation phase, the gravitational perturbations should be determined and evolved using the complete dilaton gravity background, not using the equations of general relativity. The time dependence of the dilaton will thus lead to correction terms. We are currently in the process of calculating the magnitude of these effects. This calculation will also be crucial in order to give a correct estimate of the amplitude of the slight tilt in our spectra. These calculations are also currently in progress.  We should also emphasize that, at least in the framework presented in this paper, our cosmological scenario does not solve the entropy and flatness problems of standard cosmology. In other words, we need to start with three spatial dimensions which at the time $t_R$ are much larger than the string scale (the entropy problem), and we can provide no explanation for why the current universe is so close to being spatially flat. Inflationary cosmology, in addition to providing a causal mechanism for structure formation, has the potential of solving the entropy and flatness problems. On the other hand, scalar field-driven inflationary cosmology is singular, whereas string gas cosmology has the potential of providing a non-singular cosmology."
    },
    "0608/astro-ph0608336_arXiv.txt": {
        "abstract": "We present the detection and modeling of more than 70 far--IR pure rotational lines of water vapor, including the $^{18}$O and $^{17}$O isotopologues, towards Orion~KL. Observations were performed with the \\textit{Long Wavelength Spectrometer} Fabry--Perot (LWS/FP; $\\lambda/\\Delta \\lambda \\sim$6800-9700) on board the \\textit{Infrared Space Observatory} (ISO) between $\\sim$43 and 197~$\\mu$m. The water line profiles evolve from P--Cygni type profiles (even for the H$_{2}^{18}$O lines) to pure emission at wavelengths above $\\sim$100~$\\mu$m. We find that most of the water emission/absorption arises from an extended flow of gas expanding at 25$\\pm$5~km~s$^{-1}$. Non--local radiative transfer models show that much of the water excitation and line profile formation is driven by the dust continuum emission. The derived beam averaged water abundance is 2--3$\\times$10$^{-5}$. The inferred gas temperature $T_k$=80--100~K suggests that: $(i)$ water could have been formed in the \\textit{plateau} by gas phase neutral-neutral reactions with activation barriers if the gas was previously heated (e.g.~by shocks) to $\\geq$500~K and/or $(ii)$ H$_2$O formation in the outflow is dominated by \\textit{in--situ} evaporation of grain water--ice mantles and/or $(iii)$ H$_2$O was formed in the innermost and warmer regions (e.g. the hot core) and was swept up in $\\approx$1000~yr, the dynamical timescale of the outflow. ",
        "introduction": "Star forming regions are associated with violent phenomena such as cloud collapse, molecular outflows and related shocked regions. Under these conditions the neutral gas is warm and  water vapor is predicted to be abundant \\citep{dra82} and to play a dominant role in the thermal balance \\citep{neu93}. Unfortunately, ground--based observations of H$_2$O are difficult and so the determination of water column densities is not straightforward. Nevertheless, ground--based observations of water maser lines have been performed toward Orion~KL. From VLBI observations of the 6$_{16}$--5$_{23}$ line at $\\sim$22~GHz, \\citet{gen81} determined the kinematics and the expansion velocity for the so--called \\textit{low velocity outflow} ($\\sim$18$\\pm$2~km~s$^{-1}$). The widespread nature of water vapor has been probed with maps of the 3$_{13}$--2$_{20}$ line at $\\sim$183~GHz \\citep{cer90,cer94}. This was the first time that the water abundance was estimated in the different large scale components around Orion~KL. In addition, the high excitation conditions of the \\textit{plateau} gas (a mixture of  outflows, shocks and interactions with the ambient cloud) has been revealed by observations of the 5$_{15}$--4$_{22}$ line at $\\sim$325~GHz \\citep{men90,cer99b}. Due to their maser nature, the water abundance determination from these lines is quite involved. Even the observation of the H$_{2}^{18}$O 3$_{13}$--2$_{20}$ line at $\\sim$203~GHz gives only poor estimates due to the overlap with other molecular lines \\citep{jac88}. The HDO species has been also detected and modeled towards Orion~KL \\citep{jac90,par01}, but the specific gas and dust chemistry has to be taken into account to derive $\\chi$(H$_2$O).  ISO \\citep{kes96} spectrometers have provided the opportunity to observe many IR H$_2$O lines towards bright sources such as Orion~KL. In particular, Harwitt et al. (1998) presented eight ISO/LWS/FP \\citep{cle96} water lines between $\\sim$71 and $\\sim$125~$\\mu$m involving energy levels between 300 and 800~K. They estimated $\\chi$(H$_2$O)$\\simeq$5$\\times$10$^{-4}$; however, the interpretation of these lines should include radiative pumping from IR dust photons. A larger set of weaker IR lines, including those of H$_{2}^{18}$O and H$_{2}^{17}$O, is needed to minimize the opacity effects always associated with H$_{2}^{16}$O lines. Orion~KL has also been targeted with the \\textit{Short Wavelength Spectrometer} (SWS) tracing a smaller region than that probed by the LWS beam below $\\sim$45~$\\mu$m \\citep{vd98}. Nineteen water absorption lines with energies between 200 and 750~K were detected \\citep{wri00}. These authors suggested that mid--IR water lines arise from the \\textit{low--velocity outflow} and estimated $\\chi$(H$_2$O)=2--5$\\times$10$^{-4}$ (for an assumed gas temperature of 200--350~K). Most recently, the  1$_{10}$--1$_{01}$ lines of H$_{2}^{16}$O and H$_{2}^{18}$O have been observed with SWAS ($\\chi$(H$_{2}^{16}$O)$\\simeq$3.5$\\times$10$^{-4}$; \\cite{mel00}; \\cite{sne00}) and ODIN ($\\chi$(H$_{2}^{16}$O)$\\simeq$10$^{-5}$--10$^{-4}$; \\cite{olo03}). In these observations, the H$_{2}^{16}$O ground--state line shows widespread emission over a $\\sim$5$'$ region.  In this letter we present all water lines detected in the first far--IR line survey towards Orion~KL carried out with the ISO/LWS spectrometer (Lerate et al. 2006), and the radiative transfer models that fit the water data set (involving levels up to 2000 K).   \\begin{figure*} [ht] % \\centering \\includegraphics[angle=0,width=14cm]{f1.eps} \\caption{Summary of the 70 far--IR water lines detected by the ISO/LWS-FP towards Orion KL. The ordinate corresponds to the continuum normalized flux and the abscissa to the velocity (in km~s$^{-1}$). Transition rotational numbers, rest wavelengths (in $\\mu$m), upper level energies (in K) and intrinsic line strengths are shown in each box. The central inset shows the ISO/LWS-grating observations towards Orion at a resolution of $\\lambda/\\Delta \\lambda \\sim 300$. The ordinate corresponds to the absolute flux (in W~cm$^{-2}$~$\\mu$m$^{-1}$) and the abscissa to the wavelength (in $\\mu$m). Main molecular features  between $\\sim$160 and $\\sim$197~$\\mu$m are labelled in a zoom to the grating spectrum. The full radiative transfer model for the far--IR continuum and water line spectrum is also shown in the upper inset.} \\label{fig-observations} \\end{figure*}%   \\clearpage ",
        "conclusions": "A summary of the resulting water lines is shown in Fig.~\\ref{fig-observations}. From these far--IR observations it is clear that water lines show a complicated behavior. For wavelengths above $\\sim$100~$\\mu$m, H$_{2}^{16}$O  lines  are observed in emission. However, for shorter wavelengths, lines arising from energy levels below $\\sim$1000~K and with large line--strengths, show P--Cygni profiles with emission covering a large velocity range. However, those with weak line-strengths or arising from higher energy levels, are observed in pure emission. H$_{2}^{18}$O lines also show P--Cygni type profiles below $\\sim$100~$\\mu$m. In the H$_{2}^{18}$O  case, the absorption component is deeper than in the analogous H$_{2}^{16}$O transition. Since optical depth effects  are much less severe in H$_{2}^{18}$O lines, the associated P--Cygni type profiles trace the main origin of far--IR water lines toward Orion~KL, i.e. an extended outflow. Pure emission H$_2$O lines peak around $v\\simeq$20~km~s$^{-1}$ (the source $v_{LSR}$ is $\\sim$9~km~s$^{-1}$; Scoville et al. 1993) but it is likely that the most opaque lines are redshifted due to self--absorption. On the other hand, water lines detected below $\\sim$50~$\\mu$m are observed in pure absorption with a velocity peak of $v\\simeq-$10~km~s$^{-1}$ (the same applies to most water lines observed by ISO/SWS; \\citet{wri00}).  The turnover point between absorption and emission  lines is an important clue to interpret this large data set and to determine the relations between continuum+line opacity, line strengths, spatial distribution of gas, and physical conditions. ISO observations clearly show that most of the water vapor detected in the IR arises from a flow of gas expanding at 25$\\pm$5~km~s$^{-1}$. The inferred expansion velocity is consistent with the \\textit{low--velocity outflow} originally revealed by $\\sim$22~GHz  H$_2$O maser motions \\citep{gen81}, but a contribution from the extended \\textit{high--velocity outflow} could be present \\citep{cer94}. However, no high velocity line--wing emission is observed at ISO's sensitivity and S/N. Similar conclusions have been found for the OH excited rotational lines \\citep{goi06}.  The main problem with modeling ISO data is the limited spatial resolution, which makes difficult to constrain the size and origin of the water region. Besides, any detailed fit to the data requires a detailed knowledge of the geometry and of the relative distribution of dust continuum and water lines. Fortunately,  high angular resolution maps of water at $\\sim$183~ and $\\sim$325~GHz lines do reveal the spatial and velocity distribution of water in Orion \\citep{cer94,cer99b}. In particular, Cernicharo and coworkers detected at least four different water components: the extended  \\textit{ridge} (extended quiescent gas), the \\textit{plateau} (including the \\textit{high--} and \\textit{low--velocity outflows}) and the very narrow and strong features associated with the small water bullets observed at $\\sim$22~GHz \\citep{gen81}. Besides, the newly detected far--IR water lines associated with the highest energy levels may have a contribution from the \\textit{hot core} (dense and hot inner regions). However, the large far--IR line--plus--continuum opacity  will probably hide most of the \\textit{hot core} emission. Finally,  radiative transfer effects in the most opaque lines, e.g. self--absorption and/or scattering by a  lower density diffuse halo may possibly occur at velocity scales not resolved by ISO.\\\\     In order to estimate the water abundance, and the physical conditions prevailing in the expanding gas, we have modeled the first 30 rotational levels of both ortho-- and para--H$_2$O  using the same nonlocal code for lines and dust continuum as in the analysis of H$_2$O toward Sgr~B2 \\citep{cer06}. The code has been described elsewhere \\citep{gon93} and has been recently improved to take into  account a more sophisticated description of the dust emissivity and radiative transfer (Daniel et al. 2006, in prep.). The dust continuum emission has a crucial role in the radiative excitation of light species such as H$_2$O or OH with transitions in the far--IR and has to be correctly taken into account. In our model, level populations are consistently computed in statistical equilibrium  considering collisional excitation and pumping by line and continuum photons. Collisional rates were scaled from those of H$_2$O--He collisions \\citep{gre93}. A three--component model resembling the \\textit{hot core}, the  \\textit{plateau} and the \\textit{ridge} is adopted. In this work we assume that most of the  far--IR water lines arise only from the \\textit{plateau}. Of course, a minor contribution in the lowest excitation water lines may come from the more extended regions (producing narrow line emission) that we don't model here. A low water abundance of (1--8)$\\times$10$^{-8}$ has been estimated in these regions from SWAS and ODIN observations (\\cite{sne00}; \\cite{olo03}) and thus, the expected contribution to far--IR H$_2$O lines is small. In order to have the closest view of the IR radiation field seen by water vapor we have simultaneously tried to reproduce the full SWS/LWS/ISO  continuum emission between $\\sim$10 and $\\sim$197~$\\mu$m.  We find that the continuum level for $\\lambda >$80~$\\mu$m is well reproduced  by considering an inner 10$''$ IR continuum source (the \\textit{hot core})  simulated by a grey--body, with a color temperature of 200--250K, which is optically thick in the far--IR ($\\tau_{\\lambda}=10\\times 150/\\lambda$). The dust continuum emission from this  \\textit{hot core} is attenuated by the surrounding components that we now describe. Taking into account the spatial extent of the 183 and 325~GHz water emission observed at higher spatial resolution \\citep{cer90,cer94,cer99b}, we model the \\textit{plateau} as a $\\sim$40$''$ diameter unresolved shell expanding at 25$\\pm$5~km~s$^{-1}$, where water molecules and dust grains coexist (T$_d$=T$_k$ is assumed). The radiative effects caused by an additional, only dust, 5$'$ diameter component (the \\textit{ridge}), with $T_d\\simeq$25~K and $n$(H$_2$)$\\simeq$10$^5$~cm$^{-3}$,  have been also included to fully match the SWS/LWS continuum level.  \\begin{figure} [t] % \\centering \\includegraphics[angle=0,width=8.5cm]{f2.eps} \\caption{Radiative transfer model results  discussed in the text. Selected synthetic line profiles (continuous curves) are shown over the observed ISO/LWS/FP detections (histograms). Model parameters are $\\chi$(H$_2^{16}$O)=3$\\times$10$^{-5}$, $n(H_2)$=3.5$\\times$10$^{5}$~cm$^{-3}$ and T$_k$=80~K (blue) and $\\chi$(H$_2^{16}$O)=2$\\times$10$^{-5}$, $n(H_2)$=2.5$\\times$10$^{5}$~cm$^{-3}$ and T$_k$=100~K (red). The ordinate corresponds to the continuum normalized flux and the abscissa to the velocity (in~km~s$^{-1}$). Synthetic line profiles have been convolved with a gaussian beam of 80$''$  and with a spectral resolution characterized by a Lorentzian with a width of 33~km~s$^{-1}$.} \\label{fig-models} \\end{figure}%   Dust absorption coefficients are computed from tabulated optical grain properties from  Draine \\& Lee (1984), both in the \\textit{plateau} and in the \\textit{ridge}. Due to the large optical depth of H$_2^{16}$O lines in the \\textit{plateau}, we have first modeled the H$_2^{18}$O lines to obtain a more accurate estimate  of the water abundance with less opacity constraint and faster level population convergence. A $^{16}$O/$^{18}$O abundance ratio of 500 has been adopted. After a reasonable fit to the H$_{2}^{18}$O lines, we iteratively search for the best physical conditions that enable to simultaneously  reproduce the H$_{2}^{16}$O and H$_{2}^{18}$O spectra. Fig. \\ref{fig-models} shows two representative models that qualitatively and quantitatively reproduce the majority of observed line  profiles and intensities. The best models are found for a \\textit{plateau} temperature, density and water vapor abundance of $T_k$=80--100~K, $n$(H$_2$)=2.5--3.5$\\times$10$^5$~cm$^{-3}$ and  $\\chi$(H$_2^{16}$O)=2--3$\\times$10$^{-5}$ respectively. The physical conditions inferred from water lines in the \\textit{plateau} are similar to those obtained from OH lines \\citep{goi06}, which suggests that far--IR lines from both species trace the same expanding gas. The derived water vapor abundance, $\\gtrsim$2$\\times$10$^{-5}$, is obviously an averaged value over the large LWS/ISO beam and  probably indicates that water can be locally more abundant, e.g. in the warmer interaction surfaces where the expanding gas  shocks the dense ambient material. Larger angular resolution is needed to resolve the  relative continuum and H$_2$O spatial  distribution over this complex region. Nevertheless, the inferred gas temperatures, $\\sim$100~K, are significantly below the gas temperature ($\\sim$300~K) required to activate  the gas--phase  neutral--neutral  reactions converting most of the available oxygen into water (abundances larger than $\\sim$10$^{-4}$ are then predicted). Taken into account the short dynamical timescale of the outflow, $\\sim$1000~years, water vapor could have been formed by these neutral--neutral reactions if the gas in the \\textit{plateau} was previously (or is locally) heated to $\\geq$500~K, e.g. by a C--shock passage \\citep{ber98}. However, if the \\textit{plateau} gas temperature has reached a maximum temperature of only $\\sim$100~K, other formation mechanisms are required to explain the observed water vapor abundances. In particular, the temperature inferred from ISO observations is similar to the water--ice evaporation temperature, and therefore,  \\textit{in--situ} evaporation of  water--ice grain mantles formed in an evolutionary stage  prior to the onset of the outflow(s) could now dominate the  water vapor formation in the \\textit{plateau}. Finally, the observed H$_2$O could have been also  formed in the innermost and warmest regions (e.g. the hot core) and had been swept up by the outflow.   To conclude, despite the large number of detected far--IR excited water lines, neither  high gas temperature nor high density conditions are required to populate the higher--energy water rotational levels. Radiative pumping due to the strong IR radiation field from the cloud is enough to populate these levels, at least in the average picture given by ISO observations. Future observations with  \\textit{Herschel} will allow one to map  many far--IR water resolved lines with improved angular resolution. It will then be possible to study in great detail the role of water in star forming regions."
    },
    "0608/astro-ph0608100_arXiv.txt": {
        "abstract": "With the ASTE telescope, we have made observations of three low-mass protostellar envelopes around L483, B335, and L723 in the submillimeter CS ($J$=7--6) and HCN ($J$=4--3) lines. We detected both the CS and HCN lines toward all the targets, and the typical CS intensity ($\\sim$ 1.0 K in T$_{B}$) is twice higher than that of the HCN line. Mapping observations of L483 in these lines have shown that the submillimeter emissions in the low-mass protostellar envelope are resolved, exhibit a western extension from the central protostar, and that the deconvolved size is $\\sim$ 5500 AU $\\times$ 3700 AU (P.A. = 78$^{\\circ}$) in the HCN emission. The extent of the submillimeter emissions in L483 implies the presence of higher-temperature ($\\gtrsim$ 40 K) gas at 4000 AU away from the central protostar, which suggests that we need to take 2-dimensional radiative transfer models with a flattened disklike envelope and bipolar cavity into account to explain the temperature structure inside the low-mass protostellar envelope. The position-velocity diagrams of these submillimeter lines in L483 and B335 exhibit different velocity gradients from those found in the previous millimeter observations. In particular, along the axis of the associated molecular outflow the sense of the velocity gradient traced by the submillimeter lines is opposite to that of the millimeter observations or the associated molecular outflow, both in L483 and B335. We suggest that expanding gas motions at the surface of the flattened disklike envelope around the protostar, which is irradiated from the central star directly, are the origin of the observed submillimeter velocity structure. ",
        "introduction": "Previous millimeter molecular-line observations have found 3000 -- 10000-AU-scale molecular envelopes with both rotation and infalling gas motion around low-mass protostars (\\cite{oh96b,oh97a,oh97b,ta03b}). These millimeter observations have probed structures and kinematics of molecular gas with a temperature of $\\sim$ 10 K and a density of $\\sim$ 10$^{4-5}$ cm$^{-3}$ in the envelopes (\\cite{tak00,sai01}). However, structures and kinematics of higher-temperature ($\\gtrsim$ 40 K) or higher-density gas ($\\gtrsim$ 6 $\\times$ 10$^{6}$ cm$^{-3}$) in those envelopes, which are presumably related to the region in the vicinity of the protostar and disk system (\\cite{cec00,oso03}), are poorly known. This is mainly because millimeter observations can not trace warm ($\\gtrsim$ 40 K) and dense ($\\gtrsim$ 6 $\\times$ 10$^{6}$ cm$^{-3}$) regions of envelopes due to the contamination from extended low-density ($\\sim$ 10$^{4-5}$ cm$^{-3}$) and cold ($\\sim$ 10 K) gas. Submillimeter observations, which can trace warm ($\\gtrsim$ 40 K) and dense ($\\gtrsim$ 6 $\\times$ 10$^{6}$ cm$^{-3}$) gas selectively, are required to investigate such regions and to study the kinematics.  For these purposes, we have made Submillimeter Array (SMA) \\footnote{The Submillimeter Array (SMA) is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics, and is funded by the Smithsonian Institution and the Academia Sinica.} observations of low-mass protostellar envelopes around L1551 IRS5 and IRAS 16293-2422 in the CS ($J$ = 7--6; 342.9 GHz) and HCN ($J$ = 4--3; 354.5 GHz) lines, respectively (\\cite{tak04,tak06}). These submillimeter molecular lines trace gas temperatures above $\\gtrsim$ 40 (K) and densities above $\\gtrsim$ 6 $\\times$ 10$^{6}$ cm$^{-3}$, and hence are appropriate tracers for our study. Both in L1551 IRS5 and IRAS 16293-2422, we have detected a compact ($\\sim$ 500 AU) disklike structure associated with the protostars in the submillimeter lines. However, our SMA data recover only $\\sim$ 11$\\%$ and $\\sim$ 30$\\%$ of the total flux observed with JCMT in L1551 IRS5 and IRAS 16293-2422, respectively, which implies that there is also an extended ($>$ 2000 AU) submillimeter component in the low-mass protostellar envelopes. In fact, our JCMT mapping observations of IRAS 16293-2422 in the HCN (4--3) emission have revealed an extended ($\\sim$ 3000 AU) envelope structure as well as the compact ($\\sim$ 500 AU) disklike structure which is embedded in the extended envelope. These results suggest that there is extended ($>$ 2000 AU) warm ($\\gtrsim$ 40 K) or dense ($\\gtrsim$ 6 $\\times$ 10$^{6}$ cm$^{-3}$) gas as well as the compact disklike structure in these low-mass protostellar envelopes. It seems to be difficult to make gas temperatures high enough strictly via heating from the central stars embedded in the envelopes (e.g., \\cite{lay94}), and it is a puzzle why such warm ($\\gtrsim$ 40 K) or dense ($\\gtrsim$ 6 $\\times$ 10$^{6}$ cm$^{-3}$) gas can be so extended in the low-mass protostellar envelopes. These results suggest that we need to re-consider the origin of submillimeter emissions in low-mass protostellar envelopes.  In order to clarify the extent, kinematics, and the origin of submillimeter emissions in low-mass protostellar envelopes unambiguously, we have initiated mapping observations of low-mass protostellar envelopes in the HCN (4--3) and CS (7--6) lines with ASTE (Atacama Submillimeter Telescope Experiment), a new and powerful submillimeter single-dish telescope at the Atacama Site in Chile. In this paper, we report ASTE results of HCN (4--3) and CS (7--6) observations of three nearby ($\\leq$ 300 pc) low-mass protostellar sources, that is, L483, B335, and L723. L483 is a bright ($L_{bol}$ $\\sim$ 14 $L_{\\odot}$) Class 0 object with an East (red) - West (blue) molecular outflow at a distance of 200 pc (\\cite{lad91,hat99,taf00}). Millimeter CS (2--1) and H$_{2}$CO (2$_{12}$--1$_{11}$, 3$_{12}$--2$_{11}$) spectra toward the protostar of L483 exhibit the \"infalling asymmetry\" with the stronger blueshifted emission and the absorption dip, and this source is considered to be one of the protostellar candidates with an infalling gas motion in the surrounding envelope (\\cite{mye95,mar97}). Mapping observations of L483 in centimeter and millimeter lines such as NH$_{3}$ (1,1; 2,2), N$_{2}$H$^{+}$ (1--0) and CH$_{3}$OH (2--1) have revealed that the protostellar envelope is extended along the east-west direction, and that the velocity structure is similar to that of the associated molecular outflow, suggesting that the dense gas in the outer portion of the envelope is being swept-up by the associated outflow (\\cite{gre97,ful00,par00,taf00}). In L483, detailed studies of chemical differentiation in the envelope have also been performed (\\cite{jor02,jor04}). B335 is a close ($d$ $\\sim$ 250 pc) Bok globule associated with a far-infrared source (IRAS 19347+0727; $L_{bol}$ $\\sim$ 3 $L_{\\odot}$), and its spectral energy distribution classifies this source as a Class 0 object (\\cite{dav87,bar95}). This source has an East (blue) - West (red) molecular outflow (\\cite{hir92,cha93}). Aperture synthesis observations of the millimeter H$^{13}$CO$^{+}$ emission (\\cite{sai99}), as well as the shape of the central spectra in the CS (2--1 ,3--2, 5--4) lines (\\cite{cho95,zho95,eva05}), have shown that there is an infalling gas motion in the surrounding envelope around B335, although the CS results may be severely affected by the contamination from the molecular outflow (\\cite{wil00}). Detailed millimeter continuum studies in the context of the circumstellar disk and envelope have also been performed in B335 (\\cite{ha03a,ha03b}). L723 is another Class 0 candidate with a luminosity of $\\sim$ 3$L_{\\odot}$ at a distance of 300 pc (\\cite{dav87}). In L723, there are two radio sources, VLA1 and VLA2 (\\cite{ang91}), where only VLA2 is associated with the surrounding molecular envelope as well as the quadrupolar CO outflow, suggesting that VLA2 is younger than VLA1 (\\cite{hay91,hir98}).  We note that only limited single-dish mapping observations of low-mass protostellar envelopes in submillimeter molecular lines have been reported (\\cite{hog98,hog99}), and that most submillimeter single-dish studies of low-mass protostellar envelopes are single-point observations, on the assumption that submillimeter lines in low-mass protostellar envelopes are more compact than the beam size of the single-dish telescopes ($<$ 20$\\arcsec$ $\\sim$ 4000 AU in nearby sources) (\\cite{bla94,mor95,dis95}). Our new high-sensitivity mapping observations of the low-mass protostellar envelopes in the submillimeter lines with ASTE provide us a new insight of the origin of the submillimeter emission in low-mass protostellar envelopes, as will be shown in the present paper. ",
        "conclusions": "\\subsection{Physical Conditions of Molecular Gas traced by the Submillimeter Lines\\\\ in the Low-mass Protostellar Envelopes}  Our ASTE observations have resolved the spatial and/or kinematical structures of the submillimeter HCN and CS emission in L483 and B335. The intensities and extents should trace structures of gas densities, temperatures, and molecular abundances in these envelopes. In order to investigate these physical conditions of the protostellar envelopes traced by the extended submillimeter emissions, we performed statistical equilibrium calculations of the submillimeter lines based on the Large Velocity Gradient (LVG) model (\\cite{gol74,sco74}).  For the model calculations, we employed values of the dipole moment, rotational and centrifugal constants of the HCN and CS molecule listed by Winnewisser et al. (1979). Collisional transition rates of HCN were taken from Green and Thaddeus (1974) and those of CS from Green and Chapman (1978). Rotational energy levels included in our calculations are up to $J$ = 8 (153 K) for HCN and $J$ = 13 (214 K) for CS. These energy levels are high enough to discuss physical conditions of molecular gas in low-mass protostellar envelopes (\\cite{mor95}).  In Figure 10, we show results of our calculations. The HCN and CS line intensities increase monotonically as the gas density ($\\equiv$ $n_{H2}$), molecular abundance per unit velocity gradient ($\\equiv$ $X/dv/dr$), or the gas kinetic temperature ($\\equiv$ $T_{k}$) increases. Since we have only one transition of HCN and CS, we cannot fully constrain the physical conditions traced by these submillimeter lines, but we can still make some arguments as follows. In L483, the HCN and CS abundance is estimated to be $X_{HCN}$ = 2.0 $\\times$ 10$^{-9}$ and $X_{CS}$ = 6.8 $\\times$ 10$^{-10}$, respectively (\\cite{joe04}). The observed HCN and CS line width ($\\equiv$ $dv$) at the center of L483 is $\\sim$ 2.1 km s$^{-1}$ and 1.4 km s$^{-1}$ ($\\S$3.1) and the deconvolved size ($\\equiv$ $dr$) along the major axis of the HCN and CS emission is 0.027 pc and 0.011 pc ($\\S$3.2). Hence, the $X/dv/dr$ value can be derived to be $\\sim$ 2.6 $\\times$ 10$^{-11}$ for HCN and $\\sim$ 5.6 $\\times$ 10$^{-12}$ for CS. Then, at the 20$\\arcsec$ = 4000 AU west from the central protostar of L483 where the HCN and CS lines are detected to be 0.51 (K) and 0.40 (K), respectively (Table 2), the value of $n_{H2}$ must be higher than $\\sim$ 6 $\\times$ 10$^{6}$ cm$^{-3}$ if $T_{k}$ = 10 (K), or the value of $T_{k}$ must be higher than $\\sim$ 40 (K) if $n_{H2}$ = 6 $\\times$ 10$^{5}$ cm$^{-3}$. At the central protostellar positions of L483 and B335, the same $X/dv/dr$ values provide the estimate of $n_{H2}$ $>$ 10$^{7}$ cm$^{-3}$ if $T_{k}$ $\\sim$ 10 (K), or the estimate of $T_{k}$ $>$ 40 (K) if $n_{H2}$ $\\sim$ 10$^{6}$ cm$^{-3}$, averaged over the 22$\\arcsec$ ASTE beam. These arguments imply that high-density ($\\gtrsim$ 6 $\\times$ 10$^{6}$ cm$^{-3}$) or high-temperature ($\\gtrsim$ 40 K) molecular gas could be more extended than $\\sim$ 2000 AU in radius, particularly evident at the 4000 AU west from the protostar of L483.  The presence of such an extended high-density or high-temperature gas in low-mass protostellar envelopes is also supported by other recent submillimeter-line observations of low-mass protostellar envelopes. Takakuwa et al. (2006) have made single-dish and interferometric observations of IRAS 16293-2422 in the HCN (4--3) line with the JCMT and the SMA and have found that there is an extended ($>$ 3000 AU) HCN emission plus compact ($\\sim$ 500 AU) disklike structure associated with the protostar. CS (7--6) observations of L1551 IRS5 with the SMA imply that the submillimeter CS emission is more extended than $\\sim$ 1500 AU (\\cite{tak04}). However, it is still less clear why these submillimeter molecular lines which trace high-density or high-temperature gas can be so extended. Masunaga et al. (1998, 2000) have constructed a one-dimensional radiation hydrodynamic model in the formation of an 1 M$_{\\odot}$ protostar. Their model suggests that during the main accretion phase the $T_{k}$ and $n_{H2}$ values at the radius of 4000 AU (i.e., L483 20$\\arcsec$ west) are $\\sim$ 15 K and $\\sim$ 2 $\\times$ 10$^{5}$ cm$^{-3}$, and at the radius of 2000 AU $\\sim$ 20 (K) and $\\sim$ 10$^{6}$ cm$^{-3}$, respectively. These theoretically-predicted values seem to be slightly lower to explain the observed submillimeter emissions, although the limited spatial resolution as well as the only single transitions of our ASTE data prevents us from making direct comparison between the observational and theoretical results. On the other hand, the western extension of the submillimeter emission in L483 seems to be aligned with the western component of the associated outflow (\\cite{taf00,jor04}), and there might be another heating mechanism such as shocks due to the interaction between the outflow and the ambient gas (\\cite{ume92,hir01}).  As described in $\\S$3.2 the velocity structures traced by the submillimeter lines are different from those traced by millimeter lines both in L483 and B335, which suggests that the submillimeter emissions have a different origin and that the spherically-symmetric radiation hydrodynamic model is not adequate to fully explain our observational results. We will discuss on these different velocity structures in the next section.  \\begin{figure} \\begin{center} \\FigureFile(150mm,150mm){figure10.eps} \\end{center} \\caption{Results of our LVG model calculations. Brightness temperatures of the CS (7--6) (left) and HCN (4--3) (right) lines are plotted in contours as a function of the gas density ($\\equiv$ $n_{H2}$) and the molecular abundance per unit velocity gradient ($\\equiv$ $X/dv/dr$) at each gas kinetic temperature ($T_{k}$ = 10, 20, 40, 60, and 100 K as shown in the Figure). Contour levels are shown in the Figure, and the lowest five contours correspond to the observed line intensities at the representative positions listed in Table 2.}\\label{fig:f10} \\end{figure}   \\subsection{Different Velocity Structures traced by the Submillimeter Lines\\\\ in the Low-mass Protostellar Envelopes}  Our submillimeter-line observations of the low-mass protostellar envelopes have revealed that along the axis of the associated molecular outflows the sense of the velocity gradient traced by the submillimeter lines is opposite to that traced by the previous millimeter observations or the molecular outflows themselves ($\\S$3.2). In L483, millimeter CS (2--1) and N$_{2}$H$^{+}$ (1--0) lines as well as the associated CO outflow show redshifted emissions at the east of the protostar and blueshifted emissions at the west (\\cite{hat99,jor04}), while the submillimeter CS (7--6) and HCN (4--3) lines exhibit blueshifted emissions at the east and redshifted emissions at the west of the protostar. In B335, the associated CO outflow and the millimeter C$^{18}$O (1--0) and H$^{13}$CO$^{+}$ (1--0) lines show blueshifted emissions at the east of the protostar and redshifted emissions at the west (\\cite{cha93,sai99}), whereas the submillimeter CS and HCN lines show an opposite velocity gradient along the same axis. In this subsection, the origin of these different velocity structures traced by the submillimeter lines is discussed.  One possible interpretation of the different velocity structure in the submillimeter lines is that the submillimeter lines trace the other side of the conical shell of the outflow than that traced by the millimeter lines incidentally, whose axis is close to the plane of the sky. The inclination angle of the associated outflow from the plan of the sky in L483 is estimated to be $\\sim$ 45$^{\\circ}$ and there appears a weak redshifted outflow component at the western side (\\cite{par00,taf00}), whereas the inclination angle in B335 is derived to be $\\sim$ 10$^{\\circ}$ and there are also redshifted and blueshifted outflow components at the east and west of the protostar, respectively (\\cite{hir92}). Although the main outflow component of the blueshifted and redshifted gas locates at the west and east of the central source of L483 (\\cite{hat99}) and and at the east and west of B335 (\\cite{cha93}), it is possible that the observed submillimeter emission trace those other outflow components incidentally. In such a case, the velocity gradient in the submillimeter emission shows the opposite sense to that of the main velocity gradient traced by the millimeter lines. The observed elongation of the submillimeter emission along the outflow axis in L483 favors this interpretation, and the higher-temperature ($\\gtrsim$ 40 K) gas at the western extension in L483 discussed in the previous subsection may reflect the shock interaction between the outflow and the ambient dense gas (\\cite{ume92,hir01}).   Another possible and more interesting interpretation of the observed velocity structures in the submillimeter lines is an expanding gas motion which is perpendicular to the outflow axis, as shown in Figure 11. This figure demonstrates that the submillimeter lines trace molecular gas at the surface of the flattened disklike envelope, which is being stripped away by some mechanism such as the stellar wind, while the millimeter lines trace infalling gas at the midplane of the envelope. Here, on the side of the blueshifted outflow the expanding gas is observed as a redshifted component and on the side of the redshifted outflow as a blueshifted component, while the infalling gas components show the same velocity gradient as that of the outflow. Hence, this configuration can explain the different velocity gradient observed in the submillimeter lines from that of the outflow components or infalling gas traced by the millimeter lines.  Similar velocity structures along the perpendicular direction to the polar outflow have also been found in near-infrared observations of young stellar objects. Pyo et al. (2005) have observed lower-velocity [Fe II] winds with a wide opening angle of $\\sim$ 100$^{\\circ}$ in L1551 IRS5, as well as well-collimated higher-velocity jets along the polar axis. Comparable results were also obtained for the young T Tauri star DG Tau (\\cite{pyo03}). They suggest that such a wind with a wide opening angle will be effective in sweeping up envelope material from the close vicinity of its driving source. The observed velocity structures in the submillimeter lines may trace such a gas motion in the envelope, although the spatial scale in the submillimeter results ($>$ 2000 AU) is much larger than that of the infrared observations ($\\sim$ 100 - 400 AU). In L483, the line core emission in the HCO$^{+}$ (3--2) line does show the same velocity gradient as our submillimeter results, while the HCO$^{+}$ (3--2) line wing emission traces the bipolar outflow itself (see Figure 15 in \\cite{gre97}). At the surface of the flattened disklike envelopes molecular gas is irradiated by the stellar photons and hence the gas temperature is likely to be higher than that at the midplane. Then, the warmer surface regions are selectively traced by the submillimeter observations whereas the colder, infalling regions at the midplane are traced by the millimeter lines alone. In fact, Spaans et al. (1995) proposed that the narrow $^{12}$CO and $^{13}$CO $J$=6--5 emission observed toward many low-mass young stellar objects is produced at the surface of the circumstellar envelope surrounding the cavity evacuated by the bipolar outflow, heated by the radiation field generated in the inner part of the accretion disk. Recent two-dimensional radiative transfer models in protostellar envelopes have also revealed that at the surface of the flattened disklike envelope the gas temperature becomes higher than that predicted from spherically symmetric models (\\cite{nak03,whi03}). The observed submillimeter HCN and CS emission may have such an origin.  The interpretation of the expanding flattened envelope, however, has one weak point, that is, in L483 the observed submillimeter emissions are elongated $along$ the outflow direction. If the submillimeter lines trace expanding gas motion at the surface of the disklike envelope, the elongation should be perpendicular to the axis of the associated outflow. In order to compensate for this contradiction, non polar-symmetric geometry and/or the flared shape of the envelope must be incorporated. With the present data, we cannot determine which scenario, the incidental projection effect or the expanding flattened envelope with the flared and/or non polar-symmetric shape, is more plausible.  On the other hand, we have also found an opposite velocity gradient in the submillimeter lines to that in the millimeter lines across the outflow axis in B335. The origin of the difference is less clear, but the expanding gas traced by the submillimeter lines may have a different rotational motion. In any sense, our new submillimeter mapping observations have presented us an implication that submillimeter lines trace different kinematics from that traced by millimeter lines, and that there is another velocity structure which previous millimeter observations could not find in low-mass protostellar envelopes.   \\begin{figure} \\begin{center} \\FigureFile(100mm,100mm){figure11.eps} \\end{center} \\caption{A schematic diagram to interpret the observed velocity gradient of the submillimeter emission in the low-mass protostellar envelopes. Observers' line of sight is normal to the paper. The bold black ellipse represents a flattened disklike envelope. This disklike envelope is tilted, and the near side is to the left. The associated blueshifted and redshifted molecular outflows perpendicular to the disklike envelope are shown in light blue and red colors, respectively, and the blueshifted outflow locates at the lefthand, near side. At the midplane of the disklike envelope there is an infalling gas motion as described in the black arrows. At the surface of the disklike envelope there is an expanding gas motion as shown in red and blue arrows, which is traced by the present submillimeter observations. Here, the left-half portion is redshifted while the other half blueshifted, and this velocity structure has an opposite sense to that of the outflow or the infalling gas motion. The configuration in this diagram represents the case of B335. For the case of L483 one can just flip this configuration horizontally.}\\label{fig:f11} \\end{figure}"
    },
    "0608/astro-ph0608270_arXiv.txt": {
        "abstract": "Identifying binaries among runaway O- and B-type stars offers valuable insight into the evolution of open clusters and close binary stars.  Here we present a spectroscopic investigation of 12 known or suspected binaries among field and runaway OB stars.  We find new orbital solutions for five single-lined spectroscopic binaries (HD 1976, HD 14633, HD 15137, HD 37737, and HD 52533), and we classify two stars thought to be binaries (HD 30614 and HD 188001) as single stars.  In addition, we reinvestigate their runaway status using our new radial velocity data with the UCAC2 proper motion catalogs.  Seven stars in our study appear to have been ejected from their birthplaces, and at least three of these runaways are spectroscopic binaries and are of great interest for future study. ",
        "introduction": "Nearly all O- and B-type stars are believed to form in open clusters and stellar associations, but a handful of OB stars are observed at high galactic latitudes and with large peculiar space velocities, suggesting that they have been ejected from the cluster of their birth.  There are two accepted mechanisms to explain the origin of these runaway O- and B-type stars.  In one scenario, close multi-body interactions in a dense cluster environment cause one or more stars to be scattered out of the region \\citep*{poveda1967}.  In rare cases, the ejected stars may be bound as a binary pair; \\citet{leonard1990} use $N$-body simulations of open clusters to show that a binary frequency of about 10\\% is expected from dynamical ejections.  An alternative mechanism involves a supernova explosion within a close binary, ejecting the secondary due to the conservation of momentum \\citep{zwicky1957, blaauw1961}.  The resulting neutron star may remain bound to the secondary if not enough mass is lost during the explosion.  \\citet{portegieszwart2000} predicts a binary fraction of $20-40$\\% among runaways which are ejected by such a binary supernova scenario.  The observed fraction of binaries among runaways seems consistent with either scenario ($5-26$\\%; \\citealt{mason1998}).  Observational evidence suggests that both dynamical interactions and supernovae in close binaries do produce some runaway systems.  The microquasar LS 5039 was likely ejected from the Milky Way disk about 1.1 Myr ago by the supernova explosion of the neutron star progenitor \\citep{ribo2002, mcswain2004, casares2005}.  \\citet{berger2001} find that about 7\\% of Be stars have high peculiar space velocities, likely due to energy received during the supernova of a close binary companion \\citep{mcswain2005}. On the other hand, \\citet*{gualandris2004} show that a four-body encounter 2.5 Myr ago in the Trapezium cluster resulted in three ejected systems: two single stars and one binary.  The single stars both have very large space velocities and are classified as runaways, although the ejected binary has a lower space velocity due to its higher total mass.  In a study of 44 trapezium-type clusters, \\citet*{allen2006} find 4 stars with large transverse velocities that were likely ejected by dynamical interactions.  Which process dominates the production of runaway stars? \\citet*{hoogerwerf2001} show that massive runaways tend to have enriched helium abundances and higher rotational velocities, while fewer non-runaways exhibit these traits.  While rotational mixing could explain these results, it is doubtful that mixing occurs preferentially among runaways.  It is more likely that prior mass transfer of CNO-processed material from the evolved donors has altered the surface abundances and added angular momentum to the mass gainers.  The high binding energy of close massive binaries makes it unlikely they will be disrupted by dynamical interactions, but they could be separated by a supernova.  Thus according to \\citet{hoogerwerf2001}, the evidence favors the supernova ejection scenario over dynamical cluster interactions.  However, pulsar searches by \\citet{philp1996} and \\citet*{sayer1996} have not identified any neutron stars among runaway OB stars, and further investigation of their properties is needed.  The production of runaway OB binaries is expected to be rare, but these systems can offer key insights into the evolution of close binary stars and open clusters.  If massive stars form in dense regions through the mergers of low mass protostellar cores, the expected multi-body interactions will result in a large number of high velocity runaways \\citep{bally2005}.  Ejecting a star dynamically from a young cluster will reduce the total energy imparted to the cluster via stellar winds, radiation, and supernova ejecta, whereas a proto-supernova system will remain in the cluster for a longer time and contribute more energy into its environment \\citep{dray2005}.  Finally, the relative importance of the supernova ejection scenario places limits on the mass needed to form black holes and neutron stars \\citep{dray2005}.  In this work, we examine the binary status of 12 field and runaway OB stars.  Through a radial velocity study of these stars, we determine the orbital solutions of five single-lined spectroscopic binaries (SB1s).  One is a possible triple star, and we also find two suspected SB1 systems. By fitting the spectral energy distributions of these stars, we measure their reddening and angular sizes to constrain their distances.  Finally, our radial velocities combined with proper motions allow us to measure their peculiar space velocities and reinvestigate the runaway status of these targets.  Our results reveal that only 7 of the 12 stars have high space velocities, while five stars may not be runaways.  Of the runaways, two are SB1s and a third is a suspected SB1.  We will investigate their cluster ejection scenarios in a forthcoming paper. ",
        "conclusions": ""
    },
    "0608/astro-ph0608046_arXiv.txt": {
        "abstract": "% The problems with explaining the Small Ionized and Neutral Structures (SINS) appealing to turbulence stem from inefficiency of the Kolmogorov cascade in creating large fluctuations at sufficiently small scales. However, other types of cascades are possible. When magnetic turbulence in a fluid with viscosity that is much larger than resistivity gets to a viscous damping scale, the turbulence does not vanish. Instead, it gets into a different new regime. Viscosity-damped turbulence produces fluctuations on the small scales. Magnetic fields sheared by turbulent motions by eddies not damped by turbulence create small scale filaments that are confined by the external plasma pressure. This creates small scale density fluctuations. In addition, extended current sheets create even stronger density gradients that accompany field reversals in the plane perpendicular to mean magnetic field. Those can be responsible for the SINS formation. This scenario is applicable to partially ionized gas. More studies of reconnection in the viscosity dominated regime are necessary to understand better the extend to which the magnetic reversals can compress the gas. ",
        "introduction": "Turbulence can be viewed as a cascade of energy from a large injection energy scale to dissipation at a smaller scale. The latter is being established by equating the rate of turbulent energy transfer to the rate of energy damping arising, for instance, from viscosity. Naively, one does not  expect to see any turbulent phenomena below such a scale.  Such reasoning may not be true in the presence of magnetic field, however. Consider magnetized fluid with viscosity $\\nu$ much larger than magnetic diffusivity $\\eta$, which is the case of a high magnetic Prantl number $Pr$ fluid. The partially ionized gas can serve as an example of such a fluid up to the scales of ion-neutral decoupling (see a more rigorous treatment in Lazarian, Vishniac \\& Cho 2004, henceforth LVC04). Fully ionized plasma is a more controversial example. For instance, It is well known that for plasma the diffusivities are different along and perpendicular to magnetic field lines. Therefore, the plasma the Prandtl number is huge if we use the parallel diffusivity $\\nu_{\\|}\\gg \\nu_{\\bot}$. A treatment of the fully ionized plasma as a high Prandtl number medium is advocated in Schechochihin et al (2004, henceforth SCTMM).  The turbulent cascade in a fluid with isotropic $\\eta$ proceeds up to a scale at which the cascading rate, which for the Kolmogorov turbulence, i.e. $v_l\\sim l^{1/3}$, is determined by the eddy turnover rate $v_l/l$ gets equal to the damping rate $\\eta/l^2$. Assuming that the energy is injected at the scale $L$ and the injection velocity is $V_L$, the damping scale $l_c$ is $L Re^{-3/4}$, where $Re$ is the Reylonds number $L V_L/\\eta$.  However, it is evident that magnetic fields at the scale $l_c$ at which hydrodynamic cascade would stop are still sheered by eddies at the larger scales. This should result in creating magnetic structures at scales $\\ll l_c$ (LVC04). Note, that magnetic field in this regime is not a passive scalar, but important dynamically.   Cho, Lazarian \\& Vishniac (2002) reported a new viscosity-damped regime of turbulence using incompressible MHD simulations (see an example of more recent simulations in Fig.~1). In this regime, unlike hydro turbulence, motions, indeed, do not stop at the viscosity damping scale, but magnetic fluctuations protrude to smaller scales. Interestingly enough, these magnetic fluctuations induce small amplitude velocity fluctuations at scales $\\ll l_c$.  Cho\\& Lazarian (2003) confirmed these results with compressible simulations and speculated that these small scale magnetic fluctuations can compress ambient gas to produce SINS (see also Fig. 1b). According to the model in LVC04, while the spectrum of averaged over volume magnetic fluctuations scales as $E(k)\\sim k^{-1}$, the pressure within intermittent magnetic structures increases with the decrease of the scale as $(\\delta \\hat{b})^2_k\\sim k$, while the filling factor $\\phi_k \\sim k^{-1}$, the latter being consistent with numerical simulations. The pressure of gas confining the magnetic filaments should increase accordingly, resulting in fluctuations of density increasing with the decrease of the scale. The fact that the emerging structures are filamentary allows occasional picks in the observed column densities, which could correspond to the Heiles (1987) model of SINS. \\begin{figure} \\includegraphics[width=3.6cm]{ff1.ps} \\hfill \\includegraphics[width=0.67cm]{fig1b.eps} \\hspace{0.07cm}\\hfill \\includegraphics[width=5.0cm]{fig2.eps} \\caption{ {\\it Left}: Filaments of density created by magnetic compression in the slice of data cube of the viscosity-damped regime of MHD turbulence. {\\it Right}: Spectra of density and magnetic field are similar, while velocity is damped. The resistive scale in this regime is not $L/Rm$ but $L {Rm}^{-1/2}$.(from Beresnyak \\& Lazarian, in preparation).} \\label{visc} \\end{figure}  In Fig.~1 we show the results of our high resolution compressible MHD simulations that exhibit strong density fluctuations at the scales below the one at which hydrodynamic turbulence would be damped. The testing of the LVC04 model is important by itself, even without a possible connection to SINS. According to  that model, the viscosity-damped regime is ubiquitous in turbulent partially ionized gas. Some of its discussed consequences, such as an intermittent resumption of the turbulence the fluid of ions as magnetic fluctuations reach the ion-neutral decoupling scale, are important for radio scintillations. ",
        "conclusions": "Recently GS06 attempted to explain the extreme radio-wave scattering in the direction of the Galactic Center assuming plasma can be described with high $Pr$ number and the folded fields in the spirit of SCTMM are present along the line of sight. They do not appeal directly to the SP reconnection process, but the arguments about the current sheets there implicitly assume that the SP-type reconnection. As the result, their estimates of the variation of the column density can be obtained from ours assuming that $L=l_c$ for the folded fields model. Compared to the GS06 the model in the paper is applicable to the situations when strong mean magnetic field is present. If attempting to explain the extreme scattering events with viscosity-damped turbulence model, one could use the corresponding estimates in GS06, modifying the column density fluctuation as discussed above. In the highly viscous fluid the values of $L$ and $l_c$ do not differ dramatically, anyhow.  GS06 and the model above when applied to plasmas share the same set of problems. As GS06 admitted, the SCTMM folded field structure could unwind rapidly due to the low kinematic viscosity in directions perpendicular to the magnetic field. In view of this, the formation of the SINS in partially ionized gas (e.g. HI) could be seen as a safer bet.  Another potential difficulty shared by both models is related to using SP model of current sheets for both the folded fields and the viscosity damped turbulence. Usually, one would expect that the reconnection happens faster, opening up the reconnection layers (see Shay et al. 2001 for collisionless reconnection and Lazarian \\& Vishniac 1999 for stochastic reconnection). However, the conditions for the field wandering and the outflows within both the viscosity damped turbulence at scales less than $l_c$ (see LVC04) and the folded field dynamo (see SLTMM) differ from the usual formulation of the reconnection problem. Further research should clarify what what is going on in such situations. Potentially, compression in current sheets looks as an attractive solution of the SINS phenomenon."
    },
    "0608/hep-ph0608115_arXiv.txt": {
        "abstract": "We study the effect of free streaming in a universe with cosmic strings with time-varying tension as well as with constant tension. Although current cosmological observations suggest that fluctuation seeded by cosmic strings cannot be the primary source of cosmic density fluctuation, some contributions from them are still allowed. Since cosmic strings actively produce isocurvature fluctuation, the damping of small scale structure via free streaming by dark matter particles with large velocity dispersion at the epoch of radiation-matter equality is less efficient than that in models with conventional adiabatic fluctuation. We discuss its implications to the constraints on the properties of particles such as massive neutrinos and warm dark matter. ",
        "introduction": "Cosmological observations such as cosmic microwave background (CMB), large scale structure and so on become now very precise and can constrain various cosmological parameters with unprecedented accuracy. In addition to the determination of cosmological parameters such as energy densities of baryons, dark matter, and dark energy, Hubble parameter, the scalar spectral index, reionization optical depth and so on, cosmological observations can also constrain unknowns in particle physics.  For example, neutrino masses can be severely constrained by cosmology \\cite{Spergel:2006hy,Fukugita:2006rm}. It is well known that massive neutrinos can erase small scale inhomogeneities via free streaming since massive neutrinos can have large velocity dispersion at the time of radiation-matter equality. Due to this effect, the matter power spectrum exhibits the suppression at small scales, which can be compared to observations of large scale structure to give the constraint on neutrino masses\\footnote{ The mass of neutrino can also affect the CMB power spectrum through the modification of the structure of the acoustic peaks, which can give a severe constraint on them \\cite{Fukugita:2006rm,Ichikawa:2004zi}.  }.  Another such example is warm dark matter (WDM) particles. WDM particles can also have large velocity dispersion at the epoch of radiation-matter equality like massive neutrinos. WDM scenarios have been extensively studied in the literature in connection with particle physics and astrophysics.  As for astrophysical aspects, WDM has been discussed, in particular, as a solution of the problem of small scale structure such as the missing satellite problem and the cusp problem \\cite{Colin:2000dn,Bode:2000gq}. From the viewpoint of particle physics, there exist well-motivated candidates for WDM such as a light gravitino \\cite{Pagels:1981ke,Bond:1982uy}, sterile neutrinos \\cite{Dodelson:1993je} and so on. The properties of WDM such as its mass can be constrained by cosmological observations as the same manner as the case with massive neutrinos since WDM particles also erase the small scale fluctuation via the free streaming effect. It should also be mentioned that superweakly interacting massive particles (superWIMPs) \\cite{Feng:2003xh} can also erase the small scale structure as WDM particles, thus this kind of model can also be constrained by cosmological observations by studying the damping of matter power spectrum.  Therefore, probing the damping of small scale structure can be an important test for particle physics.  Constraints on above mentioned particles such as massive neutrinos and candidates for WDM have been well studied in the framework where cosmic density fluctuation is seeded by conventional adiabatic primordial fluctuation which is motivated by inflation driven by a scalar field. However, cosmic density fluctuation can also be produced by cosmic strings. Since fluctuation seeded by cosmic strings is an incoherent actively generated isocurvature one, this kind of fluctuation cannot produce observed structure of the acoustic peaks in the CMB power spectrum but gives rise to fairly broad acoustic peaks.  Thus current observations suggest that cosmic strings cannot be the primary source of density fluctuation today. However, subdominant contribution from fluctuation seeded by cosmic strings is still allowed \\cite{Wyman:2005tu}. Furthermore cosmological scenarios with cosmic strings have been revived for recent years since there have been discussed that cosmic strings can be produced in a wide class of string theory models. In particular, cosmic strings can be formed at the end of brane inflation \\cite{Jones:2002cv,Sarangi:2002yt,Jones:2003da}.  In light of these considerations, it is interesting to study the effect of free streaming in a scenario with cosmic strings. Importantly, since cosmic strings produce fluctuation actively, the erasure of small scale inhomogeneities via free streaming can be avoided to some extent, which results in delayed damping of small scale power \\cite{Brandenberger:1987er,Brandenberger:1987kf,Bertschinger:1987,Albrecht:1992sb}. Hence one may consider that some contribution from fluctuation seeded by cosmic strings can relax the constrains on WDM or neutrino masses since the constraint on the masses mainly come from the effect of free streaming. This is the issue which we are going to consider in this paper.  In fact, the authors of Ref.~\\cite{Brandenberger:2004kc} have discussed the possibilities of relaxing the constraint on neutrino masses in models with cosmic strings with constant tension. There it was shown by a simple analytic estimate that even with the addition of fluctuation from cosmic strings the constraint on neutrino masses cannot be relaxed. In this paper, first of all, we reconfirm this result with more quantitative analysis by numerically calculating matter power spectrum and CMB anisotropy produced by cosmic strings. We also do similar calculations for WDM and judge whether the constraints on WDM can be relaxed or not.  Furthermore, recently a new class of cosmic strings has been considered, which has time-varying tension. When a scalar field constituting a string couples to another oscillating field or a string is realized in a brane configuration, cosmic strings with time-varying tension naturally appear. Cosmological evolution of such a class of cosmic strings has been investigated in Refs.~\\cite{Yamaguchi:2005gp,Ichikawa:2006rw} where it was shown that, after some relaxation time, it goes into the scaling regime like strings with constant tension \\cite{ls,gs} and global monopoles \\cite{gm}. In the scaling regime, the typical length of the cosmic string network grows with the horizon scale. Then, in case that string tension is constant, the ratio of the energy density of infinite strings to that of the background universe is also constant, which generates scale invariant density fluctuations. On the other hand, when the string tension is time-varying, the ratio of the energy density of infinite strings to that of the background universe is not necessarily constant due to the time dependence of the tension. For example, when the tension has the time dependence as $G\\mu \\propto \\tau^{n}$ with $\\tau$ being the conformal time, fluctuations at small scales are enhanced but those at large scales are suppressed for negative values of $n$. Thus the above discussion on the effects of cosmic strings on free streaming may be modified when we consider the case with time-varying tension.  The purpose of this paper is to study to what extent fluctuation from cosmic string with time-varying tension as well as with constant tension can affect to avoid the erasure of small scale inhomogeneities via free streaming by massive neutrinos and WDM. We also discuss the implications of the above mentioned phenomenon to the constraints on the masses of these particles. ",
        "conclusions": "We studied the erasure of small scale structure via free streaming in models with cosmic strings with time-varying tension as well as with constant tension. Because cosmic strings actively produce incoherent isocurvature fluctuation, free streaming effect is less efficient than that in models with the conventional adiabatic ones because of its nature of the fluctuation. Since the damping of small scale power by free streaming is an important probe to constrain the property of particles with large velocity dispersion at the epoch of radiation-matter equality such as massive neutrinos and WDM, above mentioned effects by cosmic strings can have much implications to the constraints on such particles.  We first studied this issue for massive neutrinos. For the case with constant string tension, it has been already discussed some time ago that, even if we add some contribution from isocurvature fluctuation from cosmic strings, we cannot relax the constraint on neutrino masses. In this paper, we explicitly calculate matter power spectrum and CMB anisotropy generated by strings and reach the same conclusion. Furthermore, we have studied this issue for the case with cosmic strings with time-varying tension. We have shown that even if we introduce cosmic strings with time-varying tension, we cannot alleviate the constraint on the mass without conflicting with the constraint from current CMB observations.  This is mainly because adiabatic fluctuation with massive neutrinos damps the matter power spectrum on small scales corresponding to multipoles $\\ell \\sim \\mathcal{O}(100)$ in CMB power spectrum.  We have also discussed the possibilities of relaxing constraints on the masses of WDM by adding fluctuation produced by cosmic strings with time-varying tension and those with constant tension. We explicitly showed that the damping of the matter power spectrum on small scales from adiabatic primordial fluctuation by free streaming caused by WDM can be canceled by adding fluctuation seeded by cosmic strings for both cases, namely with constant and time-varying tension by choosing appropriate parameters.  Importantly we can cancel the damping without conflicting the CMB constraints for both cases, which is not the case for massive neutrinos.  Although much more detailed study is needed to state qualitatively to what extent the constraint on WDM particles can be relaxed, in this paper, we pointed out that some contributions from cosmic strings to small scale power can affect the constraints without conflicting with the CMB constraint.  Since the time dependence of the string tension is highly model-dependent and can be more complicated than that adopted here, various possibilities can arise and have much more implications to the constraints given from observations probing the damping of small scale fluctuation via free streaming."
    },
    "0608/hep-ph0608265_arXiv.txt": {
        "abstract": "{Terascale supersymmetry has the potential to provide a natural explanation of the dominant dark matter component of the standard $\\Lambda$CDM cosmology. However once we impose the constraints on minimal supersymmetry parameters from current particle physics data, a satisfactory dark matter abundance is no longer {\\it prima facie} natural. This Neutralino Tuning Problem could be a hint of nonstandard cosmology during and/or after the Terascale era. To quantify this possibility, we introduce an alternative cosmological benchmark based upon a simple model of quintessential inflation. This benchmark has no free parameters, so for a given supersymmetry model it allows an unambiguous prediction of the dark matter relic density. As a example, we scan over the parameter space of the CMSSM, comparing the neutralino relic density predictions with the bounds from WMAP. We find that the WMAP--allowed regions of the CMSSM are an order of magnitude larger if we use the alternative cosmological benchmark, as opposed to  $\\Lambda$CDM. Initial results from the CERN Large Hadron Collider will distinguish between the two allowed regions.} ",
        "introduction": "One of the most appealing features of Terascale supersymmetry \\cite{Chung:2003fi} is the potential to provide a natural explanation of the dominant dark matter component of the standard $\\Lambda$CDM cosmology.  This explanation is driven by particle physics motivations from data which are completely unrelated to the astrophysical data that motivate $\\Lambda$CDM. In most supersymmetry (SUSY) models the dark matter candidate is the neutralino, although there are other interesting possibilities. The mass and annihilation cross section of the neutralino are determined by the SUSY model parameters. Adapted to a $\\Lambda$CDM cosmological history, the neutralino becomes a cold thermal relic, freezing out during a radiation-dominated era at a temperature equal to approximately 1/20th of its mass. SUSY model parameters thus determine the current neutralino abundance, which is usually expressed as a fraction, $\\Omega_{\\chi}$, of the critical density.  If we assume that neutralinos constitute nearly all of the non-baryonic matter at late times, then their abundance is quite constrained by astrophysical data. For example, combining the WMAP three year data with Sloan Digital Sky Survey data on large scale structure, one obtains \\cite{Spergel:2006hy} the (naive) 2-sigma limits: \\be\\label{eqn:wmapbounds} 0.095 < \\Omega_{\\chi}h^2 < 0.122 \\; . \\ee Thus the neutralino relic abundance is determined with $\\sim 10$\\%\\ accuracy.  Suppose for simplicity that we consider some subclass of Terascale SUSY models with the neutralino as the stable lightest supersymmetric partner (LSP). The most popular example is the CMSSM, inspired by minimal supergravity \\cite{Barbieri:1982eh}-\\cite{Nath:1983aw}, where the SUSY parameter set is written $m_0$, $m_{1/2}$, $A$, tan$\\,\\beta$, and sign($\\mu$).  Imposing very rough considerations of naturalness, we can require $m_0$, $m_{1/2}$ and $|A|$ to be less than, say, 3 TeV. If we assume $\\Lambda$CDM cosmology we can scan over this parameter space, placing a mark at each point which would predict a relic density satisfying the ``WMAP'' constraints  (\\ref{eqn:wmapbounds}).  The results look like Figures \\ref{fig:sm5hfs} and \\ref{fig:sm30fs}, where we have chosen $A=0$, $\\mu >0$, tan$\\,\\beta = 5$ or 30, and we only plot points which also satisfy the existing experimental bounds on the superpartner particles and the Higgs. The ``WMAP--allowed'' regions are a very small fraction of the total parameter space. More significantly, once we impose experimental bounds the WMAP--allowed points are not generic; for a Bino-like neutralino generic points predict a relic density that is much too large, while for a Higgsino-like or Wino-like neutralino generic points predict a relic density that is much too small. Getting a WMAP--allowed relic density thus requires tuning to regions where various conspiracies take place, either enhancing the neutralino annihilation (or co--annihilation) cross section, or balancing the Bino-Higgsino-Wino content of the LSP. These conspiracies can be quantified by defining simple sensitivity measures of the predicted relic density to small variations in relevant SUSY parameters \\cite{Arkani-Hamed:2006mb}. Since conspiracy is the opposite of naturalness, such an analysis \\cite{Baer:2002gm}-\\cite{Mambrini:2005cp} casts doubt upon the CMSSM, and perhaps Terascale SUSY in general, as a natural explanation of dark matter.  \\begin{figure} \\centerline{\\epsfxsize 3.0 truein \\epsfbox {sm5hfs.epsf}} \\caption{\\label{fig:sm5hfs} The WMAP--allowed points for the CMSSM, with $A=0$, $\\mu >0$, and tan$\\,\\beta =5$. Points are shown only if the corresponding Higgs mass is $\\ge 114$ GeV. The scan was performed in 2 GeV increments. \\hbox to170pt{}} \\end{figure}  This ``Neutralino Tuning Problem'' is similar to the much more famous ``Little Hierarchy Problem'' \\cite{Barbieri:2000gf,Birkedal:2004xi} of the MSSM. While conceptually distinct, these problems are related by the fact that the superpartner and Higgs mass bounds from data eliminate what would otherwise be a generic WMAP--allowed region of the CMSSM known as the ``bulk'' region. The bulk region is generic because its only distinguishing feature is that the superpartners are light, as favored by naturalness applied to electroweak symmetry breaking. Thus at least in the CMSSM the two naturalness problems are coupled.  In this paper we revisit the important question of whether we can quantitatively evaluate the claim that Terascale supersymmetry provides a natural explanation of dark matter. This issue has not been fully resolved in the literature, due to a number of difficulties which will we review in the next section. We will show that a physically meaningful approach to this question is to compare the robustness of the connection between SUSY and dark matter under alternative well-motivated cosmological scenarios.  There is ample motivation to vary our cosmological assumptions. Despite the dramatic success of the standard $\\Lambda$CDM model in confronting a wide variety of astrophysical data, this same data is largely silent on the cosmological history of the universe before big bang nucleosynthesis (BBN), {\\it i.e.}, at temperatures above about 4 MeV. WMAP data is consistent with an epoch of primordial inflation, but we are very far from being able to flesh out these hints into a data-driven model which is sufficiently robust and detailed to make definitive statements about dark matter.  A difficulty in exploring cosmological alternatives is that, because Terascale cosmology\\footnote{In this phrase `Terascale' refers to the physics responsible for the WIMP dark matter, not the freeze--out temperature.} is almost a black box, it is easy to introduce new parameters which are relevant but unconstrained, modifying the neutralino relic density predictions by orders of magnitude. For example, we can dilute the neutralino density by adding entropy to the radiation bath some time after neutralino freeze--out but before BBN. Alternatively we can increase the predicted neutralino abundance by producing them non-thermally. Our solution is to look at alternative cosmological scenarios which are reasonably motivated, make definite predictions for dark matter, and are sufficiently simple and constrained that they have no relevant adjustable parameters. The dark matter predictions of Terascale SUSY can then be compared for these cosmological benchmark scenarios.  As a first alternative cosmological benchmark, we propose the Slinky model of quintessential inflation described in \\cite{Barenboim:2005np,Barenboim:2006rx}. This model has no adjustable parameters once we require that the current radiation and dark energy fractions take their WMAP--preferred values, and that the universe is overwhelming radiation-dominated during the time of BBN. Since the same inflaton is responsible for primordial inflation and for dark energy, it is not surprising that the universe is not completely radiation-dominated at the time that neutralinos freeze out. By making the minimal assumption that inflatons do not decay to neutralinos, we have a picture in which thermally produced neutralinos are diluted by a predictable amount, over and above the dilution of standard $\\Lambda$CDM. This additional dilution of the neutralino abundance can be expressed (to an accuracy of a few percent) as a polynomial function of the square root of the neutralino freeze--out temperature.  Thus for any SUSY model we can compute the predicted neutralino relic abundance for two contrasting cosmological benchmarks, $\\Lambda$CDM and Slinky. The ratio of the WMAP--allowed regions is a physically meaningful naturalness comparison. In our CMSSM scan, we find that this ratio is always small. For the Slinky benchmark the WMAP--allowed points are more generic, without balanced mixings, mass degeneracies, or resonance-inducing mass relations.  These results give a concrete measure of the Neutralino Tuning Problem in the CMSSM. At the same time, they define the beginnings of a straightforward program to probe Terascale cosmology at colliders. ",
        "conclusions": "In this paper we have reviewed the Neutralino Tuning Problem of the CMSSM and MSSM. We have quantified this problem by an explicit comparison of the CMSSM for two different cosmological benchmark scenarios: standard $\\Lambda$CDM and Slinky. The comparison shows that the standard case requires about an order of magnitude of extra tuning to successfully account for dark matter. The WMAP--allowed regions for the Slinky case are still restricted to rather narrow strips, but this largely reflects the small ($\\pm$ 10\\%) error bar on $\\Omega_{\\chi}h^2$ that we are trying to squeeze into.  We expect that it should be possible to construct interesting new cosmological benchmarks from examples already in the literature \\cite{Giudice:2000ex}-\\cite{Drees:2006vh}. Note that a different cosmological benchmark with a larger dilution factor would look even more natural than Slinky. Suppose for example that we fix $m_0$, $A$, sign($\\mu$) and tan$\\,\\beta$, and look at \\begin{eqnarray} f(m_{1/2}) \\equiv \\vert d\\Omega_{\\chi}/dm_{1/2}\\vert \\; . \\end{eqnarray} We can regard $f(m_{1/2})$ as a rough naturalness measure; larger $f(m_{1/2})$ in a region where we obtain WMAP-allowed points means that we will obtain a narrower strip of allowed points. Generically $f(m_{1/2})$ is largest at both small and large values of $m_{1/2}$, passing through zero at some intermediate value of $m_{1/2}$ where $\\Omega_{\\chi}$ is maximized. Thus a cosmological benchmark with a larger dilution factor will pick up WMAP-allowed points in a region where $f(m_{1/2})$ is smaller. In addition, when we have a dilution factor $\\gamma$, the relevant naturalness measure is not $f(m_{1/2})$ but rather $f(m_{1/2})/\\gamma$, since it is the latter quantity which determines the width of the WMAP-allowed region in the approximation that $\\gamma$ is a slowly varying function of $m_{1/2}$.  The superpartner spectra in the standard and Slinky cases differ by at least tens of GeV. In addition, the Slinky spectra do not exhibit the mass degeneracies or resonance-inducing mass relations that occur in the standard case. Thus it should be possible to discriminate between these spectra using initial data runs from the LHC. This opportunity deserves detailed study.  Suppose that experiments at the LHC do discover a spectrum of new particles consistent with the CMSSM, but not consistent with the WMAP upper bound as derived assuming a neutralino LSP and standard cosmology. Then we will face the interesting challenge of trying simultaneously to refine our knowledge of both the underlying Terascale physics and the underlying Terascale cosmology. Though both may be complicated, it seems prudent to develop a strategy that begins with simple benchmarks. The CMSSM+Slinky benchmark described here is a step in this direction.      \\subsection*"
    },
    "0608/astro-ph0608566_arXiv.txt": {
        "abstract": "We investigate the effect of shear in the flow of charged particle equilibria that are unstable to the Coherent Synchrotron Radiation (CSR) instability. Shear may act to quench this instability because it acts to limit the size of the region with  a fixed phase relation between emitters. The results are important for the understanding of astrophysical sources of coherent radiation where shear in the flow is likely. ",
        "introduction": "When the wavelength of synchrotron radiation emitted by a bunch of relativistic particles is comparable to the size of the bunch the particles may radiate coherently. The coherent emission produces significantly greater power than the incoherent emission. Coherent synchrotron emission is seen from different astrophysical source but  most notably from the rotationally powered radio pulsars (e.g., \\cite{ManchesterTaylor1977}). In particle accelerators the coherent synchrotron emission is usually undesirable because it  causes very rapid energy loss. In this context we would also like to mention efforts to study coherent curvature radiation as a pulsar emission mechanism in the laboratory \\cite{benford1981}.  Several effects may act to stabilize a system unstable to the CSR instability. In \\cite{Schmekel:2004jb} the authors analyzed the influence of a small energy spread in a beam of charged particles in approximately circular motion. A distribution function with a single value of the canonical angular momentum was considered. The radial width of the beam was given by the amplitude of the betatron oscillations  which is non-zero for a non-zero energy spread. It was shown that the decoherence introduced by the betatron oscillations leads to a characteristic frequency spectrum, whereas the dependence on the Lorentz factor and the number density remains unaffected.   In this paper we relax the assumption of zero spread in the canonical angular momentum $P_{\\phi}$ of the equilibrium distribution. In general this leads to shear in that the average angular velocity becomes dependent on the radius of the orbit. In practice the requirement of a small spread in the canonical angular momentum may be harder to satisfy than a small energy spread. Shear is expected in astrophysical sources of coherent emission \\cite{arons1979}.    Shear itself can be the cause of instabilities, for example, the ``diocotron instability'' \\cite{Davidson2001}, but this is not the focus  of the present paper. In the case of CSR it is reasonable to expect that the shear acts to stabilize the CSR instability due to particles with different radii ``slipping away''. Using a linear perturbation analysis of the fluid equations for a laminar Brillioun flow we show that even with a spread of the canonical angular momentum $\\Delta P_{\\phi} > 0$, previous results for equilibria with constant $P_\\phi$ can be recovered treating the plasma as a relativistic cold fluid. Computer simulations of CSR emission in Brillioun flows \\cite{Schmekel:2004su}  are compatible with this picture. With the cold fluid approximation adopted here, the stability depends on the number density and the angular velocity which depends on the radius. The azimuthal velocity is approximately equal to the speed of light.  The problem of CSR emission has been investigated by several authors. Goldreich and Keeley \\cite{GoldreichKeeley1971} considered the stability of a charge distribution whose motion is confined to a thin ring with  the particle motion being one-dimensional. This calculation led to confusion as to how the proposed CSR mechanism works in detail (cf. \\cite{buschauer1978}).  As proposed in \\cite{Schmekel:2004jb} the CSR instability is related to the classical negative mass instability \\cite{Nielson1959,Briggs1966,Entis1971,Kolomenskii1959} in the sense that an increase in particle energy leads to a decrease in its angular velocity. While the classical negative mass instability is caused by the Coulomb part of the electromagnetic potential, the CSR instability is caused by the radiation field. The negative mass effect is not immediately apparent in a one-dimensional treatment based on conservation of energy and charge. Heifets and Stupakov \\cite{Heifets:2002un} effectively built in the negative mass effect by hand having a constrained radius and particle energy.   The treatments \\cite{GoldreichKeeley1971,Heifets:2002un, Schmekel:2004jb} give the growth rate in the absence of an energy spread. However, considering a non-zero energy spread requires a truly two-dimensional model. Larroche and Pellat investigated the effect of steep boundaries in the particle distribution function \\cite{larroche1987}.   Section II describes the assumed models, and Section III  approximate solutions of the equations. Section IV derives a dispersion relation for the considered perturbation. Section V discusses the results. \\begin{figure} \\includegraphics[width=3.3in]{CSRFig1.eps} \\caption{Panel (a) shows the geometry of a relativistic E-layer for the case of a uniform external axial magnetic field. Panel (b) shows the second case of a relativistic E-layer in an external toroidal magnetic field and an external radial electric field. } \\label{config} \\end{figure} ",
        "conclusions": "In order for the growth rate to vanish the expression inside the root must be real and non-negative. Using our approximations for the Bessel functions it is seen immediately that the former condition is satisfied if $m>m_2$. This explains why the drop off occurs at $m=m_2$ regardless of the value of $\\delta$ as long as the remaining real part is positive. This is the case if the field reversal parameter does not exceed the critical threshold \\begin{eqnarray} \\zeta > m \\gamma \\delta^2  ~. \\label{zetathreshold} \\end{eqnarray} If the last inequality is not satisfied complex roots can only exist if $Z$ is complex  which is the case for $m < m_2$. If however this inequality is satisfied the expression inside the root becomes negative for $m > m_2$ and an unstable solution exists, but the sharp cut-off is replaced by Eq.~\\ref{zetathreshold}.   As discussed in the last section the azimuthal mode number $m$ resulting in the largest growth rate is either given by $m=2 \\gamma^3$ or $m=\\zeta \\gamma^{-1} \\delta^{-2}$ whatever is greater. This guarantees that $(\\Delta \\tilde \\omega)^2$ is either still complex or real and negative with both leading to an unstable mode. For our allowed parameter range $2 \\gamma^3$ is typically larger. Thus, \\begin{eqnarray} \\Im(\\omega) \\le 2^{2/3} \\dot \\phi \\gamma \\sqrt{\\zeta} \\sim \\gamma^{-1/2} \\omega_p \\end{eqnarray} where $\\omega_p$ is the nonrelativistic plasma frequency.   Finally, let us consider small values of the azimuthal mode number $m$. So far we have have approximated the  Bessel functions by Airy functions which makes them easier to compute especially for large orders. For small $m$ this approximation cannot be justified and the Bessel functions have to be retained. In Figure \\ref{smallm} we solved the dispersion relation in the small $m$ regime. Also, their arguments depend on $\\Delta \\tilde \\omega$ since $m<m_1$. An accurate calculation of the growth rates of modes with azimuthal mode nubers $m$ is important for the determination of the total power radiated, because the latter decreases with increasing $m$. Even for $m=1$ Eq.~\\ref{JmAi} and \\ref{YmBi} give excellent results, presumably because the argument of the Bessel functions is close to their order for $\\gamma \\gg 1$ and $Z$ is independent of $w$ if $|w|^2 = \\left | \\left ( \\frac{m}{2} \\right )^{2/3} \\left [ \\gamma^{-2} - 2 \\Delta \\tilde \\omega - 2 (\\xi-1) \\right ] \\right |^2 \\ll 1$. The latter condition may be satisfied even if $|\\Delta \\tilde \\omega|$ violates our assumption $|\\Delta \\tilde \\omega| \\ll \\gamma^{-2}$.  Looking at Eq.~\\ref{evalint} we conclude that the quenching of the instability, i.e. the existence of the additional $\\delta^2$ term in the denominator, is due to both the non-zero thickness of the layer and $\\Omega' \\neq 0$. However, it is important to note that the instability relies on the negative mass effect and would not exist in the absence of shear in our model. \\begin{figure} \\includegraphics[width=3.3in]{CSRFig3.eps} \\caption{Growth rates as a function of azimuthal mode $m$ for $\\zeta=0.02$, $\\gamma=30$ and various energy spreads. The Bessel functions were retained in order to compute the growth rates in the low $m$ regime correctly.} \\label{smallm} \\end{figure}"
    },
    "0608/astro-ph0608085_arXiv.txt": {
        "abstract": "We investigate the propagation of ultra-high energy cosmic ray nuclei ($A=1-56$) from cosmologically distant sources through the cosmic radiation backgrounds. Various models for the injected composition and spectrum and of the cosmic infrared background are studied using updated photodisintegration cross-sections. The observational data on the spectrum and the composition of ultra-high energy cosmic rays are jointly consistent with a model where all of the injected primary cosmic rays are iron nuclei (or a mixture of heavy and light nuclei). ",
        "introduction": "The origin of the highest energy cosmic rays is among the most pressing questions in astroparticle physics \\cite{Nagano:ve,Anchordoqui:2002hs}. The ultra-high energy cosmic ray (UHECR) spectrum has been measured, using both air shower and atmospheric fluorescence techniques, to energies beyond $10^{20}$~eV \\cite{Lawrence:1991cc,Bird:1994uy,Takeda:1998ps,Abu-Zayyad:2002sf}. Protons of such high energies are not expected to be deflected significantly by galactic or extragalactic magnetic fields \\cite{Dolag:2004kp}, but should scatter inelastically on the cosmic microwave background (CMB) with an attenuation length of $\\lesssim 30$ Mpc, i.e. about the size of the local supercluster of galaxies. The observed UHECRs do not point back to any plausible sources within this range and have an isotropic sky distribution, so their sources, if astrophysical, must be cosmologically distant. If so the spectrum should be suppressed at energies above $E_{\\rm GZK} \\sim 6\\times 10^{19}$~eV (the ``Greisen-Zatsepin-Kuzmin cutoff''), if the primaries are indeed protons \\cite{Greisen:1966jv,Zatsepin:1966jv}. Presently, there are conflicting claims concerning the existence of this spectral feature. Data from AGASA (Akeno Giant Air Shower Array) show no indication of any suppression \\cite{Takeda:1998ps}, while data from the HiRes (High Resolution Fly's Eye) air fluorescence detectors are consistent with the expected cutoff \\cite{Abu-Zayyad:2002sf}. The Pierre Auger Observatory, which employs both techniques, has accumulated an exposure comparable to AGASA and about half that of HiRes. Its results are presently consistent with either possibility \\cite{Sommers:2005vs} and forthcoming data should be able to settle the issue definitively. The AGASA observation that the spectrum continues smoothly beyond the GZK cutoff has motivated many suggestions for the origin of the trans-GZK events such as decaying superheavy dark matter particles in the Galactic halo \\cite{Berezinsky:1997hy,Birkel:1998nx} and ultra-high energy neutrinos creating `Z-bursts' in local interactions with the cosmic neutrino background \\cite{Weiler:1997sh,Fargion:1997ft}. It has also been suggested that the primaries may be neutrinos with large interaction cross-sections at ultra-high energies \\cite{Domokos:1998ry,Jain:2000pu}, or perhaps new stable strongly interacting particles heavier than the proton \\cite{Albuquerque:1998va}.  Rather less exotic ways of evading the GZK cutoff are also possible. In particular, if UHECRs are not protons but consist of heavy nuclei \\cite{heavycase}, then the spectrum will be altered from the usual expectation. Although evidence for the presence of heavy nuclei in UHECRs has been around for quite some time \\cite{Bird:1993yi,yakutsk,earlyevidence}, the information available on UHECR composition is still rather imprecise~\\cite{Watson:2004ew}, the most reliable result being that the primaries are not mostly photons but hadrons \\cite{Halzen:1994gy,Shinozaki:2002ve,Ave:2000nd,Risse:2005hi}. Despite the widespread impression to the contrary, however, there is little reason to believe that these particles are protons rather than heavier nuclei.  Determination of the mass composition of UHECRs has been attempted using a variety of variables such as the rate of change with primary energy of the depth of shower maximum $X_\\mathrm{max}$, the degree of fluctuation in the depth of the shower maximum, the number of muons which reach the Earth's surface, the lateral density profile with respect to the shower axis, the width of cosmic ray shower disks, etc. Such measurements must be compared to Monte Carlo simulations, typically conducted with programs such as QGSJET, SIBYLL and DPMJET \\cite{Knapp:2002vs}.  The mass composition inferred is quite sensitive to the simulation program used. Furthermore, data from different experiments can yield rather different results regarding cosmic ray composition even when the same shower simulation is used. For example, at energies near $10^{18}$~eV, the mass composition determined from the lateral density profile of Volcano Ranch and Haverah Park data are in disagreement at $>2\\sigma$ significance \\cite{Dova:2004nq}. The situation does not improve much at higher energies. Above $10^{19}$~eV, only the Fly's Eye/HiRes and AGASA experiments have reported any information pertaining to mass composition, and these results extend to only about $2.5 \\times 10^{19}$~eV --- in this energy range the data suggest a predominantly light composition although the uncertainties are large \\cite{Abbasi:2004nz}.  So far no individual sources of UHECRs have been identified. This can be understood if there are a very large number of faint sources or, alternatively, if there are large-scale magnetic fields which deflect UHECRs sufficiently so as to to conceal their origin. As mentioned above, protons at trans-GZK energies are not expected to be significantly deflected \\cite{Dolag:2004kp} (but see ref.\\cite{Sigl:2003ay}). Heavy nuclei, by contrast, have greater electric charge and may thus have their arrival directions isotropised, even at such high energies.  There is also a theoretical prejudice to expect the UHECR spectrum to be dominated by heavy nuclei, since the maximum energy to which particles can be accelerated scales with the electric charge. Thus astrophysical candidates for UHECR accelerators such as active galactic nuclei and gamma-ray bursts, which barely meet the ``Hillas criterion''~\\cite{hillas} for containment and acceleration of $10^{20}$~eV protons, can in principle easily accelerate heavy nuclei to such energies.  Given these arguments, and the inconclusive nature of the experimental data regarding UHECR composition, it is rather surprising that most studies in this field have usually {\\em assumed} that UHECRs are protons. By contrast, relatively little attention has been paid to the propagation of ultra-high energy nuclei \\cite{Anchordoqui:1997rn,Epele:1998ia,Stecker:1998ib} although there has been a resurgence of interest in recent years \\cite{Bertone:2002ks,Yamamoto:2003tn,Khan:2004nd,Armengaud:2004yt,Allard:2005ha,Sigl:2005md,Allard:2005cx}. We revisit this problem, using up-to-date nuclear physics data, and considering a wide range of nuclei as possible UHECR primaries. We also explore different models for the cosmic infrared background (CIB), as well as the effects of possible intergalactic magnetic fields. We pay particular attention to the relationship between the composition of UHECRs at Earth and that injected at source.  We begin by considering the intergalactic propagation of ultra-high energy (UHE) protons (Section~\\ref{protons}). In Sections~\\ref{photodis}, \\ref{irb} and \\ref{prop}, we study the photodisintegration of cosmic ray nuclei, including the effect of uncertainties in the cross-sections and the CIB model used. We then proceed to discuss the composition of UHECRs observed at Earth (Section~\\ref{comp}), and the role of magnetic fields in the propagation of cosmic ray nuclei (Section~\\ref{bfields}). We propose some representative models of UHECR primary composition and discuss $X_\\mathrm{max}$ measurements in Sections~\\ref{mixedsec} and \\ref{X_{max}}. Our conclusions are presented in Section~\\ref{conc}. ",
        "conclusions": "\\label{conc}  We have studied the intergalactic propagation of a variety of heavy and intermediate mass cosmic ray nuclei at ultrahigh energies. Adopting different models for the cosmic infrared background and for the photodisintegration cross-sections has little effect on the propagated energy spectrum and composition at Earth. Of more significance is the choice of the source spectrum and the effect of intergalactic magnetic fields. Our main aim was to determine the relationship between cosmic ray composition at source and at Earth after the effects of propagation are taken into account. Somewhat surprisingly, extant data on the composition of UHECRs is consistent with the injection of even pure iron nuclei by the sources.  As the Pierre Auger Observatory continues to accumulate more exposure to UHECRs, a definitive resolution is expected soon of whether there is indeed a GZK cutoff in the spectrum. By combining the energy spectrum obtained using the surface detectors with $X_\\mathrm{max}$ measurements by the fluorescence telescopes, information can then be extracted on the sources of UHECRs using the results presented in the present paper. Observations of other types of messengers associated with the highest energy cosmic rays, such as the cosmogenic neutrino flux, will also help to determine whether the ultra-high energy cosmic rays are protons or heavy nuclei \\cite{nu} amd bring us closer to answering the long standing mystery of their origin."
    },
    "0608/astro-ph0608420_arXiv.txt": {
        "abstract": "Massive stars rarely show intrinsic X-ray variability. The only O-stars credited to be intrinsically variable are $\\theta^1$ Ori C due to effects from magnetic confinement of its wind, and $\\theta^2$ Ori A suspected of similar activity. Early \\chandra observations have shown that the most massive star system in the Orion Trapezium Cluster, $\\theta^2$ Ori A, shows rapid variability on time scales of hours. We determine X-ray fluxes from observations with \\chandra and find that the star shows very strong variability over the last 5 years. We observed a second large outburst of the X-ray source in November 2004 with the high resolution transmission grating spectrometer on-board \\chandra and compare the emissivity and line properties in states of low and high flux. In the low state X-ray emissivities indicate temperatures well above 25 MK. In the high state we find an extended emissivity distribution with high emissivities in the range from 3 MK to over 100 MK. The outburst event in stellar terms is one of the most powerful ever observed  and the most energetic one in the ONC with a lower total energy limit of 1.5$\\times10^{37}$ ergs. The line diagnostics show that under the assumption that the He-like ions are photoexcited the line emitting regions in the low states are as close as within  1 -- 2 stellar radii from the O-star's photosphere, whereas the hard states suggest a distance of 3 -- 5 stellar radii. We discuss the results in the context of stellar flares, magnetic confinement, and binary interactions. By matching the dates of all observations with the orbital phases of the spectroscopy binary orbit we find that the two outbursts are very close to the periastron passage of the stars. We argue that the high X-ray states are possibly the result of reconnection events from magnetic interactions of the primary and secondary stars of the spectroscopic binary. Effects from wind collisions seem unlikely for this system. The low state emissivity and R-ratios strengthen the predicament that the X-ray emission is enhanced by magnetic confinement of the primary wind. We also detect Fe fluorescence indicative of the existence of substantial amounts of neutral Fe in the vicinity of the X-ray emission. ",
        "introduction": "}  The number of massive stars in the Orion Nebula Cluster is small compared to other young star clusters of its class, but the few it harbors gained much attention in recent years. At the forefront are the members of the Orion Trapezium, specifically the peculiar O5.5V star $\\theta^1$ Ori C~\\citep{schulz2001,gagne2005}, which have been under some scrutiny lately for their peculiar X-ray behavior ~\\citep{schulz2003,stelzer2005}. Most of the massive stars in the ONC show magnetic properties of some form, either intrinsically or due to an unseen T Tauri companion. With respect to intrinsic magnetic peculiarity there are so far three massive OB-stars with direct evidence of a sizable magnetic field, $\\theta^1$ Ori C, HD 191612, and $\\tau$ Sco ~\\citep{donati2002, donati2006a, donati2006b}. Other massive Trapezium stars exhibit similar X-ray spectral properties to $\\theta^1$ Ori C, though in none of these stars has a magnetic field been directly measured nor has any type of X-ray variability been reported until recently~\\citep{feigelson2002, stelzer2005}. Possible magnetic properties of some sort of fossil origin have been suggested for most of them~\\citep{schulz2001, schulz2003}.  The primary star in $\\theta^2$ Ori A is the second most massive early type star in the ONC after $\\theta^1$ Ori C. It is part of a massive system that had its share of mystery in the early days of X-ray astronomy, when some studies proposed to identify the star as the optical counterpart of the X-ray source 2U 0526-06~\\citep{barbon1972, aikman1974} and later of 4U 0531-05~\\citep{bernacca1979}. One of the main properties that made it a good candidate was its high mass function. Its massive primary, a fifth magnitude O9.5V star~\\citep{abt1991} of 25~\\Msun \\citep{preibisch1999}, indicated the existence of one or more massive companions. Specifically the lack of spectral signatures from a companion star at the time hinted the possibility of the existence of a compact object. Today it is recognized that the system is unlikely to harbor a compact star but is considered to be at least a massive triple system containing a visual companion with a mass of $< 8$~\\Msun separated by $\\sim 0.4\"$~\\citep{preibisch1999}, and a spectroscopic companion to the primary with a mass of likely $< 7 - 9$~\\Msun \\citep{abt1991}. The stellar types of both companion stars are still highly uncertain and range from a massive early A-type star for the visual companion to a B-type star as early as about B5 for the spectroscopic companion. Notable is also the physical closeness of these massive companions, which amounts for the visual companion to be 174 AU and the spectroscopic companion to be 0.47 AU~\\citep{preibisch1999}, approximately 10 stellar radii off the primary. The system is likely as young as $\\theta^1$ Ori C, i.e. $\\sim$ 0.3 Myr, which is also reflected in the fact that the orbital motion of the spectroscopic binary and the O-star's rotational period are not yet synchronized. It cannot be much younger as it already arrived on the zero-age main sequence, but also not much older through its association with the Orion Trapezium.  In a recent study of X-ray emitting young stars in the Orion Nebula, ~\\cite{feigelson2002} found that $\\theta^2$ Ori A and several early B-stars show short time scale variability. Especially $\\theta^2$ Ori A exhibited the most dramatic X-ray variability with a 50$\\%$ flux drop over a time period of 10 hours superimposed by multiple small flares with a few hour durations. Until today this kind of variability in a massive system stands out as the most dramatic ever observed. The Chandra Orion Ultradeep Project (COUP) observations~\\citep{getman2005} a few years later covered the X-ray intensity of that star in several episodes for over 13 days and found it only weakly variable with the largest variation around a factor 1.8~\\citep{stelzer2005}. The observed variability in $\\theta^2$ Ori A clearly represents problems not only with the standard model of wind shocks in massive stars but also with the proposed model of magnetically confined winds in these stars~\\citep{gagne2005}. \\citet{feigelson2002} acknowledge that in some occasions large excursions in flux can be created by occasional large shocks as proposed by~\\cite{feldmeier1997} but the temperature profile during these events as observed is not consistent with such predictions. In the case of $\\theta^2$ Ori A a scenario was envisioned where magnetic reconnection events near the stellar surface occur, which tentatively and qualitatively could explain the X-ray emission observed in the O9.7Ib supergiant $\\zeta$ Ori~\\citep{waldron2001}.  In this paper we present recent observations of $\\theta^2$ Ori A observed with the high energy transmission grating spectrometer (HETGS) onboard \\chandra in a low and high flux state. ",
        "conclusions": "}  This study reveals several new aspects and details about the X-ray emission of young massive stars. $\\theta^2$ Ori A proves to be a new study case for variable emission from massive systems. The analysis confirms the findings by \\citet{feigelson2002} that this system undergoes strong excursions in X-ray flux. The outburst reported on in this study appears much more dramatic in terms of flux levels and spectral hardening even if one accounts for pileup which heavily distorted the data in the 2000 event. Before we discuss the possible nature of the variable X-ray emission we summarize the results of our analysis.  Over a period of about 5 years the X-ray source showed variability in flux of over an order of magnitude. During the last outburst the range of observed plasma temperatures changed from 3 -- 30 MK to a range of 3 -- $>$100 MK. The X-ray luminosities changed from 2.6$\\times10^{31}$ \\ergsec to 1.7$\\times10^{32}$ \\ergsec which corresponds to a change of the $L_x/L_{bol}$ ratio from $\\sim 10^{-7}$ to $\\sim 10^{-6}$. Under the assumption that all O-stars are X-ray emitters of some form and the one in $\\theta^2$ Ori A is no exception, we have to assume that the radiation field of the primary O9.5V star is the main source of destruction of the forbidden lines in the He-like triplets. The dominance of the UV field puts the X-ray source close, i.e. within a few stellar radii of the primary's photosphere. As a consequence a possible contribution from high density to the destruction of the lines is limited to the order of $ 10^{10}$ cm$^{-3}$ for the low and high state at the most. High densities may only be feasible, if the X-ray source is at a very large distance from the O-star and the entire emissivity is not associated with the O-star wind. We are not discussing this possibility as consequently the O-star itself would not be an X-ray emitter which based on the observational record of X-rays from early type stars is unlikely. The observed emissivities in these two states are then understood not only in terms of a change in the energy balance due to a dramatic change in temperatures, but also due to a vastly increased emissivity volume in the high state. Last but not least, the coincidence that the two outbursts appear close to the nearest approach of the spectroscopy companion suggests that they are related to some form of interaction.  In the following we discuss these results under various aspects. First we put the outbursts into the context of variability observed in the ONC so far with \\chandra. We then attempt to qualitatively interpret the observed X-ray emissivity with respect to intrinsic physical environments such as magnetic confinement of winds, magnetic reconnection, and binary interactions. Last we discuss possible origins of the Fe K fluorescence during the high state.   \\subsection{Hard X-ray Variability and Reconnection\\label{sec:hard}}  \\smallskip \\vbox{ \\includegraphics[angle=0,width=8.5cm]{f7.eps} \\figcaption{The number distribution of the integrated energy of all bright flares in the COUP study~\\citep{favata2005}. \\label{flare}} } \\smallskip  Figure~\\ref{flare} illustrates the extraordinary nature of the observed variability. We integrated the total energy of all bright flares of young late type stars analyzed in \\citet{favata2005} in the COUP data and sorted them by flare energy. We integrated the flares from their observed maximum using the observed decay time. All the COUP flares have well established peak luminosities, temperatures, and cooling time constants. The statistics then shows that most flares are found near a total energy of less than 10$^{36}$ ergs, with the brightest one barely reaching 10$^{37}$ ergs. In this respect the outburst in $\\theta^2$ Ori A from the year 2000 does not look particularly impressive with a total energy of 3$\\times10^{36}$ ergs. However, the lightcurve in \\citet{feigelson2002} does not appear like the flares seen in late type stars  which show a short rise and exponential decay. In both outbursts, though, the lightcurves showed their highest fluxes at the beginning of each observation with some form of decline afterwards. Thus we assumed the flux at the beginning of the lightcurve as the starting point and used the rest of the lightcurves to estimate its decay time. The integrated energy should then be considered a lower limit because we do not really know the actual peak occurrence. The second outburst in 2004 in this respect clearly distinguishes itself from all the other observations as well as generic flares (see Figure~\\ref{newlight}). The emissivity remains high over the entire exposure with some decline in flux. Under the assumption that the late 2004 event was some form of flare event, the resulting lower limit to the total energy is 1.5$\\times10^{37}$ ergs.  Flares this powerful are unprecedented in normal stellar systems and it is very likely that these high states are of a different origin. A somewhat less powerful flare has been recently uncovered in \\ros data of the massive binary $\\sigma$ Ori E~\\citep{groote2004}. For a flare event of coronal type the conditions have to be rather extreme indeed. Dissipation of large electrical currents associated with reconnection events realigning the field leads to the release of magnetic energy. Having established that the hard X-ray source during the outburst is likely located at a few stellar radii from the O-star's surface, a corresponding magnetic loop structure seems extraordinarily large. Even though the star is expected to rotate relatively fast with rotation period between 2.8 days~\\citep{preibisch1999} and 3.5 days~\\citet{kaper1998} the density of the wind at a few stellar radii is likely too low to cause the substantial field dragging necessary to explain the observed energy release. Such rotation rates produce surface velocities of about 1500 km s$^{-1}$ and we should have seen some line broadening from additional shocks or turbulence if the wind would be affected. Applying scaling laws of stellar flares from the Sun also show that reconnection events heat the ambient plasma on time scales of several minutes. Cooling times should be of the order of only a few hours. The outburst of 2004 persists for almost an entire day suggesting that there is some continuous heating.  \\subsection{Effects of Magnetic Confinement\\label{sec:magnetic}}  It is generally assumed that X-ray emission from hot stars originates from shocks in line driven winds. O-stars of all spectral and luminosity classes have so far been observed in X-rays. In hot stars like $\\theta^1$ Ori C, $\\tau$ Sco, and likely $\\zeta$ Ori energy is thought to be deposited into the wind through magnetic channeling resulting in hard X-ray emission. Since the main component of $\\theta^2$ Ori A is of O~9.5 V type, it seems warranted to associate at least some of the observed X-ray emission likely with the wind of the O-star. It has also been recognized in recent years that temperatures much in excess of 15 MK cannot be the result of standard wind shocks but rather the result of more complex structures such as magnetic confinement or interacting winds.  \\vbox{ \\smallskip \\includegraphics[angle=0,width=9.5cm]{f8.eps} \\figcaption{Schematic illustration of configurations of the binary using the orbital parameters given by~\\cite{abt1991}. The system is plotted face-on, the inclination towards the observer is estimated to be 65$^{\\circ}$. The black arrows indicate the direction of motion around the center of mass. The sizes of the stellar companions are scaled about proper to the orbit dimensions. The blue circle represents the O9.5V star, the green circle the spectroscopic companion. The configuration at orbital phase 0.90 marks the projected start of outburst activity, the one at 0.98 marks the outburst observed in November 2005, the one at 0.08 marks the start of the outburst reported by~\\citet{feigelson2002} and the projected end of activity. We used the start of the flux decline in the lightcurve of the 2000 outburst to set the end point. Symmetrically to periastron we then projected the start point of the activity. The filled red areas mark the anticipated locations of the hard outburst X-ray emission. The lighter red shade indicates the decline of the 2000 outburst. The interaction region is large compared to the size of the primary star and its moving through the line of sight should not affect the view of the X-ray source by much. This is visualized in the cartoon in lower right corner, which is now the line of sight view of the observer at 0.08 orbital phase. \\label{orbit}} } \\smallskip  In comparison with other hot stars (see Figure 5 in \\citealt{wojdowski2005}) the DEMs of $\\theta^2$ Ori A show several novel aspects. Both, the high and low states are different in their emissivity from what is expected from shocks in an unobstructed line driven stellar wind. In the low state (red curve) the X-ray spectrum and its corresponding DEM compare well with X-ray emission from $\\tau$ Sco with the difference that we see not only strong emissivity at log T[K] = 7, but also emissivity at much lower temperatures. Unobstructed winds seem rather to peak below log T[K] = 6.5 with a steady decline of emissivity above about 3 -- 5 MK. Two more facts strengthen the comparison with $\\theta^1$ Ori C and $\\tau$ Sco: the emission lines are similarly narrow and the origin of the emission is likely more than one stellar radius above the photosphere of the star. In both cases magnetic fields have been detected, most recently for $\\tau$ Sco~\\citep{donati2006b}. The emission from $\\theta^2$ Ori A in the low state then should be discussed in the context of some form of magnetic channeling of the O-star wind.  In the high state the DEM of Figure~\\ref{dems} differs from anything we have observed before, i.e. a high emissivity from log T[K] = 6.5 to 8.0 with a tail towards even higher temperatures. Peak emissivities beyond log T[K] = 7.5 are known from flares in stars with active coronae and from $\\theta^1$ Ori C. The X-ray  emissivity level in the latter star is an order of magnitude larger than observed at any level in $\\theta^2$ Ori A owing to a stronger wind and a high magnetic field of the order of $\\sim$1 kG~\\citep{gagne2005}. In $\\theta^1$ Ori C the soft emissivity is weak and the bulk of the emissivity is growing towards higher temperatures. In $\\theta^2$ Ori A the outburst DEM rises high at lower temperatures and then remains high over a large range of high temperatures.  Since a magnetic field has not yet been directly measured so far in $\\theta^2$ Ori A we rely on projections from the models presented by \\citet{babel1997} and \\citet{gagne2005} for the cases of IQ Aur and $\\theta^1$ Ori C. Clearly, the contrast between these two cases could not be larger, i.e. the mass loss rates of a few times $10^{-10}$ \\Msun yr$^{-1}$ and $10^{-7}$ \\Msun yr$^{-1}$ differ by three orders of magnitude and the surface temperatures (13000 K and 42000 K, respectively) by over a factor three. $\\theta^2$ Ori A's values are in somewhere between. Mass loss rates of less than $10^{-8}$ \\Msun yr$^{-1}$ and a stellar surface temperature of 30,000 K suggest a strong wind. Under assumption that the wind has a moderate velocity, i.e. v$_{\\infty}$ of the order of 1000 km s$^{-1}$, we would obtain a confinement parameter ($\\eta = B^2R^2/\\dot Mv_{\\infty}$) of 10 for the existence of a magnetic field of $>$ 0.6 kG.  Quantitative modeling by \\citet{uddoula2002} suggests that a confinement parameter of 10 or above should result in effective confinement. The model also predicts that most X-rays are produced in a relatively small region where the density of the hot gas is high, beyond $10^{12}$ cm$^{-3}$. To reach such high densities and temperatures, confined regions form within one ot two Alfv\\'en radii from the stellar surface where field lines are still mostly closed and perpendicular to the magnetic equatorial plane. Thus confinement places the X-ray source close to the stellar surface of the primary. For $\\theta^2$ Ori A the R-ratios in Figure~\\ref{distance} put the distance of the X-ray source less than 1 stellar radius in the case of Mg~XI and  2 -- 4 stellar radii in the case of Ne~IX above the stellar surface. The confinement model and data in the case of $\\theta^1$ Ori C place the region of emission within $\\sim$2 stellar radii~\\citep{gagne2005}.  Our data seem to rule out densities that are as high as $10^{12}$ cm$^{-3}$. However the low state spectrum indicates much lower emissivity at high temperatures than observed in $\\theta^1$ Ori C suggesting that such high densities may not be required. The R-ratios are not entirely conclusive with respect to the location of the emitting region above the stellar surface and predictions from magnetic confinement. In the low state, the Mg~XI R-ratio seems to indicate such a closeness, the Ne~IX ratio on the other hand could allow for the larger range. In the high state the emission region is more like 3 - 5 R$_{star}$ from the O-star surface. Confinement in the latter case seems difficult to achieve. Thus a simple confinement scenario as proposed for $\\theta^1$ Ori C and likely $\\tau$ Sco is feasible for the low state but is problematic for the high state.  \\subsection{Binary Interactions\\label{binary}}  \\subsubsection{Colliding Winds\\label{colliding}}  We also have to consider the fact that all stars in the $\\theta^2$ Ori A system are of early type and thus all are sources of radiatively driven winds. The proximity of the spectroscopic companion could then give rise to some form of wind interaction. Colliding wind scenarios to produce X-ray emission~\\citep{prilutskii1976, stevens1992} have been proposed since the launch of \\ein where a survey of WR stars revealed that WR-O binary system are systematically brighter in X-rays~\\citep{pollock1987}. Recent observations of O-star binaries such as HD93403 (O5.5I + O7V~\\citealt{rauw2002}), HD 159176~\\citep{debecker2004}, and HD152248 (O7III+O7.5III~\\citealt{sana2004}) involve O-stars of very early type or luminosity classes III and above or both. Other cases involving O dwarfs and B-stars such as HD 206267 and $\\iota$ Ori have so far not been conclusive in X-rays~\\citep{pittard2000, schulz2003a}. Colliding wind scenarios are usually associated with exceptionally strong winds with mass loss rates exceeding $10^{-6}$ \\Msun yr$^{-1}$ and terminal velocities of $>$ 2000 km s$^{-1}$. Unusually bright X-rays from very young O-star binaries have been observed in young star-forming regions containing many massive stars and in which strong stellar winds create a diffuse interstellar gas~\\citep{townsley2003}.  The case of $\\theta^2$ Ori A is quite different. Its O9.5V star may produce a mass loss of the order of only about 1$\\times10^{-8}$ \\Msun yr$^{-1}$. Its companion has likely a much lower rate. Recent research also suggests that mass loss rates in Galactic O-type Stars could be even smaller than generally assumed~\\citep{fullerton2006}. A colliding wind scenario is expected to differ with respect to systems like WR 140 or HD152248 where mass loss rates are over 1000 times higher. In HD152248 the separation of the two stars is always close, within 2 primary stellar radii. Figure~\\ref{orbit} shows that even at closest approach, the secondary in $\\theta^2$ Ori A is still over 6 primary radii apart from the O-star surface and wind densities should not exceed a few times $10^{9}$ cm$^{-3}$. However, wind models also suggest that the density cannot be much lower either. At much lower densities heating is provided by collisionless shocks like in WR 140~\\citep{pollock2005} but would require unrealistically high emission volumes on scales in excess of 6 AU. In the optimistic case of effective wind compression which may allow for densities up to $10^{10}$ cm$^{-3}$ the volumes would be on the scale of only a few 0.01 AU. But then we still need an average collision speed  $>> 2000$ km s$^{-1}$ and near 100$\\%$ collision efficiency, again extremely unrealistic conditions.  Furthermore, in colliding wind systems luminosity is inversely proportional to the distance from the companion to the shock interface, which at an orbital eccentricity of the order of 0.33 as suggested~\\citep{preibisch1999} would predict a slow and gradual increase and decrease over significant fractions of the binary orbit. The fact that we have low state observations intermediate to the observed high states points to more explosive origins. We therefore rule out scenarios that involve interacting winds only.  \\subsubsection{Interacting Magnetic Fields\\label{fields}}  Sect. \\ref{sec:hard} already demonstrated that simple magnetic reconnection scenarios related to an O-star's magnetic field alone seem unfeasible. On the other hand, the proximity of the outbursts to the periastron passage suggests the possibility of two body magnetic interaction events. Figure~\\ref{orbit} schematically illustrates the anticipated orbital geometry. In the two body case, the X-ray source is located in between the two stars, which near periastron are separated by about 6 -- 7 primary radii. With an integrated emissivity during the outburst of 2.1$\\times10^{55}$ cm$^{-3}$ the emitting volume has a size of (0.2 AU)${^3}$ assuming realistic plasma densities. Magnetic reconnection thus has to happen on quite large scales. We can estimate an upper limit to the magnetic energy release by calculating the heat produced by the drag of the primary field lines due to the approaching companion to  \\begin{equation} \\frac{d E_{mag}}{dt} = \\eta \\Omega_{rad} B_{surf}^2 R_d^3, \\end{equation}  \\noindent where $B_{surf}$ is the field strength of the O-star primary, $\\Omega_{rad}$ is the absolute difference between the angular velocity of the incoming secondary and the primary rotation, $R_d^3$ the size of the interaction region. The parameter $\\eta$ is an efficiency factor acounting for reconnection efficiency for specific geometries and accounting for the fact that not all magnetic field lines participate in a specific release event. A similar approach has been applied by ~\\cite{shu1997} to estimate the energy release in star-disk interactions in T Tauri stars. For a primary field of $\\sim$ 0.6 kG and an incoming speed of the companion of 250 km s$^{-1}$ we obtain a maximum energy release of $\\sim 10^{36}$ \\ergsec. This value should be compared to the observed luminosity during the outburst, which is a few times $10^{32}$ \\ergsec. This indicates that $\\eta$ can be very small. It can also mean that our basic assumptions are incorrect. However, given the large emissivity volume and a likely weak companion field this may not be unrealistic. Also given the many unknowns in this system with respect to such a complex scenario, we cannot draw more detailed conclusions. More data and modeling clearly seem warranted.  \\subsection{Origins of the Fe Fluorescence\\label{sec:fluor}}  In the high state we also detect a few photons which we attribute to Fe K fluorescence at 1.93~\\AA. The outburst produces a substantial amount of X-ray continuum flux above 7.11 keV (1.74~\\AA), the value of the photoelectric absorption edge. Fluorescence emission emerges from cool, mostly neutral, and optically thick materials located near ionizing environments. In the X-ray band fluorescence is most efficient for Fe atoms because the high atomic number mitigates autoionization through Auger electrons. Fe fluorescence has therefore the highest yields and strong Fe K lines are observed in X-ray binaries with large accretion disks \\citep{gottwald1995}, high column density material in winds~\\citep{sako1999, schulz2002}, proto- and PMS stellar disks~\\citep{tsujimoto2005}, but also in the Sun \\citep{bai1979, parmar1984}.  The fact that we observe fluorescence in such a system is a novelty, if one disregards high mass X-ray binaries, which also contain stars of early type with strong winds but also neutron stars or stellar black holes. Here strong winds from evolved supergiants are almost completely ionized by a compact X-ray source giving rise to fluorescence emission from cool, optically thick clumps dispersed throughout the wind \\citep{sako1999,schulz2002,boroson2003}. \\citet{feldmeier1995} pointed out that winds can have distinct regions of clumpy old materials specifically in the outer regions of the wind. Recent work designed to explain the low attenuation of X-ray lines quite recently led to studies involving a general porosity of stellar winds~\\citep{owocki2006}. Although the wind in the O9.5 dwarf may have such clumpiness on some scale, we do not consider these as a possible source for the fluorescence because they do not provide a sufficient mass with respect to the incident luminosity. The X-ray source in Vela X-1 is of the order of 10$^4$ times brighter than the outburst in $\\theta^2$ Ori A.  We investigated a variety of possible sources and conclude that none of the sites alone can account for the observed photons. Young systems contain a lot of debris and circumstellar material even after several million years after formation and given the projected age of the system of less then 1 Myr one should expect plenty of circumstellar matter in the system. However, the fact that the O-star entered the main sequence also means that its powerful wind has evaporated most of circumstellar material within the spectroscopic binary with a probability of some leftover disk with the visual companion. This leaves us with the optically thick photospheres of the companion as well as any possible disk structure that evolved from the wind confinement.  In order to estimate the expectation of fluorescence from the O-star's photosphere we derived an efficiency for fluorescence from the stellar surface \\citep{bai1979} as $\\epsilon \\sim 2*I_{line}/I_{c}$, where $I_c$ is the incident hard continuum flux above 7.11 keV and $I_{line}$ the measured fluorescence line flux. The model in \\citet{bai1979} was designed to estimate photospheric fluorescence from the Sun, but since its geometry only scales with the stellar radius, we may relate the efficiency to the behavior in the Sun. For $\\theta^2$ Ori A we measured an efficiency $\\epsilon = 0.064 \\pm 0.045$. For the Sun $\\epsilon$ = 0.04 is expected at the surface and drops steeply above one solar radi~\\citep{parmar1984}. Fluorescence from the photosphere alone is thus likely not sufficient to explain the observed line flux which is consistent with the non-detection Fe fluorescence in $\\theta^1$ Ori C~\\citep{schulz2003} where the hard X-ray flux was even higher. This shows that the fluorescence has to include other sources mentioned above."
    },
    "0608/astro-ph0608616_arXiv.txt": {
        "abstract": "Dynamical evolution of spiral galaxies is strongly dependent on non-axisymmetric patterns that develop from gravitational instabilities, either spontaneously or externally triggered. Some evolutionary sequences are described through which a galaxy could possibly concentrate mass and build bulges, how external gas accretion from cosmic filaments could be funneled to the galaxy disks, and intermittently driven to the galaxy center, to form nuclear starbursts and fuel an active nucleus. The frequency of both bars and lopsidedness can be used to constrain the gas accretion rate. ",
        "introduction": "\\subsection{Molecular gas in barred galaxies}  Sakamoto et al (1999) from a sample of 20 nearby galaxies were the first to notice that barred galaxies had a larger gas concentration in the central kiloparsec than unbarred galaxies, confirming that bar instabilities were able to drive gas inflow towards the center. Sheth et al (2005) from a sample of about 45 galaxies from BIMA-SONG confirm the larger CO concentration for barred galaxies, and find that the nuclear molecular gas is 4 times more abundant in early-types than late-types, suggesting the gas inflow are larger, possibly due to the existence of inner Lindblad resonances (ILR), expected when the mass is concentrated due to bulges. Jogee et al (2005) further note that barred galaxies have a large circumnuclear surface densities in molecules, and this triggers the nuclear starbursts currently observed there. The observation of non-starbursting gas-rich barred galaxies allows to distinguish several phases of evolution: the first phase reveals gas inflow, non circular motions (since the gas is partly in the bar), and the gas is accumulating at the outer ILR, there is no starburst. In the end-phase, the gas has been driven down to the inner ILR, runs in nearly circular orbits, and the galaxy experiences a nuclear starburst of 3-11 M$_\\odot$/yr. Molecular gas surface densities reach up to 500-3500 M$_\\odot$/pc$^2$, or 10-30\\% of dynamical mass in the center.  \\subsection{Star formation in the center of spiral galaxies}  Recent work on star formation in the central kpc of galaxies has shown that galaxies with pseudo-bulges are more frequently the host of starbursts. Pseudo-bulges are components intermediate between spheroids and disks, in the sense of light distribution (Sersic index $n$ between 1 and 2) of flattening and kinematics (more rotation than classical bulge, Kormendy \\& Kennicutt 2004). Fisher (2006) has recently made a star-formation study on a sample of 50 galaxies observed by Spitzer. Pseudo-bulges were identified on HST images from the presence of disky structure in the center, i.e. nuclear disks nuclear bars, and high ellipticity. The star formation was traced from the PAH emission at 8 micron wavelength. As shown in Fig \\ref{fisher}, star formation inside 1.5 kpc is clearly larger in pseudo-bulge galaxies, and among them, central starbursts are preferentially in barred galaxies, then in galaxies with oval perturbations. Nuclear star formation in the form of H$\\alpha$ nuclear rings are also preferentially found in barred galaxies (Knapen 2005, and this meeting).  \\begin{figure} \\centering \\includegraphics[width=0.7\\textwidth]{combes-f1.ps} \\caption{Specific SFR of the central 1.5 kpc, for a sample of 50 galaxies, from Fisher (2006), as a function of morphological type (left), and as a function of secular driving mechanism (B=bar, O=oval, and N=neither bar no oval). Pseudo-bulge galaxies (red circles) have clearly larger central specific SFR than classical bulge ones (crosses).} \\label{fisher} \\end{figure}  Peeples \\& Martini (2006) study the relation between bar strength (in terms of Q$_b$) and circumnuclear dust. Dust is distributed preferentially in nuclear rings for barred galaxies. SB(r) galaxies are less strongly barred and have less dust structure than SB(s) galaxies. This can be understood as an evolutionary process, since rings are formed in a second phase in barred galaxies, where first a spiral structure forms.  \\subsection{Bar length and type, $\\Omega_b$ and dynamical friction}  In a recent study of about 140 galaxies, Erwin (2005) confirms that bars are longer in early-types than in late-type galaxies, both when normalized with the optical radius R$_{25}$, or with the exponential scale-length $h$ (see Figure \\ref{erwin}). The size of bars does not depend on their strength.  These observational results can be confronted to predictions from numerical models. The length of bars is tightly related to their pattern speed, since stellar orbits support the bar until its corotation only. In the presence of a large mass of dark matter inside the bar radius, bars should slow down through dynamical friction against the dark matter particles (Debattista \\& Sellwood 2000), and bars should then be too long and slow compared to observations.  This puts constraints on the dark matter distribution in spiral galaxies, and favors maximal disks. But in some cases, the transfer of angular momentum to the dark halo can be minimized (Valenzuela \\& Klypin 2003).   \\begin{figure} \\centering \\includegraphics[width=0.8\\textwidth]{combes-f2.ps} \\caption{Sizes of bars (deprojected) relative to their disc radius R$_{25}$, as a funciton of morphological type, from Erwin (2005). The various symbols correspond to various samples. Large boxes are mean values for each type.} \\label{erwin} \\end{figure}   The length of bars is interesting under the point of view of secular evolution, since it is predicted to vary along the successive evolution phases, and is then a tracer of the past history of the galaxy.  Bars begin their life with a high pattern speed and a short length (as in late-types) and then slow down and increase their length, as in early-types. This simple scenario is slightly modified when bar destruction and reformation is taken into account: with gas accretion, new bars are generally shorter. They appear late-type just after gas accretion (Bournaud \\& Combes 2002). Tidal interactions on the contrary can produce longer bars (Holley-Bockelmann et al 2005), depending on mass and distance of the satellites.  In edge-on galaxies, the length of bars can be traced by the size of the peanut-shape. The maximum height of the peanut occurs at the z-resonance with the bar, a function of its pattern speed. When the bar slows down through angular momentum transfer with the halo, the radius of the resonance shifts outwards. Several peanut buckling can occur successively (Martinez-Valpuesta et al 2006). ",
        "conclusions": ""
    },
    "0608/astro-ph0608170_arXiv.txt": {
        "abstract": "{Precise measurements of surface abundances of extremely low metallicity stars have recently been obtained and provide new constraints for the stellar evolution models.} {Compute stellar evolution models in order to explain the surface abundances observed, in particular of nitrogen.} { Two series of models were computed. The  first series consists of 20 $M_\\odot$ models with varying initial metallicity ($Z=0.02$ down to $Z=10^{-8}$) and rotation ($\\rm \\upsilon_{ini}=0-600$\\,km\\,s$^{-1}$). The second one consists of models with an initial metallicity of $Z=10^{-8}$, masses between 9 and 85 $M_\\odot$ and fast initial rotation velocities ($\\rm \\upsilon_{ini}=600-800$\\,km\\,s$^{-1}$). } { The most interesting models are the models with $Z=10^{-8}$ ([Fe/H]$\\sim-6.6$). In the course of helium burning, carbon and oxygen are mixed into the hydrogen burning shell. This boosts the importance of the shell and causes a reduction of the CO core mass. Later in the evolution, the hydrogen shell deepens and produces large amount of primary nitrogen. For the most massive models ($M\\gtrsim 60$\\,$M_\\odot$), significant mass loss occurs during the red supergiant stage. This mass loss is due to the surface enrichment in CNO elements via rotational and convective mixing. The 85 $M_\\odot$ model ends up as a WO type wolf-Rayet star. Therefore the models predict SNe of type Ic and possibly long and soft GRBs at very low metallicities.  The rotating 20 $M_\\odot$ models can best reproduce the observed CNO abundances at the surface of extremely metal poor (EMP) stars and the metallicity trends when their angular momentum content is the same as at solar metallicity (and therefore have an increasing surface velocity with decreasing metallicity). The wind of the massive star models can also reproduce the CNO abundances of the most metal-poor carbon--rich star known to date, HE1327-2326. } {} ",
        "introduction": "Precise measurements of  surface abundances of extremely metal poor (EMP) stars have recently been obtained by \\citet{FS5}, \\citet{FS6} and \\citet{IER04}. These provide new constraints for the stellar evolution models \\citep[see][]{CMB05,Fr04,Pr05}. The most striking constraint is the need for primary $^{14}$N production in very low metallicity massive stars. Other constraints are an upturn of the [C/O] ratio with a [C/Fe] ratio \\citep[see for example Fig. 14 for C/O and Fig. 13 for C/Fe in][]{FS6} about constant or slightly decreasing (with increasing metallicity) at very low metallicities, which requires an increase (with increasing metallicity) of oxygen yields below [Fe/H]$\\sim$ -3. About one quarter of EMP stars are carbon rich (C-rich EMP, CEMP stars). \\citet{RANB05} and \\citet{BC05} propose a classification for these CEMP stars. They find two categories: about three quarter are main s--process enriched (Ba-rich) CEMP stars and one quarter are enriched with a weak component of s--process (Ba-normal). The two most metal poor stars known to date, HE1327-2326 \\citep{Fr05,Ao06} and HE 0107-5240 \\citep{Ch04} are both CEMP stars. These stars are believed to have been enriched by only one to several stars and the yields of the models can therefore be compared to their observed abundances without the filter of a galactic chemical evolution model (GCE).   The evolution of low metallicity or metal free stars is not a new subject \\citep[see for example][]{CC83,EE83,CBA84,A96}. The observations cited above have however greatly increased the interest in very metal poor stars. There are many recent works studying the evolution of metal free (or almost) massive \\citep{HW02,LC05,UN05,MEM06}, intermediate mass \\citep{SLL02,H04,SAMFI04,GS05} and low mass \\citep{PCL04,W04} stars in an attempt to explain the origin of the surface abundances observed. In this work pre-supernova evolution models of rotating single stars were computed with metallicities ranging from solar metallicity down to $Z=10^{-8}$ to study the impact of rotation and to see which initial rotation velocities can lead to the chemical enrichments observed. In Sect. 2, the model physical ingredients and the calculations are presented. In Sect. 3, the evolution of the models is described. In Sect. 4, the stellar yields of light elements are presented and compared to observations. In Sect. 5, the conclusions are given. ",
        "conclusions": "Two series of models were computed. The  first series consists of 20 $M_\\odot$ models with varying initial metallicity (solar down to $Z=10^{-8}$) and rotation ($\\rm \\upsilon_{ini}=0-600$\\,km\\,s$^{-1}$). The second one consists of models with an initial metallicity of $Z=10^{-8}$, masses between 9 and 85 $M_\\odot$ and fast initial rotation velocities. The results presented confirm the crucial role of rotation in stellar evolution and its impact in very low metallicity stars \\citep{MEM06}. The evolution of the models with $Z=10^{-8}$ ([Fe/H]$\\sim-6.6$) is very interesting. In the course of helium burning, carbon and oxygen are mixed into the hydrogen burning shell. This boosts the importance of the shell and causes a reduction of the CO core mass. Later in the evolution, the hydrogen shell deepens and produces large amount of primary nitrogen. For the most massive models ($M\\gtrsim 60$\\,$M_\\odot$), significant mass loss occurs during the red supergiant stage assuming that CNO elements are important contributors to mass loss. This mass loss is due to the surface enrichment in CNO elements via rotational and convective mixing. The models predict the production of WR stars for an initial mass higher than 60 $M_\\odot$ at $Z=10^{-8}$ and the 85 $M_\\odot$ model becomes a WO. Therefore SNe of type Ib and Ic are predicted from single massive stars at these low metallicities. The 85 $M_\\odot$ model retains enough angular momentum to produce a GRB but the calculations did not include the effects of magnetic fields.  The stellar yields are presented for light elements. These yields were used in a galactic chemical evolution model and successfully reproduce the early evolution of CNO elements \\citep{CH06}. A scenario is also proposed to explain the abundances of the most metal poor star known to date HE1327-2326 \\citep{Fr05}. In this scenario, a first generation of stars (PopIII) pollutes the interstellar medium to very low metallicities ([Fe/H]$\\sim$-6). Then a PopII.5 star \\citep{Br05} like the 40 $M_\\odot$ model calculated in this study pollutes only with its wind the interstellar medium out of which HE1327-2326 forms.   There are still many questions or issues that could not be treated in this work. It is necessary to determine over which metallicity range the large primary production and the other specific features of the $Z=10^{-8}$  models occur. The impact of the yields on the initial composition and therefore on the evolution of the next stellar generations also needs to be studied. It will also be interesting to follow the evolution of Na, Mg and Al since the high yields of $^{22}$Ne seem to indicate that there could an overproduction of these elements in the wind \\citep[see Table \\ref{ytot} and the $Z=10^{-5}$ model presented in][]{MEM06}. Since $^{22}$Ne is also a neutron source, s--process calculations are also planned. The effects of magnetic fields \\citep{YL05,WH06,ROTBIII} on the results will be studied in the near future. The dependence of the mass loss rates on the metallicity, especially in the RSG stage need to be further studied to see how the results of \\citet[mass loss in the RSG phase independent of metallicity]{VL05} can be extrapolated to very low metallicities."
    },
    "0608/astro-ph0608036_arXiv.txt": {
        "abstract": "Simple theoretical calculations have suggested that small body impacts onto Pluto's newly discovered small satellites, Nix and Hydra, are capable of generating time-variable rings or dust sheets in the Pluto system. Using HST/ACS data obtained on 2006 February 15 and 2006 March 2, we find no observational evidence for such a ring system and present the first constraints on the present-day I/F and optical depth of a putative ring system. At the 1500-km radial resolution of our search, we place a 3$\\sigma$ upper limit on the azimuthally-averaged normal I/F of ring particles of 5.1x10$^{-7}$ at a distance of 42,000~km from the Pluto-Charon barycenter, the minimum distance for a dynamically stable ring \\citep{Sternetal94, Nagyetal06}; 4.4x10$^{-7}$ at the orbit of Nix; and 2.5x10$^{-7}$ at the orbit of Hydra. For an assumed ring particle albedo of 0.04 (0.38), these I/F limits translate into 3$\\sigma$ upper limits on the normal optical depth of macroscopic ring particles of 1.3x10$^{-5}$ (1.4x10$^{-6}$), 1.1x10$^{-5}$ (1.2x10$^{-6}$), 6.4x10$^{-6}$ (6.7x10$^{-7}$), respectively. Were the {\\it New Horizons} spacecraft to fly through a ring system with optical depth of 1.3x10$^{-5}$, it would collide with a significant number of potentially damaging ring particles. We therefore recommend that unless tighter constraints can be obtained, {\\it New Horizons} cross the putative ring plane within 42,000~km of the Pluto-Charon barycenter, where rings are dynamically unstable. We derive a crude estimate of the lifetime of putative ring paritcles of 900 years. ",
        "introduction": "The discovery of Pluto's two small ($\\sim$100~km diameter) satellites, Nix and Hydra, in 2005 \\citep{Weaveretal06} raised the possibility that Pluto may possess a time-variable ring system \\citep{Sternetal06}. Previously, \\cite{Durda:stern00} demonstrated that collisions between small Kuiper belt debris and larger objects, such as Pluto's satellites, are a common occurrence. The characteristic ejecta velocity resulting from these collisions will be of order 1--10\\% of the impactor velocity, or 10--100~m s$^{-1}$. Since the escape velocities of Nix and Hydra are between 30 and 90~m s$^{-1}$, depending on their precise size and density, ejecta from such collisions can escape from Pluto's small satellites but will generally remain gravitationally bound to Pluto, thus forming rings. This is in contrast to the situation at Charon, where most such ejecta will fall back onto the surface, due to Charon's $\\sim$500~m s$^{-1}$ escape velocity. Such rings will be highly time-variable, as collisions on Nix and Hydra compete with loss processes. \\cite{Sternetal06} calculated a characteristic ring optical depth of $\\tau$=5x10$^{-6}$ for ring material between Nix and Hydra. This is considerably more tenuous than the main rings of Saturn, Uranus, and Neptune, but comparable to the optical depth of Jupiter's rings \\citep{Burnsetal84} and the faint dust rings of Uranus \\citep{Showalter:lissauer06}.  No prior observational constraints on the reflectance or optical depth of rings in the Pluto system have been published. We therefore undertook a search for such material using existing datasets acquired by the \\hst Advanced Camera for Surveys (ACS) in early 2006. In addition to the discovery potential of this search, detection of or constraints on Plutonian rings can yield useful information about both the potential hazards of ring material to the NASA {\\it New Horizons} Pluto mission, now en route, and the small body population orbiting in the 30--50 AU region of the Kuiper belt traversed by Pluto.  In what follows we describe the dataset we used to constrain the amount of ring material in the Pluto system, the analysis techniques we employed, and the results we obtained. We then go on to interpret our results, which are upper limits, as they apply to ring system hazards for {\\it New Horizons} and the lifetime of ring particles. ",
        "conclusions": "We have used existing \\hst ACS observations of the Pluto system to derive the first constraints on the normal I/F and optical depth of rings having a radial extent greater than 1500~km. We find a 3$\\sigma$ upper limit of $I/F_{_\\perp}$=5.1x10$^{-7}$ and, assuming a macroscopic ring particle albedo of 0.04 (0.38), a 3$\\sigma$ upper limit $\\tau_{_\\perp}$=1.3x10$^{-5}$ (1.4x10$^{-6}$). Although higher ring optical depths are possible for tightly confined rings less than 1500~km wide, this is unlikely given the apparent lack of additional small satellites to shepherd the rings. The present-day optical depth limits on the rings of Pluto are usefully constraining, but they remain a factor of three larger than the characteristic ring optical depth of 5x10$^{-6}$ predicted by \\cite{Sternetal06}. We note that the optical depth constraint derived from the \\hst observations is not sufficient to ensure safe passage of the {\\it New Horizons} spacecraft through the ring plane at distances of 42,000--65,000~km. Finally, in the case of a ring in steady state, we estimate a ring particle lifetime of 900 years."
    },
    "0608/astro-ph0608200_arXiv.txt": {
        "abstract": "{The detection of magnetic chemically peculiar (CP2) stars in open clusters of extragalactic systems can give observational answers to many unsolved questions. For example, one can study the influence of different global as well local environments on the lack of and presence of peculiarities.} {The mean percentage of CP2 stars in the Milky Way is of the order of 5\\% for the spectral range from early B- to F-type, luminosity class V objects. The origin of the CP2 phenomenon seems to be closely connected to the overall metallicity and global magnetic field environment. The theoretical models are still only tested by observations in the Milky Way. It is therefore essential to provide high quality observations in rather different global environments.} {The young clusters NGC 2136/7 were observed in the $\\Delta a$ photometric system. This intermediate band photometric system samples the depth of the 520\\,nm flux depression by comparing the flux at the center with the adjacent regions with bandwidths of 11\\,nm to 23\\,nm. The $\\Delta a$ photometric system is most suitable for detecting CP2 stars with high efficiency, but is also capable of detecting a small percentage of non-magnetic CP objects. Furthermore, the groups of (metal-weak) $\\lambda$ Bootis, as well as classical Be/shell stars, can be successfully investigated.} {We present high precision photometric $\\Delta a$ observations of 417 objects in NGC 2136/7 and its surrounding field, of which five turned out to be bona fide magnetic CP stars. In addition, we discovered two Be/Ae stars.} {From our investigations of NGC 1711, NGC 1866, NGC 2136/7, their surroundings, and one independent field of the LMC population, we derive an occurrence of classical chemically peculiar stars of 2.2(6)\\% in the LMC, which is only half the value found in the Milky Way. The mass and age distribution of the photometrically detected CP stars is not different from that of similar objects in galactic open clusters.} ",
        "introduction": "The classical chemically peculiar (CP) stars of the upper main sequence (luminosity classes V and IV) are targets of detailed investigations since their first description by Maury (1897). They provide excellent test objects for astrophysical processes like diffusion, convection, and stratification in stellar atmospheres in the presence of rather strong magnetic fields. These mechanisms for stars, both in the Milky Way and its surrounding stellar systems, contribute important knowledge to stellar evolution under different local circumstances, e.g., different metallicities.  CP stars can be detected very efficiently by applying the $\\Delta a$ photometric system (cf. Paunzen et al. 2005a), which measures the characteristic broadband absorption feature located around 520\\,nm, which is most certainly a consequence of the non-solar elemental abundance distribution of CP and related objects in the presence of a strong stellar magnetic field (Kupka et al. 2004). It samples its depth, comparing the flux at the center (521\\,nm, $g_{\\rm 2}$), with the adjacent regions (503\\,nm, $g_{\\rm 1}$ and 551\\,nm, $y$), using bandwidths from 11\\,nm to 23\\,nm. The respective index $a$ was introduced as $$ a = g_{\\rm 2} - (g_{\\rm 1} + y)/2. $$ The intrinsic peculiarity index $\\Delta a$ is defined as the difference between the individual $a$-values and those (=\\,$a_{\\rm 0}$) of non-peculiar stars of the same color. The locus of the $a_{\\rm 0}$-values has been called, the normality line.  Virtually all peculiar objects with magnetic fields (CP2 and CP4 stars, Preston 1974) have positive $\\Delta a$ values in excess of +100\\,mmag whereas Be/Ae/shell and metal-weak (e.g., $\\lambda$ Bootis group) stars exhibit significantly negative ones up to $-$35\\,mmag (Paunzen et al. 2005a). In general, some spectroscopic binary systems can mimic a positive $\\Delta a$ value of up to +16\\,mmag, which is well below the observed values for the stars presented in this paper. Later type, evolved objects of luminosity classes III to I might also show the same behavior, but can be easily identified within the color-magnitude diagram and do not influence the overall statistics.  In this paper we present our efforts to detect chemically peculiar stars in the field of NGC 2136/7, a binary cluster system, located in the Large Magellanic Cloud (LMC). In total, we observed 417 objects of which five turned out to be bona fide magnetic CP stars. These observations, added to our statistical analysis in the LMC presented in Paunzen et al. (2005b), show an overall occurrence of 2.2(6)\\% for chemically peculiar stars in the LMC. ",
        "conclusions": "We present the results of our search for CP stars of the upper main sequence in the LMC clusters NGC2136/7 and the surrounding field applying the intermediate band $\\Delta a$ photometry, which measures the characteristic broadband absorption feature located around 520\\,nm.  In total, 417 objects were measured on 69 individual frames observed with two different telescopes. We report the detection of five classical chemically peculiar objects and two Be/Ae stars. In addition, three peculiar objects were found that deserve further attention.  We conclude from our investigations of NGC 1711, NGC 1866, NGC 2136/7, their surroundings, and one independent field of the LMC population that the occurrence of classical chemically peculiar stars is 2.2(6)\\% for chemically peculiar stars in the LMC.  The age and mass distributions, derived by applying appropriate isochrones, do not alter from those of CP stars in galactic open clusters.  This provides a valuable observational source for understanding the CP phenomenon of the upper main sequence in a different global environment than in our Milky Way."
    },
    "0608/astro-ph0608214_arXiv.txt": {
        "abstract": "The main objective of the present work is to ckeck if the star formation efficiency plays a relevant role in the evolution of the relative abundance N/O.  In order to explore this idea, we analyze the evolution of the nitrogen-to-oxygen ratio as predicted by a set of computed theoretical models. These models consist of simulated galaxies with different total masses which are evolved assuming different collapse time scales and different star formation efficiencies. The combinations of these two parameters produce different star formation histories, which in turn have, as we show, an important impact on the resulting N/O ratio.  Since we want to check the effect of variations in these efficiencies on our models results, the same stellar yield sets are used for all of them. The selected yields have an important contribution of primary nitrogen proceeding from low and intermediate mass stars, which implies that N is ejected with a certain delay with respect to O. It allows to obtain, as we demonstrate, a dispersion of results in the N/O-O/H plane when star formation efficiencies vary which is in general agreement with observations.  The model results for the N/O abundance ratio are in good agreement with most observational data trends.  In particular, the behavior shown by the extragalactic H{\\sc ii} regions is well reproduced with present time resulting abundances. Furthermore, the low N/O values estimated for high-redshift objects, such as those obtained for Damped Lyman Alpha (DLA) galaxies, as well as the higher (and constant) values of N/O observed for irregular and dwarf galaxies or halo stars, can be simultaneously obtained with our models at the same low oxygen abundances $\\rm 12+log(O/H)\\sim 7$.  We therefore conclude that, even though there seems to be a general believe that abundance ratios depend mostly on stellar yields, these are not the only parameter at work when both elements are ejected by stars of different mass range, and that differences in the star formation history of galaxies and regions within them are a key factor to explain the data in the N/O-O/H plane. ",
        "introduction": "The nitrogen abundances have posed an important debate which lasts up to now. First, because it is created in two types of stars, what rises the doubt about which proportion is produced in massive stars and how much N is ejected by low and intermediate mass (LIM) stars. Second, the possible primary or secondary origin of nitrogen implies some questions about what is the contribution of each type produced in each star. The question is then how much N is created in each stellar mass range and what fraction of it is primary?  From an observational point of view, the need for a primary component of N appeared very soon.  \\citet{edm78} and \\citet{allo79} already concluded, from the analysis of the then available extragalactic H{\\sc ii} region data, that N behaves, at least partially, as a primary element. This was later supported by N measurements for Galactic metal-poor halo stars \\citep{bar83,tom84,lai85,car87}, and by additional extragalactic H{\\sc ii} region observations \\citep{mcc85,vila93}.  The number of N data points has increased extraordinarily since then \\citep[see][and references therein]{hen00,con02} showing the same basic trend: 1) a strong slope for the N/O data against oxygen abundance for metal-rich regions, a well explained behaviour by the secondary character expected from the CNO cycle production, and 2) a flat line for the data of low mass and irregular galaxies and halo stars with $\\rm 12 + \\log{(O/H)}<8$. Some very recent observations of halo stars \\citep{isr04,spi05} have increased the number of data in the low abundance end. They show a large scatter in N/O at low metallicities, and some high N/O ratios, implying an important primary N production at very early evolutionary times.  Since low metallicity galaxies have been considered to be young objects undergoing their first burst of star formation \\citep{thuan95}, it was assumed that they have had no time for their LIM stars to evolve and eject the nitrogen produced by them. Moreover, the dispersion found for their N/O abundances was really low. This way, a number of authors have claimed that massive stars are the main primary N producers. There are two facts against this argument: 1) these galaxies are not so young as it was thought, since recent observations indicate that they host stellar populations which are at least $10^{7}-10^{8}$ yr old, reaching even 1 Gyr is some cases \\citep{leg00,tols03,cai03,vzee04,thuan05}; 2) if LIM stars with masses in the range 4--8 $\\rm M_{\\odot}$ eject some amount of primary N, this can be in the interstellar medium (ISM) within a time as short as $\\sim$ [50 -- 200] Myr, since the mean-lifetimes of these stars are in this range. Moreover, low oxygen abundance do not necessary imply short time-scales. In fact, we claim that it is essential to change the usual scheme that identify oxygen abundances with evolutionary time, a misunderstander concept that lead to erroneous conclusions.  On the other hand, there is no clear mechanism which may produce primary N in massive stars. Some authors working on stellar evolution have searched for new possibilities, the most plausible one being the effect of rotation on stars.  \\citet{mey02} have computed some yields from stellar models including rotation obtaining that intermediate and massive stars may produce primary N, mostly at very low metallicity. However, \\citet{chia05} have included these yields in a Galactic chemical evolution (GCE) model finding that an extra-- production of N in low metallicity massive stars by a large factor, between 40 and 200 along the mass range, is still necessary to explain the data for very metal-poor halo stars since the yields used do not produce a sufficient amount of primary N. Recently, \\cite{chia06} use new (and still unpublished) yields for metal-poor massive stars which rotate rapidly thus producing large amounts of primary nitrogen with which it is possible to increase the N/O abundance up to similar levels of the Galactic halo stars. These preliminary yields are still uncertain since they are computed assuming a non-observed very high rotation velocity for these low-metallicity stars. \\footnote{In fact, as \\cite{hir05} stated, {\\sl the ratio $v_{ini}/v_{c}$ decreases at very low metallicity for the initial velocity of 300 km $^{-1}$ because stars are more compact at lower metallicity}.}.  The situation has become more complicated because abundances of N for DLA galaxies have been derived \\citep[][and references therein]{pet02,pro02,cen03}. When these recent observations are included in the sample, a large dispersion appears, in evident disagreement with the hypothesis of all the primary N being produced by massive stars.  DLA data show lower values of N/O than those, populating the flat slope, obtained for the irregular and blue compact galaxies and for the halo stars of similar oxygen abundance. These high redshift objects have oxygen abundances which correspond to slightly evolved objects as those observed in the local universe, but in this case some values of N/O are much smaller than the local ones. This fact appears incompatible with primary N being produced by massive stars, as those very metal-poor ones rotating at very high velocity used in \\cite{chia06}, since in that case it results difficult to conceive which mechanism could produce a different final yield (and therefore a different point on the N/O-O/H plane) for each DLA object.  From the theoretical point of view, a secondary production is expected in most stars, as corresponds to the CNO cycle, while the primary contribution should arise from intermediate mass stars ($4 \\leq M/M_{\\odot} \\leq 8$).  These stars may suffer, during the Asymptotic Giant Branch (AGB) phase, dredge-up episodes and hot bottom burning processes, due to which some primary N may be created. This fact has been well known for some decades \\citep{rv81,ser86}. In fact, most of yield sets \\citep[as][]{vhoek97, mar98} presently used for the chemical evolution models include a contribution of primary N produced by these processes in LIM stars. The AGB phase is still far from being well understood, and there exist large uncertainties about the exact amount of primary nitrogen produced through the HBB process. Probably due to that reason, the up to now computed models do not seem be able to reproduce most of the data.  The idea that a constant N/O may only be obtained with a primary N from massive stars is indeed an oversimplification which does not take into account the stellar lifetimes.  \\citet{vila93} and \\citet{pil03}, by analyzing the N/O values in galaxies of different morphological type, concluded that a long-time-delayed contribution to the N production must exist. \\citet{vila93} suggested that most data fall in a region limited by two extreme closed box models defined by distinct delays, produced by different star formation efficiencies, and/or ages of galaxies, thus explaining the observed dispersion.  This way, \\citet{hen00} and \\citet{pran03} --hereinafter HEN and PRAN, respectively -- computed realistic chemical evolution models using yields from \\citet{vhoek97} for LIM stars, with a contribution of primary N. Although their models produce a level of N/O higher than observations for Z low, they show that, actually, a N/O {\\sl vs} O/H relation much flatter than usual may be obtained with their models, only decreasing the efficiencies of star formation. The same finding was also obtained by \\cite{lar01} using several sets of yields, all of them with a proportion of primary N ejected by LIM stars. These results lead to the idea that stellar yields are not the only driver of the N/O ratio, at least when LIM stars produce a quantity of primary N, and that most observations in the plot N/O-O/H might be reproduced if different star formation rates in the different galaxies or regions are considered  \\citet{gav05} and \\citet[][hereinafter GAV05 and GAV06, respectively]{gav06} have recently calculated some chemical evolution models for the Milky Way Galaxy (MWG), by using different sets of stellar yields. They demonstrated in those works that a model using the set of yields for LIM stars from \\citet{gav05} joined to the massive stellar yields from \\cite{woo95} is able to reproduce most of trends of MWG data.  In the cited model, a flat behaviour of N/O {\\sl vs} O/H appears when the star formation rate is low during a long time, as it occurs in the Galactic halo or in the the outer regions of the Galactic disc. In this case, since the evolution is slow, a time longer than 40 Myr is necessary to reach oxygen abundances $12 +log(O/H) \\sim 7$, which is sufficient for LIM stars to eject their products to the ISM. This, combined with low O/H abundances, produces high N/O values, similar to those observed, for low abundances.  Most of the data of the Galaxy halo and disc, in particular those of nitrogen --see their Fig.12--, are successfully reproduced with that model.  Furthermore, the model results for the radial regions other than the Solar Neighbourhood suggest that actually the dispersion observed in MWG data can be a consequence of different star formation histories (produced by different collapse timescales to form the disc and different efficiencies to form stars) in the Galactic radial regions. This supports the mentioned findings from the above cited authors: different evolutionary tracks and end points in the O/H-N/O plane are possible simply by increasing or decreasing the star formation rate, although other parameters probably also play a role.  The objective of this work is, therefore, to explore if different star formation efficiencies may produce different tracks in the N/O-O/H plane and to check, with a wide grid of chemical evolution models, if the results are acceptable as compared with the observational data trends. In order to do this task we need a large number of models with the same basic assumptions about the evolution and where only the star formation be varied.  Therefore, we analyze the results obtained for the evolution of nitrogen with the wide grid of models shown in \\citet{mol05}, which simulates 44 galaxies with different total masses and 10 possible star formation efficiencies. The chosen sets of yields used for these computations were those from WW for massive stars and those from GAV05 for LIM stars, which were already calibrated as explained above.  An advantage of using these yields is that the primary N proceeds from LIM stars with which the effect of variable star formation is probably more evident. \\footnote{Since the objective of this work is to explore if different star formation efficiencies produce different tracks in the N/O-O/H plane we must use the same ingredients for all our realizations, except the parameter we want to check. This is similar to the usually applied method when the stellar yields are the free parameters. In that case using the same ingredients, except the stellar yields, for some Galactic chemical evolution (GCE) models, it is possible to determine what set of yields is the most adequate.}  In Sect.~2 we analyze the N/O-O/H results obtained with simple chemical evolution models, by describing how elemental abundances of N and O evolve. In Sect.~3 we summarize the resulting abundances for the multiphase chemical evolution model with a contribution of primary nitrogen proceeding from LIM stars variable with Z.  Sect.~4 is devoted to the results obtained when the same code is applied to a wide grid of simulated galaxies with different masses and star formation efficiencies.  Our conclusions are in Sect.5. ",
        "conclusions": "The evolutionary track followed by a given region or galaxy in the N/O {\\sl vs} O/H plane depends strongly on the star formation history of the region. Strong bursting star formation histories would produce high oxygen abundances soon and, hence, an early secondary behaviour, thus reproducing most of the spiral H{\\sc ii} region observations. On the contrary, a low and continuous star formation rate keeps the oxygen abundance low for a long time; thus a large quantity of primary N may be ejected reproducing the flat slope shown by dwarf galaxy data.  The resulting dispersion of our realizations, similar to the one observed, appears as a consequence of the different star formation histories in each modeled region and in each galaxy. These are obtained by combining various star and molecular cloud formation efficiencies and collapse time scales.  The stellar yields for LIM also play a role: if they have the adequate contribution of primary N depending on metallicity, it is easy to reproduce adequately the generic trend of the observational points in the N/O {\\sl vs} O/H plane, in particular a flat slope for low metallicity regions and, simultaneously, lower values of N/O at similar oxygen abundances for regions or galaxies where the star formation is strong and occurred early in the evolution.  The combination of the new yield set from \\citet{gav05} and the self-consistent chemical evolution models by \\citet{mol05} are able to reproduce the data corresponding to H{\\sc ii} regions in different galaxies and Galactic stars of different ages. This arises naturally just by assuming different input parameters (mostly star formation efficiencies and infall rates) for the different galaxies and regions in the same generic model, without the need to appeal to {\\sl ad hoc} models to reproduce each kind of object.  They may also to reproduce the existing data on galaxies at 0 $\\leq$ z $\\leq$ 4, and predict the existence of a {\\sl second plateau}, found in the N/O {\\sl vs} O/H relationship for $\\rm log(N/O) \\sim -2.5$ and galaxies with $z\\ge 3$. Our models predict that objects at redshifts $z \\le 2.5$ can not be discriminated from $z=0$ objects in the log(N/O) {\\sl vs} log(O/H) plane; however, for $z\\ge 3$ objects, this discrimination is possible.  This characteristic is a consequence of the rapid passage of model results by this zone of the plane, when the time evolution is considered, since 300 Myr after the beginning of the formation of stars, the value log(N/O) increases from -2.5 to -1.6 with oxygen abundances $\\rm 12+log(O/H)=7-8.5$"
    },
    "0608/astro-ph0608678_arXiv.txt": {
        "abstract": "The lack of an observed supernova associated with GRB 060614 appears to require a new paradigm for the formation of (a subset of) long-duration GRBs. This requirement is based on the presumed low redshift of the burst, which was inferred from the spatial coincidence of the afterglow with a $z=0.125$ galaxy. We explore the possibility that this low-redshift galaxy is a chance superposition along the line of sight to GRB 060614. We examine the galaxy distribution of the field of GRB 060614 and find that the probability of a chance association with a galaxy at least as bright as the putative host is only $\\sim 0.5 - 1.9$\\%. However, for the current ensemble of $\\approx 180$ \\textit{Swift} GRBs it is likely that several such coincidences have occurred, and given the ``non-standard\" nature of GRB 060614 it is not implausible that this is one such occurrence. Thus the conclusion that GRB 060614 requires a revision to the formation paradigm for long-duration GRBs should be approached with caution. ",
        "introduction": "The evidence that long-duration GRBs are associated with supernovae (SNe) is very strong. From observed low-luminosity GRBs with SNe that dominate the bursts' optical emission (e.g. GRB 980425/SN 1998bw and 060218/SN 2006aj; \\citealt{Galama+98,Cobb+06,Ferrero+06,Mirabal+06,Modjaz+06,Sollerman+06}) to high-luminosity GRBs with SNe embedded in the bursts' optical afterglow (e.g. GRB 030329/SN 2003dh; \\citealt{Hjorth+03,Stanek+03,Bloom+04}) has emerged a paradigm in which long-duration GRBs are produced by the core collapse supernovae of  massive stars.  The study of this association is difficult at high redshifts, due to the intrinsic faintness of the SNe when compared with the luminosities produced at all wavelengths by ultra-relativistic GRB jets.  Any possible low-redshift GRB is, therefore, of great importance and is scrutinized by many observers for photometric or spectroscopic evidence of a SN. Such was the case for GRB 060614 \\citep{Parsons+06}, a burst that was typical in most respects and was followed by a relatively bright and long-lasting optical afterglow \\citep{Holland06,Brown+06}.  Initial observations of GRB 060614 detected afterglow light in all \\textit{Swift} UVOT bands, indicating that the GRB had a moderate redshift of $z<1.7$, due to the absence of absorption from the Lyman-$\\alpha$ forest \\citep{Holland06,SX06}.  From the spectral parameters of the GRB, \\cite{PA06} calculated the pseudo-redshift of GRB 060614 to be $pz= 1.45\\pm0.85$, consistent with the earlier estimate \\citep[see also][]{SX06}. Spectroscopic observations of the afterglow of GRB 060614 taken less than two days post-burst did not yield a spectroscopic redshift \\citep{Fugazza+06a}, as no absorption or emission features were detected on top of the power-law spectrum of the afterglow.  More spectra were obtained at the afterglow position of GRB 060614 as the afterglow continued to fade away. In a spectrum taken $\\sim5$ days post-burst \\citep{Price+06}, a single strong emission line was detected and interpreted to be due to H$\\alpha$ from the GRB host galaxy at $z=0.13$. The redshift of this potential host galaxy was confirmed to be $z=0.125$ by \\cite{Fugazza+06b}. With the possibility of this burst being at very low redshift, several observers continued careful follow-up campaigns \\citep[e.g.][]{BH06,CB06,Fynbo+06a}, anticipating the brightening expected from a SN at that redshift. However, in time, all decaying afterglow light curves leveled off to a constant magnitude, interpreted to be the magnitude of the host galaxy, and no SN brightening was observed \\citep{CB06,DellaValle+06,Fynbo+06b,Gal-Yam+06}.  The lack of an observed SN from this seemingly low-redshift GRB can be explained in several ways: \\begin{enumerate} \\item A SN did occur but was not detected due to high line-of-sight extinction toward the GRB, most likely due to dust in the GRB's host galaxy.  However there are no indications of the presence of a supernova in IR observations \\citep{CB06} and also no indication of the presence of this reddening in the UV/optical/IR observations of the afterglow \\citep{Holland06,DellaValle+06,Fynbo+06b,Gal-Yam+06}. \\item A SN did occur but was not detected because the SN was underluminous by more than 5 magnitudes.  This would be unexpected, however, as all previous GRB-related SNe have luminosities that cluster within a magnitude of the peak brightness of SN 1998bw \\citep{ZKH04,Cobb+06,Ferrero+06}. While local non-GRB Type Ibc SNe do span a larger range of $\\pm\\sim2$ magnitudes from the peak of SN 1998bw \\citep{Richardson+06}, even the dimmest known Ibc SN should have been detectable in the absence of strong reddening. \\item No stellar core collapse occurred, or the core collapse did not result in a typical GRB ``hypernova\". These possibilities have now been suggested by several authors \\citep{DellaValle+06,Fynbo+06b,Gal-Yam+06} and require a paradigm shift in our understanding of the formation mechanism of GRBs or, at the very least, require the introduction of a new class of GRBs. \\item The proposed host galaxy of GRB 060614 is not the true host galaxy but, instead, represents the chance coincidence of a low-redshift galaxy intersecting the line of sight toward what is actually a moderate-redshift GRB (as suggested by \\citealt{SX06}). The SN, therefore, did occur but is undetectable given the redshift of the burst. \\end{enumerate}  In some ways, the fourth possibility is the most appealing, since it requires no new formation channel for long-duration GRBs, and is consistent with some features of the GRB itself \\citep{SX06}. However, it requires a spatial coincidence between a foreground galaxy and the GRB.  Here we analyze a deep optical image of the field of GRB 060614 to study the statistics of such a coincidence empirically.  In agreement with \\cite{SX06}, we find that a chance superposition is not implausible. While statistical arguments can never rigorously exclude or require such a coincidence, the drastic nature of the alternatives suggests that the possibility of a chance superposition should be seriously considered. ",
        "conclusions": "Although the probability for a chance superposition for any given GRB is small, we must also consider that a large number of GRB fields have now been observed in this manner.  If we assume that the galaxy density in the field of GRB 060614 is not unusual (as appears to be the case --- Cobb et al. in prep), then given the $\\approx 180$ bursts observed by \\textit{Swift}, a chance overlap of a galaxy brighter than $I\\sim22$ and the precise position of a GRB would be expected for $\\approx 1 - 4$ GRBs. The probability that no chance overlaps would be observed in the ensemble is quite small, at 3.2\\% to 41\\%, assuming a probability of 1.9\\% to 0.5\\% for a chance superposition in any given image.  Most such superpositions would not be noticed, however, because the typical foreground galaxy would be at a redshift comparable to that of the GRB.  In contrast, GRB 060614 has generated great interest because the galaxy implies that the GRB is at an unusually low redshift.  The strongest host galaxy claims can only be made for GRBs with detected optical afterglows.  Of the $\\sim80$ optical afterglows detected following \\textit{Swift} bursts, the detection of no chance overlaps would not be unexpected at a probability of 22\\% to 67\\%.  Note, however, that the position of GRBs without optical afterglows are determined from \\textit{Swift} XRT observations of X-ray afterglows, which have a spatial precision of $\\sim4\"$. Taking this larger positional uncertainty into account would significantly increase the probability of chance superpositions.  One additional feature of GRB 060614 is that it appears to be at the periphery of the galaxy \\citep{Gal-Yam+06}.  This is not expected for standard long-duration GRB formation mechanisms, since the star formation rate at the outskirts of galaxies is expected to be low.  Indeed \\cite{Gal-Yam+06} invoke the location of the GRB within the galaxy as a first clue to the nature of a new GRB formation paradigm.  The chance superposition hypothesis is consistent with an off-center location of the GRB. With deeper observations it may be possible to confirm or refute the presence of a background galaxy, although we note that there is a $\\sim1/4$ chance that the center of the true host galaxy will lie directly behind that of the $z=0.125$ galaxy \\citep{Bloom+02}.  We note that any superposition with a foreground galaxy requires that the source be lensed.  However, the low luminosity and hence inferred low mass for this particular galaxy suggests that the lensing induced magnification is less than 5\\%, which is inconsequential.  This lensing estimate is obtained using standard assumptions and modeling the galaxy as a singular isothermal sphere.  It has been suggested that GRB 060505 is also a low-redshift GRB with no SN \\citep{Fynbo+06b}, although the observational data are more limited both in the gamma-rays \\citep{SX06} and for the afterglow, for which no optical spectrum has been reported.  Clearly the existence of a second ``060614-like\" object would significantly diminish the probability of chance superpositions.  But this argument requires that the two objects be members of the \\textit{same} putative new category, and it is not clear that this could be the case.  In particular, there are significant differences in the GRB characteristics of the two sources (GRB 060505 was much shorter--4 s versus 120 s--and weaker than GRB 060614; \\citealt{Palmer+06,Parsons+06}) and GRB 060505 appears to be located in a star forming region of a galactic spiral arm, whereas GRB 060614 seems to be outside the main region of star formation.  If, as seems likely, the inferred low redshifts of these two objects require different formation mechanisms from each other, then an explanation of either or both of them as a chance superposition remains plausible.  We emphasize that an analysis of this kind is by nature not definitive. From this argument alone, it is not possible to claim that GRB 060614 \\textit{is} a chance superposition between a low-$z$ galaxy and a moderate-$z$ GRB; we only claim that there is a possibility of such an occurrence. The fact that galaxies as bright as the purported host of GRB 060614 cover several percent of the field, and that well over a hundred GRBs have now been observed by \\textit{Swift}, suggests that a chance superposition cannot be excluded.  Indeed, over the ensemble of \\textit{Swift} bursts, the probability that such a coincidence would be observed in at least one case is high. The unusual combination of a low-redshift putative host and the lack of an associated supernova make GRB 060614 an intriguing candidate for a chance superposition. Thus the conclusion that GRB 060614 requires a ``new paradigm\" for GRB formation should be approached with caution."
    },
    "0608/astro-ph0608352_arXiv.txt": {
        "abstract": "We have carried out a sensitive high-resolution imaging survey of stars in the young (6--8\\,Myr), nearby (97\\,pc) compact cluster around $\\eta$~Chamaeleontis to search for stellar and sub-stellar companions. Given its youth and proximity, any sub-stellar companions are expected to be luminous, especially in the near infrared, and thus easier to detect next to their parent stars. Here, we present VLT/NACO adaptive optics imaging with companion detection limits for 17 $\\eta$\\,Cha cluster members, and follow-up VLT/ISAAC near-infrared spectroscopy for companion candidates. The widest binary detected is $\\sim$0\\farcs2, corresponding to the projected separation 20\\,AU, despite our survey being sensitive down to sub-stellar companions outside 0\\farcs3, and planetary mass objects outside 0\\farcs5. This implies that the stellar companion probability outside 0\\farcs3 and the brown dwarf companion probability outside 0\\farcs5 are less than 0.16 with 95\\,\\% confidence. We compare the wide binary frequency of $\\eta$\\,Cha to that of the similarly aged TW~Hydrae association, and estimate the statistical likelihood that the wide binary probability is equal in both groups to be $< 2 \\times 10^{-4}$. Even though the $\\eta$\\,Cha cluster is relatively dense, stellar encounters in its present configuration cannot account for the relative deficit of wide binaries. We thus conclude that the difference in wide binary probability in these two groups provides strong evidence for multiplicity properties being dependent on environment. In two appendices we derive the projected separation probability distribution for binaries, used to constrain physical separations from observed projected separations, and summarize statistical tools useful for multiplicity studies. ",
        "introduction": "Most solar-type stars reside in binaries \\citep{duq91}, yet their formation is not well understood. It is generally accepted that most stars form in groups rather than in isolation \\citep{ada01}. Intriguingly, the multiplicity of young stars in many regions seems to be systematically higher than that of their main-sequence counterparts \\citep[e.g.][and references therein]{mat00}. The reason for this multiplicity overabundance is not entirely understood, although two principal scenarios have been proposed. One is that the multiplicity fraction is sensitive to the initial conditions of the star formation process. That would imply that different regions show a variety of multiplicity properties, and indeed, there seems to be some evidence for the multiplicity fraction being anticorrelated with the stellar density of the region \\citep{pat01}. Another possibility is that young stars start out with a high fraction of multiples that subsequently are disrupted due to dynamical evolution \\citep[e.g.][]{rei00}. Star-forming regions are generally too dispersed for binaries to be disrupted by interactions between members. The alternative is that many stars form not just in binaries, but in unstable higher-order multiples that get disrupted with time, as indeed suggested by numerical simulations \\citep{bat02,del04}. For a recent review on the current status of the field, see \\citet{duc06}.  The nearby (97\\,pc), young \\citep[6--8\\,Myr,][and references therein]{jil05} $\\eta$\\,Chamaeleontis cluster was found by the X-rays emitted from its members, as revealed by \\textit{ROSAT}, together with their common space motion, as revealed by \\textit{Hipparcos} \\citep{mam99}. The cluster is mainly populated by late-type stars (K3-M6, see Table~\\ref{t:targets}), and shows no evidence for extinction \\citep{mam00}. These properties all contribute to make the cluster an excellent laboratory for investigating brown dwarf (BD) and planet formation, and their evolution \\citep[e.g.][]{son04,luh04a,hai05,lyo06,jay06}. In particular, using adaptive optics (AO) systems on large telescopes, it is possible to reach high contrast-ratio sensitivities close in, to detect any wide ($>$0\\farcs5) companion down to planetary masses.  In a previous high-resolution survey using speckle interferometry and AO on 2.2--3.6\\,m telescopes, \\citet{koh02} searched 13 members for companions ($\\eta$\\,Cha~1--12 and 15, see Table~\\ref{t:targets}). They found two resolved binaries, $\\eta$\\,Cha~1 and $\\eta$\\,Cha~9, and one suspected unresolved binary, $\\eta$\\,Cha~12. $\\eta$\\,Cha~12 has also been suspected to be binary due to its elevation in color-magnitude diagrams \\citep{law01,luh04b}.  In the present paper we report on a deep AO search for faint companions to 17 confirmed members of the $\\eta$\\,Cha cluster. \\S\\,\\ref{s:obs} details the observations. \\S\\,\\ref{s:analysis} describes analysis and results, including how astrometry and photometry were performed, how the Strehl ratio was measured and the contrast sensitivity estimated. In \\S\\,\\ref{s:discuss} we discuss the companion candidates and constrain the orbit of the previously observed 0\\farcs2 binary $\\eta$\\,Cha~1 (Fig.~\\ref{f:ec01}). We then go on to determine limits on the companion probability, use this to discuss the apparent deficit of binaries in the $\\eta$\\,Cha cluster, and estimate the likelihood of binaries being disrupted by stellar encounters. \\S\\,\\ref{s:conclusions} contains an enumerated list of conclusions.  In the Appendices, we first derive the (time-averaged) probability distribution for the projected separation of a binary, given its semi-major axis. We then continue with some useful properties of multiplicity statistics, in particular how to estimate binomial confidence intervals, and how to test if two outcomes are derived from the same binomial distribution. ",
        "conclusions": "\\label{s:conclusions} We summarize our conclusions as follows:  \\begin{enumerate} \\item We found no new companions to stars in the $\\eta$\\,Cha cluster, despite being sensitive down to the star/BD limit outside 0\\farcs3 (30\\,AU) and down to the BD/planet limit outside  0\\farcs5 (50\\,AU).  \\item We have constrained the orbit of $\\eta$\\,Cha\\,1 further, but are unable to usefully constrain a dynamical mass. Resolved radial velocities of the two components, or a longer astrometric time baseline, are required for an accurate mass estimate.  \\item The 95\\,\\% upper limit for the wide ($>$30\\,AU) binary probability $w_{\\eta\\,\\mathrm{Cha}}$ in $\\eta$\\,Cha is $p_{95\\%}$=0.18. This contrasts to the wide binary probability $w_{\\mathrm{TWA}}=0.58^{+0.13}_{-0.14}$ of TWA. The likelihood that  $w_{\\eta\\,\\mathrm{Cha}} = w_{\\mathrm{TWA}}$ is less than $2\\times10^{-4}$.  \\item Multiplicity properties depend on the initial conditions of the formation environment.  \\end{enumerate}"
    },
    "0608/astro-ph0608487_arXiv.txt": {
        "abstract": "We report on the {\\it Hubble Space Telescope} program to observe periodic comet 9P/Tempel 1 in conjunction with NASA's Deep Impact mission. Our objectives were to study the generation and evolution of the coma resulting from the impact and to obtain wide-band images of the visual outburst generated by the impact.  Two observing campaigns utilizing a total of 17 \\hst\\ orbits were carried out: the first occurred on 2005 June 13--14 and fortuitously recorded the appearance of a new, short-lived fan in the sunward direction on June 14. The principal campaign began two days before impact and was followed by contiguous orbits through impact plus several hours and then snapshots one, seven, and twelve days later.  All of the observations were made using the Advanced Camera for Surveys (ACS).  For imaging, the ACS High Resolution Channel (HRC) provides a spatial resolution of 36 km (16 km pixel$^{-1}$) at the comet at the time of impact.  Baseline images of the comet, made prior to impact, photometrically resolved the comet's nucleus.  The derived diameter, 6.1 km, is in excellent agreement with the $6.0 \\pm 0.2$~km diameter derived from the spacecraft imagers. Following the impact, the HRC images illustrate the temporal and spatial evolution of the ejecta cloud and allow for a determination of its expansion velocity distribution.  One day after impact the ejecta cloud had passed out of the field-of-view of the HRC. ",
        "introduction": "The Deep Impact mission \\citep{A'Hearn:2005a,A'Hearn:2005b} successfully placed a 364 kg, about half copper, impactor, onto the surface of comet 9/P Tempel 1 at a relative velocity of 10.3~\\kms\\ on 2005 July 4 at 05:52:03 UT (as seen from Earth).  The event was observed by cameras aboard the mother spacecraft and by a large number of Earth- and space-based telescopes as part of an extensive campaign to study the comet prior to, during, and in the course of several days following the impact \\citep{Meech:2005}.  The {\\it Hubble Space Telescope} (\\hst) provided the highest spatial resolution images from Earth (36 km, 16 km pixel$^{-1}$) using the Advanced Camera for Surveys (ACS) High Resolution Channel (HRC), and also ultraviolet imaging using the Solar Blind Channel (SBC) of ACS to study the evolution of gaseous species, particularly CO, released by the impact.  This paper describes the \\hst\\ campaign and presents the images obtained and the preliminary interpretation of the data in terms of the physical properties of the observed ejecta.  We also compare the ejecta produced by the impact with that from a ``natural outburst'' that was serendipitously observed on 2005 June 14. ",
        "conclusions": "We have presented an overview of the visible imaging obtained with the Advanced Camera for Surveys High Resolution Channel on the {\\it Hubble Space Telescope} before, during and following the encounter of the Deep Impact spacecraft with comet 9P/Tempel~1.  In addition, we serendipitously observed a natural outburst of the comet on 2005 June 14, whose characteristics may be compared with those of the ejecta produced by the impact.  From the data we have obtained a measurement of the photometric size of the nucleus, determined the ejecta expansion velocity distribution, and followed the temporal evolution of the morphology of the visible brightness of the ejecta cloud.  The high spatial resolution and image quality of the HRC images should make possible detailed kinematical modeling of the ejecta plume that may serve to constrain the physical properties of the ejecta particles.  \\begin{center}{\\bf ACKNOWLEDGMENTS}\\end{center}  This work is based on observations with the NASA/ESA {\\it Hubble Space Telescope} obtained at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy (AURA), Inc., under NASA contract NAS 5-26555. We thank Ian Jordan, Ron Gilliland, and Charles Proffitt (STScI) for the planning and successful execution of the \\hst\\ program; Eddie Bergeron, Max Mutchler, Zolt Levay (STScI) and Ken Anderson (JHU) for the rapid response and production of properly corrected images; Cheryl Gundy, Lisa Frattare, Ray Villard, Mario Livio, and countless others at STScI for the July~4 logistics; and STScI/JPL for coordination of impact time before the end of \\hst\\ visibility.  This work was supported by grant GO-10144.01-A from the Space Telescope Science Institute."
    },
    "0608/astro-ph0608164_arXiv.txt": {
        "abstract": "We report on the first census of INTEGRAL/IBIS detections ($\\gtrsim 4\\sigma$ significance) above 100 keV based on the Core Program and public Open time observations up to April 2005. There are 49 sources detected in the 100-150 keV band of which 14 are also seen in the 150-300 keV range. The low energy sample is dominated by X-ray binary systems of both low and high mass, but also includes 10 active galaxies. Of the binary systems that are detected above 150 keV, more than 50\\% are associated with black hole candidates, often reported as microquasars. The present survey results are then used to construct LogN-LogS curves for galactic and extragalactic objects in the 100-150 keV band: above a 1 mCrab sensitivity limit we expect that around 200 galactic sources and almost 350 active galaxies populate the sky above 100 keV. While the contribution of individual point sources to the total Galactic emission has been estimated to be around 70-80$\\%$ between 100-300 keV, we find that active galaxies detected above 1 mCrab account for only about 3$\\%$ of the cosmic hard X-ray background in the 100-150 keV band. ",
        "introduction": "One of the primary goals of the INTEGRAL observatory \\citep{wink03}, and in particular of the IBIS imager \\citep{uber03}, is to provide a survey of the high energy sky ($>$20 keV). To this end a number of teams have made an effort to provide a systematic analysis of the data sets, as available to them at the time, so as to provide the general community with consolidated source information: so far 2 general surveys \\citep{bird04, bird06}, several surveys of specific areas of the sky \\citep{revn04a,revn06,molk04,lebr04} and a few AGN catalogues \\citep{beck06,bass06} have been published. A common factor throughout all of these activities is that the source search concentrated on the 18-100 keV regime where the IBIS sensitivity to point sources is optimal. Emission at higher energies has not been specifically investigated for or characterised in terms of source typology and emission mechanism except for a few objects. In this letter we focus our search above 100 keV and report the first census of the sky above this energy; as a result we provide for the first time the LogN-LogS relationship for galactic and extragalactic sources in the 100-150 keV band as well as an estimate of the galactic source contribution to the total Galactic emission and of the AGNs to the cosmic hard X-ray background. ",
        "conclusions": "A further use of the present catalogue is to estimate the contribution of galactic point sources to the total Galactic emission and that of AGN to the cosmic diffuse background in the energy bands explored in the present paper. The total galactic point source contribution has been estimated by summing the 100-150 keV fluxes from all sources detected within the central radius excluding all AGNs and Sco-X1 which is located at high galactic latitude. The flux obtained amounts to $2.17\\pm0.01\\times10^{-4}$ and $4.29\\pm0.10\\times10^{-5}$ ph cm$^{-2}$s$^{-1}$keV for 100-150 and 150-300 keV respectively; this implies a contribution of 83\\% and 68\\% in the two bands in agreement with previous and current INTEGRAL results \\citep{bouc05,lebr04}. The percentage contribution does not increase much as we go to fainter fluxes as sources a factor of 100 weaker will bring the point source component to 85\\% contribution. The above estimates further confirm previous indications that point sources largely dominate the total Galactic emission. In a similar way, the extragalactic logN-logS distribution derived from the present survey  provides a total AGN emissivity above 1 mCrab of $\\sim 0.003$ keV cm$^{-2}$ keV $^{-1}$ s$^{-1}$ sr$^{-1}$ in the  100-150 keV or approximately 3\\% of the total intensity of the extragalactic hard X-ray background in the same energy band \\citep{grub99}. Extrapolation of our LogN-LogS by a factor of 100 towards fainter fluxes will account for 15\\% of the extragalactic background above 100 keV; this value increases to 20\\% if we adopt a LogN-LogS slope of 1.5. In short we just start to uncover the source populations responsible for the hard X-ray background. As the INTEGRAL mission lifetime is extended (so far confirmed till 2010), more sky will be surveyed and more exposure will be accumulated resulting in an increase in the number of objects detected. At a limiting flux of 1 mCrab (which is likely reachable in the entire lifetime of the mission) around 200 galactic objects may be detected and more than 300 AGN identified at high energy."
    },
    "0608/astro-ph0608602_arXiv.txt": {
        "abstract": "One of the prime pieces of evidence for dark matter is the observation of large overdense regions in the universe. Since we know from the cosmic microwave background that the regions that contained the most baryons when the universe was $\\sim400,000$ years old were overdense by only one part in ten thousand, perturbations had to have grown since then by a factor greater than $(1+z_*)\\simeq 1180$ where $z_*$ is the epoch of recombination. This enhanced growth does not happen in general relativity, so dark matter is needed in the standard theory. We show here that enhanced growth can occur in alternatives to general relativity, in particular in Bekenstein's relativistic version of MOdified Newtonian Dynamics (MOND). The vector field introduced in that theory for a completely different reason plays a key role in generating the instability that produces large cosmic structures today. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608328_arXiv.txt": {
        "abstract": "The discovery of nearly 200 extrasolar planets during the last decade has revitalized scientific interest in the physics of planet formation and ushered in a new era for astronomy.  Astronomers searching for the small signals induced by planets inevitably face significant statistical challenges.  For example, radial velocity (RV) planet searches (that have discovered most of the known planets) are increasingly finding planets with small velocity amplitudes, with long orbital periods, or in multiple planet systems. Bayesian inference has the potential to improve the interpretation of existing observations, the planning of future observations and ultimately inferences concerning the overall population of planets. The main obstacle to applying Bayesian inference to extrasolar planet searches is the need to develop computationally efficient algorithms for calculating integrals over high-dimensional parameter spaces.  In recent years, the refinement of Markov chain Monte Carlo (MCMC) algorithms has made it practical to accurately characterize orbital parameters and their uncertainties from RV observations of single-planet and weakly interacting multiple-planet systems.  Unfortunately, MCMC is not sufficient for Bayesian model selection, i.e., comparing the marginal posterior probability of models with different parameters, as is necessary to determine how strongly the observational data favor a model with $n+1$ planets over a model with just $n$ planets.  Many of the obvious estimators for the marginal posterior probability suffer from poor convergence properties.  We compare several estimators of the marginal likelihood and feature those that display desirable convergence properties based on the analysis of a sample data set for HD 88133b \\citet{Fischer05}.  We find that methods based on importance sampling are most efficient, provided that a good analytic approximation of the posterior probability distribution is available.  We present a simple algorithm for using a sample from the posterior to construct a mixture distribution that approximates the posterior and can be used for importance sampling and Bayesian model selection.  We conclude with some suggestions for the development and refinement of computationally efficient and robust estimators of marginal posterior probabilities. ",
        "introduction": "Recent collaboration between astronomers and statisticians has led to a better understanding of the particular challenges associated with Bayesian analysis of dynamical planet detections.  In this paper, we briefly review the state of the art of Bayesian parameter estimation, model selection, and experimental design in the context of extrasolar planet searches.  Then, we discuss recent work on Bayesian model selection, demonstrating the properties of several estimators of the marginal posterior probability using an actual set of data from an RV planet search.  \\subsection{Observational Data}  In RV surveys, the velocity of the central star is precisely monitored for periodic variations which could be caused by orbiting companions (see Fig.\\ 1, left).  Each individual observation can be reduced to an estimate of the observational uncertainty ($\\sigma_{k}$) and a measurement of the star's RV ($v_k$) relative to the $j_k$th velocity reference at time $t_k$. Because each RV measurement is based on calculating the centroid of thousands of spectra lines and averaged over hundreds of sections of the spectrum, the observational uncertainties of most current echelle based RV surveys can be accurately estimated and are nearly Gaussian \\citep{Butler96}.  There may also be intrinsic stellar variability (``jitter'') that we model as an addition source of uncorrelated Gaussian noise with variance $s^2$ and add to the measurement uncertainties in quadrature. If the velocity observations ($\\vec{v}_{\\vec{\\theta}}$) were generated by the model specified by $\\mathcal{M}$ and model parameters $\\vec{\\theta}$, then the probability of measuring the observed velocities is \\begin{equation} L(\\vec{\\theta}) \\equiv p(\\vec{v}|\\vec{\\theta},  \\mathcal{M}) = \\prod_{k} \\frac{1}{\\sqrt{2\\pi\\left(\\sigma_k^2+s^2\\right)}} \\exp \\left[ -\\frac{\\left(v_{\\vec{\\theta}}(t_k,j_k)-v_k\\right)^2}{2\\left(\\sigma_k^2+s^2\\right)} \\right], \\label{EqnLikelihood} \\end{equation} assuming that the errors in individual observations and the ``jitter'' are both normally distributed and uncorrelated.  \\begin{figure}[!ht] \\plottwo{Ford_vels.eps}{Ford_pgram.eps} \\label{FigData} \\caption{Left: The 17 RV observations of HD 88133 published in \\citet{Fischer05} plotted versus the time of observation.  Right: The log posterior probability marginalized over all model parameters except the orbital period, assuming a single planet on a circular orbit (i.e., the Bayesian generalization of the periodogram) and the observational data shown on the left.} \\end{figure}   \\subsection{Theoretical Model}  Specifying the mass and six phase space coordinates of each body in a planetary system at a specified time provides a complete description of the system. In practice, it is convenient to choose the osculating Keplerian orbital elements (orbital period, $P$, orbital eccentricity, $e$, inclination relative to the plane of the sky, $i$, argument of periastron measured from the plane of the sky, $\\omega$, longitude of ascending node, $\\Omega$, and mean anomaly, $M$) for each planet in Jacobi coordinates, since the mean anomaly is the only one of these orbital elements that changes with time for a planet on an unperturbed Keplerian orbit. The observed stellar velocity is the sum of the line of sight velocity of the center-of-mass and the projection of the reflex velocity due to any planetary companions onto the line of sight.  For multiple planet systems, it can be important to use complete n-body simulations to model the planetary motions accurately (e.g., GJ876; Laughlin et al.\\ 2005; Rivera et al.\\ 2005).  However, in most cases, the mutual planetary perturbations are negligible on time scales comparable to the duration of observations.  In such cases, the RV perturbations due to a multiple planet system can be modeled as the linear superposition of multiple non-interacting Keplerian orbits, $ v_{\\vec{\\theta}}(t,j) = C_j + \\sum_p \\Delta v_{p}(t), $ where $C_j$ is the $j$th velocity reference.  While there is a single mean line of sight velocity of the center of motion, it is important to use separate constants, $C_j$ for each observatory/spectrograph pair, since the velocities are measured differentially relative to a reference velocity that is unique to each observatory.  The perturbation to the stellar RV ($\\Delta v_{p}$) due to a planet on a Keplerian orbit is given by $ \\label{rveqn} \\Delta v_{p}(t) = K_p \\left[ \\cos\\left(\\omega_p+T_p\\right) + e_p \\cos(\\omega_p) \\right] $ where $p$ labels the planet, $K$ is the velocity semi-amplitude and $T$ is the true anomaly, which implicitly depends on time.  The true anomaly ($T$) is related to the eccentric anomaly ($E$) via the relation $ \\tan \\left( T/2 \\right) = (1+e)^{1/2}(1-e)^{-1/2} \\tan^{1/2} \\left( E/2 \\right). \\label{EccTrueEqn} $ The eccentric anomaly is related to the mean anomaly ($M$) via Kepler's equation $ \\label{KeplerEqn} E(t) - e \\sin\\left(E(t)\\right) = M(t) - M_o = 2\\pi t /P - M_o $ where $M_o$ is a constant, the orbital phase at $t=0$. Unfortunately, RV observations alone are not sensitive to the orbital inclination relative to the plane of the sky ($i$) or the longitude of ascending node ($\\Omega$).  Therefore, RV observations by themselves can only measure the minimum mass ratio, $m_{\\min}/M_* = m \\sin i/M_*$, that is a function of $K$, $P$, $e$, and $M_*$, where $M_*$ is the stellar mass, typically estimated by independent astronomical observations.   \\subsection{Bayesian Framework}  To quantitatively analyze the available observational constraints, we employ the techniques of Bayesian inference.  By treating both the observation and the model parameters as random variables, Bayesian inference is able to address statistical questions in a mathematically rigorous fashion.  The joint probability, $p(\\vec{v}, \\vec{\\theta} | \\mathcal{M})$, can be expressed as the product of the likelihood ($L(\\vec{\\theta}) \\equiv p(\\vec{v} | \\vec{\\theta}, \\mathcal{M})$), the probability of the observables given the model parameters), and a prior probability distribution function ($p(\\vec{\\theta} | \\mathcal{M})$) which is based on previous knowledge of the model parameters.  Note that each of the probability distribution functions (PDFs) is conditioned on the assumption of a model, $\\mathcal{M}$, that includes the meaning of the model parameters, $\\vec{\\theta}$, and their relationship to the observational data, $\\vec{v}$.  Bayes's theorem allows one to compute a posterior probability density function, $p(\\vec{\\theta} | \\vec{v}, \\mathcal{M})$, which incorporates the knowledge gained by the observations $\\vec{v}$, \\begin{equation} p(\\vec{\\theta}| \\vec{v}, \\mathcal{M} ) = \\frac{ p( \\vec{\\theta} | \\mathcal{M} ) p(\\vec{v} | \\vec{\\theta}, \\mathcal{M} ) }{ \\int p( \\vec{\\theta} | \\mathcal{M} ) p( \\vec{v}| \\vec{\\theta}, \\mathcal{M} ) \\,d\\vec{\\theta} }. \\label{BayesEqn} \\end{equation} This paper is particularly interested in the case when there are multiple viable models (e.g., no planet model, one planet model, two planet model, etc.) for the current data set.  In this case, the posterior for a given model and set of parameters is given by $p(\\vec{\\theta}, \\mathcal{M} | \\vec{v} ) = p(\\mathcal{M}) p(\\mathcal{M}|\\vec{v}) p(\\vec{\\theta} | \\vec{v}, \\mathcal{M} )$, where $p(\\mathcal{M})$ is the prior probability of model $\\mathcal{M}$ and the marginal posterior probability for model $\\mathcal{M}$ is given by \\begin{equation} \\label{EqnMarginal} m(\\vec{v}) \\equiv p(\\mathcal{M} | \\vec{v} ) =  \\int p( \\vec{\\theta} | \\mathcal{M} ) p( \\vec{v}| \\vec{\\theta}, \\mathcal{M} ) \\,d\\vec{\\theta}. \\end{equation} We briefly review recent progress in sampling from $p(\\vec{\\theta}| \\vec{v}, \\mathcal{M} )$ via MCMC in \\S1.5 and introduce and compare algorithms for evaluating of $m(\\vec{v})$ in \\S2.  \\subsection{Choice of Priors}  In Table~\\ref{TabPriors}, we list the priors used for each model parameter in this work.  While the choice of these and other parameters can be fine-tuned for a given problem, we suggest these choices as a starting point for the Bayesian analysis of radial velocity data sets.  Physical and geometric considerations lead to natural choices for the prior PDFs for most of the model parameters (see Table~\\ref{TabPriors}).  A few of the priors merit closer attention.  The cutoff at a minimum orbital period is chosen to be less than the smallest orbital period of known extrasolar planets, but this is somewhat larger than the theoretical limit (the Roche limit occurs at $\\sim0.2$d for a 10$M_{\\mathrm{Jup}}$ planet around a 1$M_{\\odot}$ star).  The cutoff at a maximum orbital period is much longer than any known extrasolar planet, but is chosen to be roughly where perturbations from passing stars and the galactic tide would disrupt the planet's orbit.  The cutoff at a maximum velocity semi-amplitude is chosen to be a function of orbital period, so that the cutoff corresponds to a constant planet-star mass ratio.  In this paper, we set $K_{\\max} = 2129$m/s, which corresponds to a maximum planet-star mass ratio of 0.01.  This choice is primarily based on the observed distribution of extrasolar planet masses.  Clearly, there are stellar binaries with much larger velocity amplitudes, but this limit can be considered the definition of a planet.  While the possibility of arbitrarily small masses (and hence velocity amplitudes) prevents a physical justification for a lower cutoff, $K_{\\min}$, it is not possible to detect or constrain the orbital parameters of a sufficiently low-mass planet.  To keep the prior distribution normalizable, we impose breaks in the priors for $K$ and $s$ at $K_0$ and $s_0$, chosen to be at a velocity amplitude less than the smallest detectable velocity amplitude.  For a data set with $N_{\\mathrm{obs}}$ RV measurements with typical measurement precision of $\\bar{\\sigma}$, we suggest setting $K_0 \\le \\bar{\\sigma} \\sqrt{50/N_{\\mathrm{obs}}}$ (Cumming 1999).  For the reference prior, we chose $s_0 = K_0 = 1$m/s, somewhat arbitrarily, but motivated by the current state of the art in RV planet searches \\citep{N2K6}.  For systems where a planet is clearly detected, the shape of the posterior will not be sensitive to our assumptions about the prior for small values of $K$, but for planets which are marginally detected, this choice may become significant. If the posterior distribution has significant probability near $K\\simeq K_o$, then one should check how sensitive any conclusions are to the choice of $K_o$.  \\begin{table}[!ht] \\caption{SAMSI Exoplanet Working Group Reference Priors} \\smallskip \\begin{center} {\\small \\begin{tabular}{ccccc} \\tableline \\noalign{\\smallskip} Model Parameter & Variable & Prior Distribution & Minimum & Maximum \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} Amplitude of jitter & $s$ & $\\frac{(s+s_0)^{-1}}{\\log\\left(1.+\\frac{s_{\\max}}{s_0}\\right)}$ & 0 m/s & $K_{\\max}$ \\\\ \\smallskip\\\\ \\multicolumn{5}{c}{Parameters for each velocity reference} \\\\ Velocity offset & $C_j$ & $\\frac{1}{C_{\\max}-C_{\\min}}$ & -$K_{\\max}$ & $K_{\\max}$ \\\\ \\smallskip\\\\ \\multicolumn{5}{c}{Parameters for each planet} \\\\ Orbital period & $P_i$ & $\\frac{P_i^{-1}}{\\log(P_{\\max}/P_{\\min})}$ & 1 d & $10^3$ yr \\\\ Velocity semi-amplitude & $K_i$ & $\\frac{(K_i+K_0)^{-1}}{\\log\\left[1.+\\frac{K_{\\max}}{K_0}\\left(\\frac{P_{\\min}}{P_i}\\right)^{1/3}\\right]}$ & 0 m/s & $K_{\\max}$ $\\left(\\frac{P_{\\min}}{P_i}\\right)^{\\frac{1}{3}}$ \\\\ Orbital eccentricity & $e_i$ & 1  & 0 & 1 \\\\ Argument of periastron & $\\omega_i$ & $\\frac{1}{2\\pi}$ & 0 & $2\\pi$ \\\\ Orbital phase  & $M_i$ & $\\frac{1}{2\\pi}$ & 0 & $2\\pi$ \\\\ \\noalign{\\smallskip} \\tableline \\end{tabular} } \\end{center} \\label{TabPriors} \\end{table}    \\subsection{Previous Research}  Identifying the correct orbital solution from a set of RV (or other dynamical) observations is challenging due to the necessity of considering a very large parameter space of possible solutions (see Fig.\\ 1, right).  For RV planet searches, there are at least five model parameters per planet and one model parameter per observatory. To make the global search problem tractable, RV data sets are traditionally searched for sinusoidal signals (potential planets) using a periodogram.  The advantage of the periodogram is that it is extremely tractable computationally ($O(n\\log n)$) and can be quite useful for identifying potential orbital periods.  Then Levenberg-Marquardt (LM) minimization \\citep{NR} is applied starting from several initial guesses of orbital solutions near each of the potential orbital periods identified by the periodogram.  If the quality of the fit is consistent with the combination of the observational uncertainties \\citep{Butler96} and the expected intrinsic stellar variability \\citep{Wright05}, then estimates of the uncertainties in orbital parameters are obtained by repeatedly finding the best-fit orbital parameters (with LM) to several synthetic data sets generated via bootstrap resampling (with replacement) of the observational data \\citep{NR}.  These methods work quite well for analyzing stars with a single planet on a low-eccentricity, short-period (relative to the duration of the observations) orbit when the velocity perturbation is large (relative to the measurement uncertainties).  Since such planets are the easiest to discover, they are common among the known sample of planets, and the traditional frequentist methods have proven quite valuable in their discovery.  In recent years, RV searches have become increasingly sensitive to planets with small velocity amplitudes and/or long orbital periods, as well as planets in multiple planet systems.  Bayesian inference can help with each of these challenges and has the potential to significantly improve the sensitivity of detections and accuracy of orbital determinations.  Recent work has begun to develop the framework and computational tools to make this happen.  For example, \\citet{Cumming04} discussed the relationship between the periodogram method and a Bayesian analysis that assumes any planet is on a circular orbit.  \\citet{Ford06Adapt} combined brute force Monte Carlo (to integrate over orbital period) and the Laplace approximation (to integrate over the remaining model parameters) to render Bayesian model selection practical for planets assumed to be on a circular orbit (e.g., short-period planets prone to tidal circularization). MCMC has been applied to estimate orbital parameters and their uncertainty and to help understand the situations where traditional frequentist methods leave significant room for improvement \\citep{Ford05Mcmc,DriscolMTh}.  \\citet{Gregory05Book} and \\citet{Ford06Fast} have automated the determination of parameters of the candidate transition PDFs.  \\citet{Ford06Fast} has also identified non-linear candidate transition PDFs that dramatically accelerate the convergence of MCMC.  These advances make it computationally feasible for MCMC to characterize the allowed orbital solutions for single planets or weakly-interacting multiple planet systems \\citep[e.g.][]{Ford05UpsAnd}.  \\citet{Gregory05AAT} has applied parallel tempering to allow MCMC to explore multiple orbital solutions widely separated in parameter space. \\citep{Loredo04} developed the theoretical framework for applying Bayesian adaptive experimental design to dynamical extrasolar planet searches, and \\citet{Ford06Adapt} developed computationally practical algorithms for applying these techniques to adaptively schedule radial velocity observations. Bayesian techniques are just beginning to be applied to analyze the population of extrasolar planets \\citep{Ford06Tidal,Loredo06}. ",
        "conclusions": "In this paper, we have reviewed several methods for calculating the marginal posterior probabilities in the context of RV planet searches. One the positive side, we found that several algorithms were able to accurately calculate the marginal posterior probability for a simple test case, where there was a single dominant peak in the posterior probability distribution.  However, all of the estimators based on sampling from either the prior or posterior had serious short comings.  The method of partial linearization can be a useful tool for rapidly computing relatively low accuracy estimates of $m(\\vec{v})$ for data sets with one planet or even $\\sim1-3$ planets on low eccentricity orbits.  However, it rapidly becomes computationally intractable when there are multiple planets with significant eccentricities.  Restricted Monte Carlo ($\\widehat{m}_{RMC}(\\vec{v})$) can be useful for planets with large eccentricities, but is computationally feasible only once the orbital parameters are relatively well constrained.  Parallel tempering is able to estimate $m(\\vec{v})$ even for multimodal posterior distributions, but for our test data set $\\widehat{m}_{PT}(\\vec{\\theta})$ converged more slowly than all of the other algorithms tested (except basic Monte Carlo).  For our test case, we found no regime where the harmonic mean ($\\widehat{m}_{NR}(\\vec{v})$) or the weighted harmonic mean ($\\widehat{m}_{WHM,h}(\\vec{v})$) would be the most desirable estimator.  The new ratio estimator ($\\widehat{m}_{RE}(\\vec{v})$) performed very well for our test case, but we recommend proceeding with caution, based on preliminary tests with more complex data sets.  Based on our tests, the most promising methods are based on importance sampling using an analytic density that mimics the posterior (e.g., $\\widehat{m}_{IS,N}(\\vec{v})$ or $\\widehat{m}_{DIS,g^*,\\vec{\\beta}^*}(\\vec{v})$).  When the posterior has a single dominant peak that can be reasonably approximated by a multivariate normal distribution, then simple importance sampling ($\\widehat{m}_{IS,N}(\\vec{v})$) provides a very efficient tool for estimating marginal posteriors.  In cases where the posterior is more complex (e.g., multiple peaks and/or non-linear parameter correlations), then importance sampling can still be useful when combined with a mixture distribution based on a sample from the posterior that can be readily calculated via standard MCMC. Refinements to the basic importance sampling algorithm (e.g., $\\widehat{m}_{DIS,g^*,\\vec{\\beta}^*}(\\vec{v})$) can provide increased robustness and offer a tool for diagnosing when the mixture distribution is sufficient.  We hope that future research will improve our understanding of these estimator's theoretical and real-life properties, as well as lead to additional refinements.  In particular, we hope to investigate how the estimator $\\widehat{m}_{DIS,g^*,\\vec{\\beta}^*}(\\vec{v})$ performs on more widely dispersed posterior distributions and on higher dimensional problems (e.g., multiple planet systems).  In a sense, we can consider the problem of Bayesian model selection to have been reduced to the problem of constructing an analytic approximation to a probability density based only on a set of samples from the distribution.  Unfortunately, we recognize our algorithm for constructing importance sampling densities is not yet sufficiently robust to be applied generally.  Therefore, we would like to see additional research that would improve the robustness and computational efficiency of these algorithms.  Fortunately, we recognize several ways our current algorithm could be improved.  For example, rather than centering the mixture components on random samples from the posterior, it might be possible to make do with a smaller number of mixture components.  Perhaps methods making use of Voronoi tessellations and/or quasi-Monte Carlo methods could be beneficial in constructing good mixture distributions with fewer components, improving the computational efficiency of these methods.  Finally, we note that this field is still young, and additional research is needed to explore a wide range of methods for estimating $m(\\vec{\\theta})$, including methods on importance sampling, parallel tempering, reversible jump MCMC, MC$^3$, and nested sampling."
    },
    "0608/astro-ph0608091_arXiv.txt": {
        "abstract": "We describe a model-independent integral constraint which defines an upper limit to the allowed degree of evolution in the ratio of black hole (BH) mass to host galaxy luminosity or mass, as a function of redshift.  Essentially, if the BH/host ratio is excessive at redshift $z$, then it would imply that the total mass density in BHs above some $\\mbh^{\\rm min}$ is larger at that redshift than at $z=0$, which is impossible.  This argument requires no knowledge of host or BH properties, only a lower limit to the observed luminosity density in the brightest galaxies at some $z$.  We calculate this upper limit from a variety of luminosity and mass functions in different bands from redshifts $z=0-2$.  We show that it is consistent with passive evolution of spheroid populations (with a fixed $M_{\\rm BH}/M_{\\rm host}$ relation) in all cases, and provides tighter constraints than have generally been obtained previously, ruling out at $\\gtrsim6\\sigma$ observational and theoretical estimates suggesting that $M_{\\rm BH}/M_{\\rm host}$ was significantly larger at high redshifts than locally, although relatively weak (factor $\\lesssim2$ by $z=2$) evolution is still allowed. We discuss a variety of possible ``loopholes'' or changes in the BH/host populations and correlations, and show that they typically lower the upper limits and strengthen our conclusions. ",
        "introduction": "\\label{sec:intro}  Recent discoveries of tight correlations between the masses of supermassive black holes (BHs) in the centers of nearby galaxies and either the luminosity \\citep[e.g.,][]{KormendyRichstone95}, mass \\citep{Magorrian98}, or velocity dispersion \\citep{FM00,Gebhardt00} of their host spheroids demonstrate a fundamental link between the growth of BHs and galaxy formation.  Determining the evolution of these correlations with redshift is critical for informing analytical models \\citep[e.g.,][]{SR98} and simulations \\citep{DSH05,Robertson05} which follow the co-formation of BHs and bulges, as well as theories which relate the evolution and statistics of BH formation and quasar activity to galaxy mergers \\citep[e.g.,][]{H06a,H06e} and to the remnant spheroid population \\citep{H06b}.  Likewise, the significance of observations tracing the buildup of spheroid populations \\citep[e.g.,][]{Cowie96,Wolf03,Bell04,Bundy05b} and associations between spheroids in formation, mergers, and quasar hosts \\citep[e.g.,][]{H06d} depends on understanding the evolution of BH/host correlations.  Efforts to directly infer these correlations at redshifts $z>0$ are, however, difficult and limited by the small numbers of objects with observable hosts.  Consequently, different groups have reached seemingly contradictory conclusions, with O[{\\small III}] velocity dispersion \\citep{Shields03} and BH clustering \\citep{AdelbergerSteidel05} measurements finding no evolution in these correlations out to redshifts $z\\sim4$, while CO velocity dispersion \\citep{Walter04,Shields05} and host $R$-band luminosity \\citep{Peng06} observations imply significant evolution (with BHs becoming over-massive relative to their hosts) occurring at $z\\gtrsim1$. Spectral template fitting \\citep{Treu04,Woo06} suggests strong evolution at much lower redshifts $z=0.37$, and comparison of radio quasar and galaxy populations \\citep{McLure06} indicates evolution of the form $\\propto (1+z)^{2}$. In addition to the difficulty of identifying host properties at these redshifts, the measurements generally rely upon virial line-width relations \\citep[e.g.,][and references therein]{Kaspi05} to determine BH masses, which have not been tested at high redshifts and require that the BH be visible as an optical, broad-line AGN.  It is therefore of interest to develop additional constraints on the evolution of BH/host correlations, especially ones which are independent of systematics in host measurements and selection effects in considering only AGN populations. For example, \\citet{Merloni04} compare the integrated quasar luminosity and stellar mass densities, and infer a much weaker evolution of the form $\\propto(1+z)^{0.5}$, but this is subject to systematic uncertainties in quasar radiative efficiencies and (as discussed below) in observational completeness and galaxy type segregation at low luminosities.  In this paper, we formulate and apply a model-independent integral constraint which provides an upper limit to evolution in these correlations as a function of redshift, and show that it rules out various proposed forms of evolution at high significance.  We adopt a $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$, $H_{0}=70\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$ cosmology, but the choice of $H_{0}$ factors out in our analysis and the limits depend only weakly on $\\Omega_{\\rm M}$ and $\\Omega_{\\Lambda}$. ",
        "conclusions": "\\label{sec:discuss}  Our model-independent integral measure provides a constraint on the evolution in the correlation between BH mass and host luminosity or mass.  The upper limits we derive are competitive with or stronger than those from direct, but difficult observations of high-redshift BHs and their hosts, and can be applied in any band (including those for which direct estimates of these correlations at $z>0$ do not exist).  Our upper limits are consistent with passive evolution (with reasonable formation redshifts $z_{f}\\lesssim6$) at all redshifts and in every band, as well as hydrodynamic simulations of BH and spheroid co-formation at different redshifts \\citep{Robertson05}, and various estimates of these correlations at low redshifts \\citep[$z\\lesssim1$,][]{Shields03,Peng06}. The weak evolution inferred by \\citet{Merloni04} cannot be ruled out for sufficiently low $z_{f}\\lesssim3$, but the mass function data prefer a no-evolution case at $\\sim2\\,\\sigma$. However, the best-fit evolution estimates (open points in Figure~\\ref{fig:limits}) by e.g.\\ \\citet{Treu04,Shields05,Peng06,McLure06,Woo06} are ruled out at $\\gtrsim6\\sigma$ with respect to the upper limits we measure over each corresponding redshift interval. (Note that this does not necessarily mean their measurements are inconsistent with our upper limits at $6\\sigma$, as the error bars in the direct observations are large and generally inconsistent with passive evolution at only $\\sim2-3\\sigma$.) It should be noted, however, that there is a factor $\\sim2$ systematic normalization uncertainty in the virial BH mass estimators these authors adopt, and lowering e.g.\\ the estimates of \\citet{Peng06} by such a factor (i.e.\\ allowing for a factor of $\\sim2$ increase in $\\mbh/M_{\\ast}$ by $z\\sim2$) makes them consistent with our upper limits.  Given this apparent contradiction, it is natural to ask whether there are effects which could increase the upper limits ($\\Gamma_{\\rm max}$) in Figure~\\ref{fig:limits}.  Observational incompleteness will only lead to a higher ``true'' luminosity density at some $z$, and lower $\\Gamma_{\\rm max}$ (assuming the luminosity density at $z=0$ is not substantially less complete than that at higher redshift).  If we have somehow severely overestimated the mass in large spheroids at high $z$, i.e.\\ if these red systems are really all dusty starbursts, this could increase $\\Gamma_{\\rm max}$, but many of the samples shown in Figure~\\ref{fig:limits} are morphologically, not color, selected, and observations find that only a small fraction $\\lesssim10\\%$ of the luminosity density in red galaxies at $z\\sim1$ comes from dusty, intrinsically blue sources (too small to change our conclusions) \\citep{Bell04b}. However, this may rise to $\\sim30-50\\%$ at $1\\lesssim z\\lesssim 2$ \\citep[e.g.,][]{Cimatti02}, shifting the limits at these redshifts from the color-selected samples of \\citet{Fontana04} and \\citet{Giallongo05} up by a factor $\\sim1.5-2$.  Furthermore, the most meaningful upper limits come from high luminosities (typically $\\sim0.5-1$\\,mag above the LF break at each $z$), where a significant fraction of the most massive galaxies are in place by $z\\sim1$ and this constraint is most sensitive to a change in $\\mu(z)$, but also where the observations are most well-constrained. For this reason, it makes no difference which local LF we use to compare in Equation~(\\ref{eqn:limit}), as they all agree well at these luminosities (see e.g.\\ Figures~8 \\& 9 of \\citet{Bell03}).  We have also directly tested this with the local LFs in \\citet{Kochanek01,Bell03,Blanton03,Nakamura03} and find the results are similar.  Moreover, it makes no difference whether or not we use the cumulative (all galaxy) LFs (i.e.\\ assume a correlation between $\\mbh$ and total galaxy light or mass) instead of spheroid LFs as above, since spheroids dominate the luminosity density above these bright luminosities. In any case, the contribution of late-types tends to evolve in a similar fashion to that of early types \\citep[see e.g.\\ the observations in Figure~\\ref{fig:limits} and][]{Drory04}.  In even more detail, allowing for late types but accounting for their bulges (with more detailed morphology-dependent LFs) makes negligible difference, as e.g.\\ \\citet{Koo05} find most massive bulges are already in place at $z\\sim1$, giving a similar constraint but with more objects and thus a stronger $1\\sigma$ upper limit.  Changing the form of the correlation also does not alter our results. Our derivation of Eqn.~\\ref{eqn:limit} can be performed for any correlation of the form $\\mbh=\\mu(z)\\,f(\\lhost)$ or $\\mbh=f[\\mu(z)\\,\\lhost]$, and any relationship where $\\mbh$ is an increasing function of host mass or luminosity will give a similar result.  Furthermore, all non-linear forms of the correlation $f(\\lhost)$ which have been proposed observationally or theoretically steepen at high luminosity (high $\\mbh$), increasing the implied mass of the most massive BHs and their relative weight in the BH mass density, and directly considering these non-linear \\citep{KormendyRichstone95,MarconiHunt03,HaringRix04} or even log-quadratic \\citep{Wyithe06} correlations in every case {\\em lowers} $\\Gamma_{\\rm max}$.  But what if the shape or slope of $f(\\lhost)$ changes with redshift? This would be particularly interesting, in the implication that the ratio $\\mbh/M_{\\ast}$ evolves differently for systems of different mass. The exact effect on our derivation depends on the ``pivot point'' about which the relationship slope changes. For example, given that most of the $\\mbh-\\lhost$ measurements in Figure~\\ref{fig:limits} have been restricted to high-$\\mbh$, keeping these points fixed but finding a steeper slope at high redshift would imply a lower observed $\\Gamma(z)$ at lower masses, which could relieve tension between our upper limits and the observations, although it would also {\\em raise} the implied $\\Gamma(z)$ at higher masses. In such a case, our integral constraint still obtains, it simply must be applied as a function of $\\mbh^{\\rm min}$ or $\\ldummy$, as $\\Gamma(z)$ becomes a function of mass as well as redshift. The limits shown in Figure~\\ref{fig:limits} would then refer to the specific $\\mbh^{\\rm min}$ to which we perform our integration (generally $\\sim10^{8}-10^{9}\\,M_{\\sun}$).   Allowing for scatter in the BH/host relationship $f(\\lhost)$ likewise does not alter our results.  Non-zero scatter $\\sigma(z)$ will introduce a term $\\sim\\exp{\\{0.5\\,\\ln^{2}{10}\\, (\\sigma^{2}(z)-\\sigma^{2}(0)) \\}}$ in the left-hand side of Equation~(\\ref{eqn:limit}) (the appropriate exact factor will be a non-trivial function of $\\ldummy$, but since the relevant constraints come from somewhat above the LF break, this factor is accurate to $\\sim10\\%$ and the qualitative behavior is identical).  A larger scatter implies more objects at the highest $\\mbh$ for a given host LF, so if the scatter in these relations were significantly smaller at $z$, it would increase $\\Gamma_{\\rm max}$. However, every observational indication, especially from the observations in conflict with a low $\\Gamma_{\\rm max}$, is for an increasing scatter with $z$ \\citep{Shields03,Shields05,Treu04,Peng06,Woo06}.  Moreover, if the mean $\\mbh$-host relation were higher in the past ($z\\sim1$), we still know observationally that there is a non-negligible population (especially at the highest masses of interest) which has only passively evolved since such times (no star formation beyond the $\\sim4\\%$ level, and $M_{\\rm BH}$ can only increase) \\citep[see e.g.,][]{Kauffmann03,Bell04,Faber05}, so these systems must have been at or below the present relation at that redshift (spheroid-spheroid ``dry'' mergers, which by definition conserve the mean $\\mbh$-host mass/luminosity relations, cannot change this).  Therefore, the scatter must increase with $z$ in such a scenario, which gives a {\\em lower} $\\Gamma_{\\rm max}$.  Since the limits on $\\Gamma(z)$ rule out $\\mbh$ increasing strongly in the mean (i.e.\\ for the general population) relative to the host bulge or total galaxy, it is next worth asking whether the discrepant $\\Gamma(z)$ from some of the observations in Figure~\\ref{fig:limits} could represent some particular (short-lived) stage of development. The simplest possibility is that once an object has formed its bulge or spheroid, it will lie on the appropriate $z=0$ (passively evolved) relation, but that the observations see disks which have grown their BHs but have not yet grown a bulge. In such a scenario, $\\Gamma(z)$ in spheroids is given by the passive evolution tracks, and there is an independent, potentially much larger $\\Gamma(z)$ for ``pre-spheroids.''  This pre-spheroid population, however, cannot constitute all disks or blue ``star-forming'' galaxies.  First, at least some disks have survived since $z\\sim1$, and this would demand they have overmassive BHs relative to the total galaxy light/mass, where in fact the opposite is observed \\citep[i.e.\\ $\\mbh$ correlates with bulge, not total galaxy mass; e.g.,][]{KormendyRichstone95}.  Second, this scenario trivially modifies our calculation of $\\Gamma_{\\rm max}$ -- we integrate over spheroids with $\\Gamma(z)$ from the passive evolution tracks and determine the upper limit to $\\Gamma(z)$ in the remaining disk population. In Figure~\\ref{fig:limits} it is clear that the previous $\\Gamma_{\\rm max}$ is already marginally consistent with the passive evolution tracks; adding an additional source of high-mass BHs is not possible. Regardless, integrating down to a lower $\\ldummy$ (where disks begin to dominate the population) gives a similar constraint to those in Figure~\\ref{fig:limits}, since there is little room for ``additional'' BH mass density and, as noted above, the disk population evolves similarly to spheroids at the bright end (in the sense that the evolution is weak).  The remaining possibility is that objects with BHs larger than our $\\Gamma_{\\rm max}$ upper limits are an even more limited sub-class of galaxies. The natural expectation might be that AGN or bright quasars in particular obey this discrepant relation, since essentially all of the estimates of the high-redshift relation are from quasars. Locally at least, bright quasars constitute a tiny fraction of galaxies, so they can have much higher BH masses than their inactive counterparts (i.e.\\ higher $\\mu_{\\rm QSO}/\\mu(z=0)$) without contributing much total BH mass. However, this phase would have to be short-lived, much like the quasar phase itself, or else it is equivalent to saying that a large fraction of non-spheroids obey this higher $\\mu(z)$ and violates the integral constraints as above. Therefore, the host must ``catch up'' to have the appropriate stellar mass for the $z=0$ (passively evolved) relation in a short period of time. However, cosmological cooling and infall generally requires a non-negligible fraction of a Hubble time, implying (if this were the mechanism to add bulge mass) that too large a fraction of galaxies would lie on the discrepant relation (given the integral constraints above), and invoking subsequent mergers or interactions to transport gas would imply that the quasar somehow ``knows'' the system is going to interact before the interaction begins.  This suggests that the mass required for the bulge to ``catch up'' is present, in the form of gas ``about'' to be turned into stars \\citep[see also e.g.][]{Croton05}, but this hypothesis encounters a number of observational conflicts. To reach the appropriate $z=0$ relation given the observations in Figure~\\ref{fig:limits} which imply strong evolution in $\\mbh/M_{\\ast}$, the systems would have to be $\\gtrsim75\\%$ gas at $z\\gtrsim1$, which for the appropriate total mass range at these redshifts with these concentrations gives quiescent (lower-limit) star formation rates $\\sim10-100$'s of $M_{\\sun}\\,{\\rm yr^{-1}}$ \\citep[e.g.,][]{Kennicutt98,SDH05b}. The typical quasar hosts would be late-type galaxies, and in the subsequent, brief period of rapid star formation the BH would be overmassive with the galaxy ``catching up'' (and could not significantly accrete without preventing this ``catching up'').  All of these predictions appear to be opposite the observations, with quasar spectra dominated by {\\em post}-starburst activity \\citep{VandenBerk06}, typically early-type, bulge-dominated hosts \\citep[especially in the brightest quasars;][]{Dunlop03,Floyd04}, and observed starbursts usually showing high accretion rate BHs which are at that moment undermassive relative to their hosts \\citep[e.g.,][]{Borys05}.  Moreover, if star formation is ongoing in these objects to grow their bulges, it implies a lower mass-to-light ratio and suggests that the discrepancy between their $\\mbh$-host relations is even larger than that in Figure~\\ref{fig:limits}, amplifying these difficulties.  Furthermore, by the highest luminosities and redshifts $z\\sim1$, the active fraction of BHs with masses $M_{\\rm BH}\\gtrsim10^{9}\\,M_{\\sun}$ is $\\gtrsim10\\%$ \\citep[relative to the total number of such objects at $z\\sim0$;][]{MD04}, which creates some tension if these are a different population from the ``finished'' large spheroids at that redshift, which can already account for the $z=0$ mass density of such large BHs (but is perfectly consistent if these are the same population, and quasars have large bulges).  Finally, associating this discrepant $\\Gamma(z)$ with quasars connects it not with the cosmological buildup of galaxies but with a particular evolutionary ``snapshot'' or event. This makes it difficult to explain why quasars at $z\\gtrsim1$ would appear in bulge-less systems as a ``signpost'' of bulge formation about to occur, while quasars of the same luminosity at $z\\lesssim1$ and low-luminosity AGN (which form a continuum in $L_{\\rm QSO}$) would appear in normal, massive, bulge-dominated systems."
    },
    "0608/astro-ph0608058_arXiv.txt": {
        "abstract": "We investigate the self-consistent electrodynamic structure of a particle accelerator in the Crab pulsar magnetosphere on the two-dimensional poloidal plane, solving the Poisson equation for the electrostatic potential together with the Boltzmann equations for electrons, positrons and gamma-rays. If the trans-field thickness of the gap is thin, the created current density becomes sub-Goldreich-Julian, giving the traditional outer-gap solution but with negligible gamma-ray luminosity. As the thickness increases, the created current increases to become super-Goldreich-Julian, giving a new gap solution with substantially screened acceleration electric field in the inner part. In this case, the gap extends towards the neutron star with a small-amplitude positive acceleration field, extracting ions from the stellar surface as a space-charge-limited flow. The acceleration field is highly unscreened in the outer magnetosphere, resulting in a gamma-ray spectral shape which is consistent with the observations. ",
        "introduction": "\\label{sec:intro} The Energetic Gamma Ray Experiment Telescope (EGRET) aboard the Compton Gamma Ray Observatory has detected pulsed signals from at least six rotation-powered pulsars (e.g., for the Crab pulsar, Nolan et al. 1993, Fierro et al. 1998). Since interpreting $\\gamma$-rays should be less ambiguous compared with reprocessed, non-thermal X-rays, the $\\gamma$-ray pulsations observed from these objects are particularly important as a direct signature of basic non-thermal processes in pulsar magnetospheres, and potentially should help to discriminate among different emission models.  The pulsar magnetosphere can be divided into two zones (fig.~\\ref{fig:sidev}): The closed zone filled with a dense plasma corotating with the star, and the open zone in which plasma flows along the open field lines to escape through the light cylinder. The last-open field lines form the border of the open magnetic field line bundle. In all the pulsar emission models, particle acceleration and the resultant photon emissions take place within this open zone.  \\begin{figure} \\epsscale{0.6} \\plotone{f01.eps} \\caption{ Schematic figure (sideview) of the two representative accelerator models. The small filled circle represents the neutron star \\label{fig:sidev} } \\end{figure}  On the spinning neutron star surface, an electro-motive force, $\\mbox{EMF}\\approx \\Omega^2 B_\\ast r_\\ast^3/c^2 \\approx 10^{16.5}$~V, is exerted from the magnetic pole to the rim of the polar cap. In this paper, we assume that both the spin and magnetic axes reside in the same hemisphere; that is, $\\mbox{\\boldmath$\\Omega$}\\cdot\\mbox{\\boldmath$\\mu$}>0$, where $\\mbox{\\boldmath$\\Omega$}$ represents the rotation vector, and $\\mbox{\\boldmath$\\mu$}$ the stellar magnetic moment vector. This strong EMF causes the magnetospheric currents that flow outwards in the lower latitudes and inwards near the magnetic axis (left panel in fig.~\\ref{fig:current}). The return current is formed at large-distances where Poynting flux is converted into kinetic energy of particles or dissipated (Shibata~1997).  \\begin{figure} \\epsscale{0.8} \\plottwo{f02a.eps}{f02b.eps} \\caption{ Schematic picture of electric current in the pulsar magnetosphere. {\\it Left}: Global electric current due to the EMF exerted on the spinning neutron star surface when $\\mbox{\\boldmath$\\Omega$}\\cdot\\mbox{\\boldmath$\\mu$}>0$. {\\it Right}: Current (downward arrows) derived in the inner-slot-gap (shaded region). \\label{fig:current} } \\end{figure}  Attempts to model the particle accelerator % have traditionally concentrated on two scenarios: Polar-cap models with emission altitudes of $\\sim 10^4$cm to several neutron star radii over a pulsar polar cap surface (Harding, Tademaru, \\& Esposito 1978; Daugherty \\& Harding 1982, 1996; Dermer \\& Sturner 1994; Sturner, Dermer, \\& Michel 1995), and outer-gap models with acceleration occurring in the open zone located near the light cylinder (Cheng, Ho, \\& Ruderman 1986a,b, hereafter CHR86a,b; Chiang \\& Romani 1992, 1994; Romani and Yadigaroglu 1995). Both models predict that electrons and positrons are accelerated in a charge depletion region, a potential gap, by the electric field along the magnetic field lines to radiate high-energy $\\gamma$-rays via the curvature and inverse-Compton (IC) processes.  In the outer magnetosphere picture of Romani (1996), he estimated the evolution of high-energy flux efficiencies and beaming fractions to discuss the detection statistics, by considering how pair creation on thermal surface flux can limit the acceleration zones. Subsequently, Cheng, Ruderman and Zhang (2000, hereafter CRZ00) developed a three-dimensional outer magnetospheric gap model, self-consistently limiting the gap size by pair creation from collisions of thermal photons from the polar cap that is heated by the bombardment of gap-accelerated charged particles. The outer gap models of these two groups have been successful in explaining the observed light curves, particularly in reproducing the wide separation of the two peaks commonly observed from $\\gamma$-ray pulsars (Kanbach 1999; Thompson 2001), without invoking a very small inclination angle. In these outer gap models, they consider that the gap extends from the null surface, where the Goldreich-Julian charge density vanishes, to the light cylinder, beyond which the velocity of a co-rotating plasma would exceed the velocity of light, adopting the vacuum solution of the Poisson equation for the electrostatic potential (CHR86a).  However, it was analytically demonstrated by Hirotani, Harding, and Shibata (2003, HHS03) that an active gap, which must be non-vacuum, possesses a qualitatively different properties from the vacuum solution discussed in traditional outer-gap models. For example, the gap inner boundary shifts towards the star as the created current increases and at last touch the star if the created current exceeds the Goldreich-Julian (GJ) value at the surface. Therefore, to understand the particle accelerator, which extends from the stellar surface to the outer magnetosphere, we have to merge the outer-gap and polar-cap models, which have been separately considered so far.  In traditional polar-cap models, the energetics and pair cascade spectrum have had success in reproducing the observations. However, the predicted beam size of radiation emitted near the stellar surface is too small to produce the wide pulse profiles that are observed from high-energy pulsars. Seeking the possibility of a wide hollow cone of high-energy radiation due to the flaring of field lines, Arons (1983) first examined the particle acceleration at the high altitudes along the last open field line. This type of accelerator, or the slot gap, forms because the pair formation front (PFF), which screens the accelerating electric field, $\\Ell$, in a width comparable to the neutron star radius, occurs at increasingly higher altitude as the magnetic colatitude approaches the edge of the open field region (Arons \\& Scharlemann 1979). Muslimov and Harding (2003, hereafter MH03) extended this argument by including two new features: acceleration due to space-time dragging, and the additional decrease of $\\Ell$ at the edge of the gap due to the narrowness of the slot gap. Moreover, Muslimov and Harding (2004a,b, hereafter MH04a,b) matched the high-altitude slot gap solution for $\\Ell$ to the solution obtained at lower altitudes (MH03), and found that the residual $\\Ell$ is small and constant, but still large enough at all altitudes to maintain the energies of electrons, which are extracted from the star, above 5~TeV.  It is noteworthy that the polar-slot gap model proposed by MH04a,b is an extension of the polar-cap model into the outer magnetosphere, assuming that the plasma flowing in the gap consists of only one sign of charges. This assumption is self-consistently satisfied in their model, because pair creation in the extended slot gap occurs at a reduced rate and the pair cascade due to inward-migrating particles does not take place. In the polar-slot gap model, the completely charge-separated, space-charge-limited flow (SCLF) leads to a negative $\\Ell$ for $\\mbox{\\boldmath$\\Omega$}\\cdot\\mbox{\\boldmath$\\mu$}>0$. However, we should notice here that the electric current induced by the negative $\\Ell$ (right panel of figure~\\ref{fig:current}) contradicts with the global current patterns (left panel of figure~\\ref{fig:current}), which is derived by the EMF exerted on the spinning neutron-star surface, if the gap is located near the last-open field line. (Note that the return current sheet is not assumed on the last-open field line in the slot gap model.)  On these grounds, we are motivated by the need to contrive an accelerator model that predicts a consistent current direction with the global requirement. To this aim, it is straightforward to extend recent outer-gap models, which predict opposite $\\Ell$ to polar-cap models, into the inner magnetosphere. Extending the one-dimensional analysis along the field lines in several outer-gap models (Hirotani and Shibata 1999a,~b,~c; HHS03), Takata, Shibata, and Hirotani (2004, hereafter TSH04) and Takata et al. (2006, hereafter TSHC06) solved the Poisson equation for the electrostatic potential on the two-dimensional poloidal plane, and revealed that the gap inner boundary is located inside of the null surface owing to the pair creation within the gap, assuming that the particle motion immediately saturates in the balance between electric and radiation-reaction forces.  In the present paper, we extend TSH04 and TSHC06 by solving the particle energy distribution explicitly, and by considering a super-GJ current solution with ion emission from the neutron star surface. In \\S~\\ref{sec:basic_eqs}, we formulate the basic equations and boundary conditions. We then apply it to the Crab pulsar in \\S~\\ref{sec:appl}, and compare the solution with MH04 in \\S~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary} To conclude, we investigated the self-consistent electrodynamic structure of a particle accelerator in the Crab pulsar magnetosphere on the two-dimensional surface that contains the magnetic field lines threading the stellar surface on the plane in which both the rotation and magnetic axes reside. We regard the trans-field thickness, $h_{\\rm m}$, of the gap as a free parameter, instead of trying to constrain it. For a small $h_{\\rm m}$, the created current density, $j_{\\rm e}$, becomes sub-Goldreich-Julian, giving the traditional outer-gap solution but with negligible $\\gamma$-ray flux. However, as $h_{\\rm m}$ increases, $j_{\\rm e}$ increases to become super-GJ, giving a new gap solution with substantially screened acceleration electric field, $\\Ell$, in the inner part. In this case, the gap extends towards the neutron star with a small-amplitude positive $\\Ell$, which extracts ions from the stellar surface. It is essential to examine the pitch-angle evolution of the created particles, because the inward-migrating particles emit $\\gamma$-rays, which governs the gap electrodynamics through pair creation, via synchro-curvature process rather than pure-curvature one. The resultant spectral shape of the outward-directed $\\gamma$-rays is consistent with the existing observations; however, their predicted fluxes appear insufficient. The pulsar wind at the light cylinder is magnetically dominated: Along the field lines threading the gap, magnetization parameter, $\\sigma_{\\rm L}$, is about $10^2$.   \\subsection{How to obtain sufficient gamma-ray flux} \\label{sec:SR} The obtained $\\gamma$-ray fluxes in the present work, are all below the observed values. Nevertheless, this problem may be solved if we extend the current analysis into a three-dimensional configuration space. As mentioned in \\S~\\ref{sec:Boltz_gamm}, we assume $k_\\varphi=0$ and neglect the aberration of photon emission directions. However, in a realistic three-dimensional pulsar magnetosphere, $\\gamma$-rays will have angular momenta and may be emitted in a limited solid angle as suggested by the caustic model (e.g., Dyks \\& Rudak~2003; Dyks, Harding, \\& Rudak 2004), which incorporate the effect of aberration of photons and that of time-of-flight delays. In particular, in the trailing peak of a highly inclined rotator, photons emitted at different altitudes $s$ will be beamed in a narrow solid angle to be piled up at the same phase of a pulse (Morini~1983; Romani \\& Yadigaroglu~1995), resulting in a $\\gamma$-ray flux which is an order of magnitude greater than the present values. Thus, the insufficient $\\gamma$-ray fluxes does not suggest the inapplicability of the present method. It is noteworthy that the meridional propagation angles of the emitted photons (e.g., different curves in figs.~\\ref{fig:C70s1}, \\ref{fig:C70s2}, and \\ref{fig:incl}) can be readily translated into the emissivity distribution in the gap as a function of $s$. Therefore, in this work, we do not sum up the $\\gamma$-ray fluxes emitted into different meridional angles taking account of the aberration of light.  \\subsection{Stability of the Gap} \\label{sec:stability} Let us discuss the electrodynamic stability of the gap, by considering whether an initial perturbation of some quantity tends to be canceled or not. In the present paper, we consider that the soft photon field is given and unchanged when gap quantities vary. Thus, let us first consider the case when the soft photon field is fixed. Imagine that the created pairs are decreased as an initial perturbation. It leads to an increase of the potential drop due to less efficient screening by the discharged pairs, and hence to an increase of particle energies. Then the particles emit synchro-curvature radiation efficiently, resulting in an increase of the created pairs, which tends to cancel the initial decrease of created pairs.  Let us next consider the case when the soft photon field also changes. Imagine again that the created pairs are decreased as an initial perturbation. It leads to an increase of particle energies in the same manner as discussed just above. The increased particle energies increase not only the number and density of synchro-curvature $\\gamma$-rays, but also the surface blackbody emission from heated polar caps and the secondary magnetospheric X-rays. Even though neither the heated polar-cap emission nor the magnetospheric emission are taken into account as the soft photon field illuminating the gap in this paper, they all work, in general, to increase the pair creation within the gap, which cancels the initial decrease of created pairs more strongly than the case of the fixed soft photon field.  Because of such negative feedback effects, solution exists in a wide parameter space. For example, the created current density is almost unchanged for a wide range of $h_{\\rm m}$ (e.g., compare the dash-dotted and dash-dot-dot-dotted curves in fig.~\\ref{fig:C70d}). On these grounds, although the perturbation equations are not solved under appropriate boundary conditions for the perturbed quantities, we conjecture that the particle accelerator is electrodynamically stable, irrespective whether the X-ray field illuminating the gap is thermal or non-thermal origin.  \\subsection{Local vs. global currents} \\label{sec:current} Let us briefly look at the relationship between the locally determined current density $j_{\\rm gap}$ (eq.~[\\ref{eq:jgap}]) and the globally required one, $j_{\\rm global}$. It is possible that $j_{\\rm global}$ is constrained independently from the gap electrodynamics by the dissipation at large distances (Shibata~1997), which provides the electric load in the current circuit, or by the condition that the magnetic flux function should be continuous across the light cylinder, as discussed in recent force-free argument of the trans-field equation (Contopoulos, Kazanas, Fendt 1999; Goodwin et al. 2004; Gruzinov 2005; Spitkovsky 2006). In either case, $j_{\\rm global}$ will be more or less close to unity (i.e., typical GJ value). On the other hand, as demonstrated in \\S~\\ref{sec:superGJ}, for super-GJ cases, $j_{\\rm gap}$ is automatically regulated around $0.9$ for a wide parameter range. Thus, the discrepancy between $j_{\\rm gap}$ and $j_{\\rm global}$ is small provided that $h_{\\rm m}$ is large enough to maintain the created current density at a super-GJ value. The small imbalance $j_{\\rm global}-j_{\\rm gap}$ may have to be compensated by a current injection across the outer boundary (if the gap terminates inside of the light cylinder, charged particles could be injected from the outer boundary), or by an additional ionic emission from the stellar surface (if the imbalance leads to an additional residual $\\Ell$ at the surface). In any case, the injected current is small compared with $j_{\\rm gap}$; thus, it will not change the electrodynamics significantly, even though the gap active region may be shifted to some degree along the magnetic field lines, as demonstrated by Hirotani and Shibata (2001a, b; 2002), TSH04, and TSHC06.  \\subsection{Created pairs in the inner magnetosphere} \\label{sec:pairs} Let us devote a little more space to examining the particle flux along the open field lines that do {\\it not} thread the gap (i.e., $h_{\\rm m}<h<1$). Since $\\Ell$ vanishes, the created, secondary pairs emit synchrotron photons, which are capable of cascading into tertiary and higher-generation pairs by $\\gamma$-$\\gamma$ or $\\gamma$-B collisions. Examining the cascade, we can calculate the rate of pair creation, which takes place mainly in the inner magnetosphere. Denoting that the pair creation rate is $\\kappa_{\\rm w}(h)\\cdot(\\Omega B/2\\pi e)$ per unit area per second, we find $\\kappa_{\\rm w}= 2.2 \\times 10^4$, $2.1 \\times 10^4$, $1.9 \\times 10^4$, $1.8 \\times 10^4$, $1.6 \\times 10^4$, $1.1 \\times 10^4$, $1.0 \\times 10^4$, $0.99\\times 10^4$, $0.95\\times 10^4$, $0.96\\times 10^4$, $1.0 \\times 10^4$ at $h=0.20125$, $0.2025$, $0.20375$, $0.205$, $0.305$, $0.405$, $0.505$, $0.605$, $0.705$, $0.805$, and $0.905$, respectively, for $\\inc=70^\\circ$, $\\mu= 6.0\\times 10^{30}\\mbox{ G cm}^3$, and $h_{\\rm m}=0.200$. Thus, the averaged creation rate becomes ${\\bar\\kappa}_{\\rm w}= 1.4 \\times 10^4(\\Omega B/2\\pi e)$ pairs per unit area per second, giving $\\dot{N}_{\\rm pair}= 3.8 \\times 10^{38} \\mbox{s}^{-1}$ as the pair creation rate in the entire magnetosphere. It should be noted that this $\\dot{N}_{\\rm pair}$ appears less than the constraints that arise from consideration of magnetic dissipation in the wind zone (Kirk \\& Lyubarsky~2001, who derived $10^{40}\\,\\mbox{s}^{-1}$), and of Crab Nebula's radio synchrotron emission (Arons 2004, who derived $10^{40.5}\\,\\mbox{s}^{-1}$).  Due to strong synchrotron radiation, these inwardly created particles quickly lose energy to fall onto the lowest Landau level, preserving longitudinal momentum $\\sim m_{\\rm e}c$ per particle. These non-relativistic particles possess momentum flux of $ 2\\dot{N}_{\\rm pair} m_{\\rm e}c/(\\pi R_{\\rm PC}^2) \\sim 1.0\\times 10^{12} \\mbox{dyn cm}^{-2}$, where $R_{\\rm PC}\\sim \\sqrt{r_\\ast^3/\\rlc}$ denotes the polar cap radius. On the other hand, surface X-ray field has the upward momentum flux of $6.0 \\times 10^8 (kT/100\\mbox{eV})^4 \\mbox{dyn cm}^{-2}$. Thus, the created pairs will not be pushed back by resonant scatterings. They simply fall onto the stellar surface with non-relativistic velocities. The luminosity of $e^-$-$e^+$ annihilation line is about $3\\times10^{26}\\mbox{\\,ergs s}^{-1}$, which is negligible (e.g., compared with the $\\gamma$-ray luminosity $\\sim 10^{34.5}\\mbox{\\,ergs s}^{-1}$). On these grounds, for the Crab-pulsar, we must conclude that the present work fails to explain the injection rate of the wind particles, in the same way as in other outer-gap models.  \\subsection{Comparison with Polar-slot gap model } \\label{sec:SG} It is worth comparing the present results with the polar-slot-gap model proposed by Muslimov and Harding (2003; MH04a; MH04b), who obtained a quite different solution (e.g., negative $\\Ell$ in the gap) solving essentially the same equations under analogous boundary conditions for the same pulsar as in the present work. The only difference is the transfield thickness of the gap. Estimating the transfield thickness to be $\\Delta l_{\\rm SG} \\sim h_{\\rm m} r_\\ast \\sqrt{r/\\rlc}$, which is a few hundred times thinner than the present work, they extended the solution (near the polar cap surface) that was obtained by MT92 into the higher altitudes (towards the light cylinder). Because of this very small $\\Delta l_{\\rm SG}$, emitted $\\gamma$-rays do not efficiently materialize within the gap; as a result, the created and returned positrons from the higher altitudes do not break down the original assumption of the completely-charge-separated SCLF near the stellar surface.  To avoid the reversal of $\\Ell$ in the gap (from negative near the star to positive in the outer magnetosphere), or equivalently, to avoid the reversal of the sign of the effective charge density, $\\rho_{\\rm eff} = \\rho-\\rhoGJ$, along the field line, MH04a and MH04b assumed that $\\rho_{\\rm eff}/B$ nearly vanishes and remains constant above a certain altitude, $s=s_{\\rm c}$, where $s_{\\rm c}$ is estimated to be within a few neutron star radii. Because of this assumption, $\\Ell$ is suppress at a very small value and the pair creation becomes negligible in the entire gap. In another word, the enhanced screening is caused not only by the proximity of two conducting boundaries, but also by the assumption of $\\partial(\\rho_{\\rm eff}/B)/\\partial s=0$ within the gap (see eq.~[\\ref{eq:CHR}]). To justify this $\\rho/B$ distribution, MH04a and MH04b proposed an idea that $\\rho$ should grow by the {\\it cross field motion} of charges due to the toroidal forces, and that $\\rho_{\\rm eff}/B$ is a small constant so that $c\\rho_{\\rm eff}/B$ may not exceed the flux of the emitted charges from the star, which ensures the {\\it equipotentiality} of the slot-gap boundaries (see \\S~2.2 of MH04a for details).  The cross-field motion becomes important if particles gain angular momenta as they migrate outwards to pick up energies which is a non-negligible fraction of the difference of the cross-field potential between the two conducting boundaries. Denoting the fraction as $\\epsilon$, we obtain $\\Gamma m_{\\rm e}c^2 \\dot\\varphi \\Omega (\\varpi/c)^2 =\\epsilon\\,eB \\Delta l_{\\rm SG}$ (Mestel~1985; eq.~[12] of MH04a). If we substitute their estimate $\\Delta l_{\\rm SG} \\sim r_\\ast/20$, we obtain $\\epsilon \\sim 0.33 (\\dot\\varphi/\\Omega)\\gamma_7 B_{6}^{-1} r_6^{-3} (\\varpi/\\rlc)^2$, where $\\gamma_7=\\Gamma/10^7$, $B_{6}=B_\\ast/10^{6}\\,\\mbox{G}$, and $r_6=r_\\ast/10\\,\\mbox{km}$; therefore, the cross-field motion becomes important in the outer magnetosphere within their slot-gap model. A larger value of $\\Delta l_{\\rm SG}$ is incompatible with the constancy of $\\rho_{\\rm eff}/B$ due to the cross-field motion in the higher altitudes.  As for the equipotentiality of the boundaries, it seems reasonable to suppose that $c\\vert\\rho_{\\rm eff}\\vert/B < c\\vert\\rho_\\ast\\vert/B_\\ast$ should be held at any altitudes in the gap, as MH04a suggested, where $\\rho_\\ast$ denotes the real charge density in the vicinity of the stellar surface. However, the assumption that $\\rho_{\\rm eff}/B$ is a small positive constant may be too strong, because it is only a sufficient condition of $c\\vert\\rho_{\\rm eff}\\vert/B < c\\vert\\rho_\\ast\\vert/B_\\ast$.  In the present paper, on the other hand, we assume that the magnetic fluxes threading the gap is unchanged, considering that charges freely move along the field lines on the upper (and lower) boundaries. As a result, the gap becomes much thicker than MH04a,b; namely, $\\Delta l_{\\rm SG} \\sim 0.5 h_{\\rm m} \\rlc$, which gives $\\epsilon < 10^{-3}$. Therefore, we can neglect the cross-field motion and justify the constancy of $\\rho/B$ in the outer region of the gap, where pair creation is negligible. In the inner magnetosphere, $\\rho_{\\rm eff}/B$ becomes approximately a negative constant, owing to the discharge of the copiously created pairs. Because of this negativity of $\\rho_{\\rm eff}/B$, a positive $\\Ell$ is exerted. For a super-GJ solution, we obtain $j_{\\rm e}+j_{\\rm ion} \\sim 0.9 > \\rho_{\\rm eff}/(\\Omega B/2\\pi)$, which guarantees the equipotentiality of the boundaries. For a sub-GJ solution, a problem may occur regarding the equipotentiality; nevertheless, we are not interested in this kind of solutions.  In short, whether the gap solution becomes MH04a way (with a negative $\\Ell$ as an outward extension of the polar-cap model) or this-work way (with a positive $\\Ell$ as an inward extension of the outer-gap model) entirely depends on the transfield thickness and on the resultant $\\rho_{\\rm eff}/B$ variation along the field lines. If $\\Delta l_{\\rm SG} \\sim r_\\ast/10$ holds in the outer magnetosphere, $\\rho_{\\rm eff}/B$ could be a small positive constant by the cross-field motion of charges (without pair creation); in this case, the current is slightly sub-GJ with electron emission from the neutron star surface, as MH04a,b suggested. On the other hand, if $\\Delta l_{\\rm SG} > \\rlc/40$ holds in the outer magnetosphere, $\\rho_{\\rm eff}/B$ takes a small negative value by the discharge of the created pairs (see fig.~\\ref{fig:C70nine}); in this case, the current is super-GJ with ion emission from the surface, as demonstrated in the present paper. Since no studies have ever successfully constrained the gap transfield thickness, there is room for further investigation on this issue."
    },
    "0608/astro-ph0608003_arXiv.txt": {
        "abstract": "Infrared (IR) emission spectra are calculated for dust heated by starlight, for mixtures of amorphous silicate and graphitic grains, including varying amounts of polycyclic aromatic hydrocarbon (PAH) particles. The models are constrained to reproduce the average Milky Way extinction curve.  The calculations include the effects of single-photon heating. Updated IR absorption properties for the PAHs are presented, that are consistent with observed emission spectra, including those newly obtained by {\\it Spitzer Space Telescope}. We find a size distribution for the PAHs giving emission band ratios consistent with the observed spectra of the Milky Way and other galaxies.  Emission spectra are presented for a wide range of starlight intensities. We calculate how the efficiency of emission into different IR bands depends on PAH size; the strong 7.7$\\mum$ emission feature is produced mainly by PAH particles containing $<10^3$ C atoms. We also calculate how the emission spectrum depends on $U$, the starlight intensity relative to the local interstellar radiation field. The submm and far-infrared emission is compared to the observed emission from the local interstellar medium.  Using a simple distribution function, we calculate the emission spectrum for dust heated by a distribution of starlight intensities, such as occurs within galaxies. The models are parameterized by the PAH mass fraction $\\qpah$, the lower cutoff $\\Umin$ of the starlight intensity distribution, and the fraction $\\gamma$ of the dust heated by starlight with $U>\\Umin$. We present graphical procedures using Spitzer IRAC and MIPS photometry to estimate the parameters $\\qpah$, $\\Umin$, and $\\gamma$, the fraction $f_{\\rm PDR}$ of the dust luminosity coming from photodissociation regions with $U>100$, and the total dust mass $\\Mdust$. ",
        "introduction": "}  Interstellar dust in galaxies absorbs energy from starlight; this absorbed energy is then reradiated at infrared (IR) and far-IR wavelengths. The unprecedented sensitivity of Spitzer Space Telescope \\citep{Werner+Roellig+Low_etal_2004} allows this IR emission to be measured for a wide range of galaxy sizes and morphologies. The spectral properties of the IR emission from dust allow one to infer the composition of the dust, the size distribution of the dust particles, the intensity of the starlight that is heating the dust, and the total mass of dust.  Deducing the dust properties and starlight intensity is by no means direct or straightforward. To start, one needs to assume a provisional dust model with the physical nature of the dust (composition, geometry, and size distribution) fully specified. One then tries to constrain the grain properties by comparing model predictions with observations (e.g. interstellar extinction, scattering, polarization, IR and microwave emission, and interstellar depletions).  Various models have been proposed for interstellar dust. The models fall into three broad categories: the silicate-graphite model \\citep{Mathis+Rumpl+Nordsieck_1977, Draine+Lee_1984, Kim+Martin+Hendry_1994} and its natural extension -- the silicate-graphite-PAH model \\citep{Siebenmorgen+Krugel_1992, Li+Draine_2001b, Weingartner+Draine_2001a}; the silicate core-carbonaceous mantle model \\citep{Desert+Boulanger+Puget_1990, Jones+Duley+Williams_1990, Li+Greenberg_1997}, and the composite model which assumes the dust to be low-density aggregates of small silicates and carbonaceous particles \\citep{Mathis+Whiffen_1989, Mathis_1996, Zubko+Dwek+Arendt_2004}. The core-mantle model is challenged by the nondetection of the 3.4$\\mum$ C--H aliphatic hydrocarbon polarization on sightlines where the 9.7$\\mum$ Si--O silicate band is observed to be polarized \\citep{Chiar+Adamson+Whittet_etal_2006, Mason+Wright+Adamson+Pendleton_2006} [see also \\citet{Adamson+Whittet+Chrysostomou_etal_1999, Li+Greenberg_2002}]. The original composite model proposed by \\citet{Mathis+Whiffen_1989}, with $\\sim$\\,80\\% of the grain volume consisting of vacuum, may have too flat a far-IR emissivity to explain the observational data \\citep{Draine_1994}. The updated version of the composite model of \\citet{Mathis_1996} assumes a vacuum fraction $\\sim$\\,45\\% in order to make economical use of the available carbon to account for the observed extinction while satisfying the ``subsolar'' interstellar abundance budget suggested by \\citet{Snow+Witt_1996}. However, \\citet{Dwek_1997} argues that this dust model emits too much in the far-IR in comparison with observations. The latest version of the composite model put forward by \\citet{Zubko+Dwek+Arendt_2004} aims at reproducing the observed interstellar extinction and IR emission within a ``subsolar'' abundance budget. The ability of composite grain models to reproduce the observed extinction within ``subsolar'' abundance constraints was recently challenged by \\citet{Li_2005} based on the Kramers-Kronig relations.\\footnote{% \\newtext{ Total (gas~+~dust) interstellar abundances remain uncertain. The most recent determination of the solar oxygen abundance $({\\rm O/H})_\\odot=10^{-3.34\\pm0.05}$ \\citep{Asplund+Grevesse+Sauval_etal_2004} is close to the value for B stars in the Orion association ${\\rm O/H}=10^{-3.31\\pm0.03}$ \\citep{Cunha+Hubeny+Lanz_2006}. The warm interstellar medium gas has ${\\rm O/H}=10^{-3.41\\pm0.01}$ \\citep{Cartledge+Lauroesch+Meyer+Sofia_2006}, consistent with $\\sim$20\\% of the oxygen in silicate grains. For ``refractory'' elements such as Mg, Si, and Fe, however, partial separation of dust and gas by processes such as radiation pressure, ambipolar diffusion, and gravitational sedimentation in the process of star formation could result in stellar abundances differing from interstellar abundances \\citep{Snow_2000, Draine_2004a}. Photospheric abundances may therefore not be a reliable indicator of interstellar abundances for the elements that form refractory grains. } }  In this paper we calculate the IR emission expected from a specific physical model of interstellar dust -- the silicate-graphite-PAH model. In this model, the dust is assumed to consist of a mixture of carbonaceous grains and amorphous silicate grains, with size distributions that are consistent with the observed wavelength-dependent extinction in the local Milky Way \\citep{Weingartner+Draine_2001a}, including different amounts of polycyclic aromatic hydrocarbon (PAH) material.  \\newtext{Spectroscopic observations with {\\it ISO} revealed the rich spectrum of PAHs (see, e.g., \\cite{Beintema+vandenAncker+Molster_etal_1996}; \\cite{Tielens+Hony+vanKerckhoven+Peeters_1999}; \\cite{Joblin+Abergel+Bregman_etal_2000}). } The model \\newtext{presented here} posits \\newtext{``astro}-PAH\\newtext{''} absorption properties that are consistent with spectroscopic observations of PAH emission from dust in nearby galaxies \\citep{Smith+Draine+Dale_etal_2006}. The model calculations show how the IR emission depends on the PAH abundance, and also on the intensity of starlight heating the dust. Comparison of models to infrared photometry obtained with {\\it Spitzer Space Telescope} allows one to estimate PAH abundances, starlight intensities, and total dust masses.  The plan of this paper is as follows:  With new laboratory data [particularly the near-IR absorption spectra of PAH ions measured by \\citet{Mattioda+Hudgins+Allamandola_2005}] and Spitzer spectroscopic data [with new PAH features discovered \\citep{Smith+Dale+Armus+Draine_etal_2004, Werner+Uchida+Sellgren_etal_2004, Smith+Draine+Dale_etal_2006} and the familiar 7.7, 8.6, and 11.3$\\mum$ features resolved into several subfeatures \\citep{Smith+Draine+Dale_etal_2006} we update in \\S\\ref{sec:PAH Cross Sections} the cross sections $\\Cabs(\\lambda)$ that we previously assumed for the PAH particles\\footnote{ Note that ``PAH particles'', ``PAH grains'', and ``PAH molecules'' are synonymous. } \\citep{Li+Draine_2001b}. We distinguish the PAHs by 2 charge states: either neutral (PAH$^{\\rm o}$) or ionized (PAH$^{\\pm}$). We assume that multiply-charged (both positively and negatively) PAHs all have the same cross sections as those of singly-charged cations. \\newtext{The resulting ``astro-PAH'' absorption cross sections do not represent any specific material, but should approximate the actual absorption properties of the PAH mixture in interstellar space.}  In \\S\\ref{sec:Heating of Dust by Starlight} we show examples of temperature distribution functions for both neutral and ionized PAHs heated by starlight. These temperature distribution functions are then used to calculate emission spectra for individual particles in \\S\\ref{sec:Single-Grain Emission Spectra}, which are in turn used to prepare a plot (Figure \\ref{fig:band efficiencies}) showing how the efficiency of emitting into different emission features depends on the PAH size.  When we observe emission from a region in the Milky Way or another galaxy, there will always be a mixture of dust types and sizes present. We consider the emission from a specific set of dust mixtures, as described in \\S\\ref{sec:Dust Mixtures}. All the dust mixtures considered here are consistent with the observed ``average'' extinction curve for diffuse regions in the local Milky Way, but they differ from one another in the assumed abundance of small PAH particles. In \\S\\ref{sec:fir and submm emission} we show that the far-IR and submm emission calculated for the model is consistent with the COBE-FIRAS observations of emission from dust in the local Milky Way. In \\S\\ref{sec:irac band ratios} we show that the calculated IRAC band ratios appear to be consistent with the spectrum of diffuse emission from the Milky Way, as determined by \\citet{Flagey+Boulanger+Verstraete_etal_2006} from Spitzer Space Telescope observations.  The long wavelength emission from the dust model depends on the intensity of the starlight heating the dust.  It will often be the case that the region observed (e.g., an entire star-forming galaxy) will include dust heated by a wide range of starlight intensities. In \\S\\ref{sec:alpha-gamma dust models} we describe a simple parametric model for the distribution of the dust mass between regions with starlight intensities ranging from low to very high. We show in \\S\\ref{sec:dust mass estimation} how observations in the 3 MIPS bands (24$\\mum$, 71$\\mum$, and 160$\\mum$) plus the 3.6$\\mum$ and 7.9$\\mum$ IRAC bands\\footnote{% \\newtext{ The IRAC and MIPS bands are often referred to using nominal wavelengths 3.6, 4.5, 5.8, 8.0, 24, 70, and 160$\\mum$. However, to two significant digits, the effective wavelengths of IRAC bands 3 and 4 are $5.7$ and $7.9\\mum$ (IRAC Data Handbook Version 3.0), and that of MIPS band 2 is $71\\mum$ (MIPS Data Handbook Version 3.2) and we will use these wavelengths to refer to the bands. } } can be used to estimate the parameters describing the distribution of starlight intensities as well as the fraction $\\qpah$ of the total dust mass that is in PAHs, thereby allowing estimation of the total dust mass $\\Mdust$ in the emitting region. A reader interested primarily in applying the results of this paper to interpretation of IRAC and MIPS observations may wish to proceed directly to \\S\\ref{sec:dust mass estimation}.  The results of the paper are discussed in \\S\\ref{sec:discussion} and summarized in \\S\\ref{sec:summary}. ",
        "conclusions": "}  \\subsection{Origin of near-IR Continuum Opacity of PAHs\\label{sec:pah_continuum}}  The present models assume that the small-particle end of the grain size distribution is dominated by PAH particles, as assumed by WD01 and Li \\& Draine (2001b), and the PAH particles are therefore assumed to have a ``continuum'' component to $\\Cabs$ in order to be able to reproduce 2--6$\\mum$ emission seen by ISO \\citep{Lu+Helou+Werner+Dinerstein+Dale_2003} and emission in the IRAC 4.5$\\mum$ filter \\citep{Helou+Roussel+Appleton+Frayer+Stolovy_etal_2004}. Continuum emission also appears to be present at other wavelengths, e.g., between the 12.6 and 16.45$\\mum$ features. In the 2--6$\\mum$ wavelength range, even very small carbonaceous grains are arbitrarily taken to have $\\Cabs$ equal to 1\\% of the continuum that would be produced by free electron absorption in graphite, added to the discrete absorption bands expected for PAHs. The nature of this continuum absorption is unclear. Perhaps some PAHs, e.g., tubular PAHs with appropriate chirality \\citep{Zhou+Sfeir+Zhang_etal_2006} have zero bandgap as in graphite.  We are assuming that the PAHs are responsible for the 2--5\\um\\ continuum. \\citet{An+Sellgren_2003} find that the 2\\um\\ continuum emission and 3.29\\um\\ PAH emission have different spatial distributions in NGC~7023. Perhaps the 2\\um\\ emission is strongest in regions where the PAHs have been dehydrogenated, thereby suppressing the 3.29\\um\\ emission.  The assumption that the continuum emission originates in PAHs is consistent with the fact that the 9.8$\\mum$ and 18$\\mum$ silicate features are generally not seen in emission, except from regions (dusty winds, compact H~II regions, dust near the Trapezium stars in Orion) where the starlight is thought to be sufficiently intense to heat $a\\approx0.1\\micron$ grains to $T\\gtsim 200\\K$. However, \\citet{Li+Draine_2001a} showed that the observed emission from the diffuse interstellar medium in the Milky Way would in fact allow larger amounts of ultrasmall silicate grains than had been estimated previously, because the 9.7$\\mum$ silicate emission feature can be ``hidden'' between the 8.6$\\mum$ and 11.3$\\mum$ PAH features.  In summary, the origin of the 2--5\\um\\ continuum emission remains uncertain.  \\subsection{Near-IR Absorption of PAH Ions: Astrophysical Implications\\label{sec:pah_cabs_ir}}  The far-red to near-IR opacity of PAH ions described in \\citet{Li+Draine_2001b} was derived from earlier experimental data for a small number of small PAH ions [see \\citet{Salama+Bakes+Allamandola+Tielens_1996} and references therein]. More recent laboratory measurements were carried out by \\citet{Mattioda+Hudgins+Allamandola_2005} for 27 PAH cations and anions ranging in size from C$_{14}$H$_{10}$ to C$_{50}$H$_{22}$. The newly-measured absorption of PAH ions at 0.77--2.5$\\mum$ is considerably higher than previously reported. This further supports our earlier conclusion \\citep{Li+Draine_2002b} that visible and near-IR photons are able to excite PAHs to temperatures high enough to emit the mid-IR bands, in agreement with the detection of the PAH features in \\newtext{interstellar} regions lacking UV photons \\citep{Uchida+Sellgren+Werner_1998, Pagani+Lequeux+Cesarsky_etal_1999},\\footnote{% \\newtext{ PAH emission features have also been detected in UV-poor dust debris disks around main-sequence stars SAO\\,206462 [Sp=F8V, $\\Teff\\approx 6250\\K$; \\cite{Coulson+Walther_1995}] and HD\\,34700 [Sp=G0V, $\\Teff\\approx 5940\\K$; \\cite{Smith+Clayton+Valencic_2004}]. However, in an extensive {\\it Spitzer} IRS spectroscopic survey of 111 T Tauri stars in the Taurus star-forming region, \\citet{Furlan+Hartmann+Calvet_etal_2006} found that the PAH emission bands are not seen in dust disks around T Tauri stars of spectral type later than G1. } } and in agreement with the observation by \\citet{Sellgren+Luan+Werner_1990} that the ratio of the IRAS 12$\\mum$ emission (to which the PAH features are the dominant contributor) to the total far-IR surface brightness is independent of the effective temperature $T_{\\rm eff}$ of the exciting stars, for 24 reflection nebulae for $5000\\K\\le T_{\\rm eff} \\le 33000\\K$.   This also relates to our understanding of the origin of interstellar PAHs. At present, the origin and evolution of interstellar PAHs are not very clear. Suggested sources for interstellar PAHs include the injection (into the interstellar medium) of PAHs formed in carbon star outflows \\citep{Jura_1987, Latter_1991}. However, the PAH emission features are commonly not seen in the mid-IR spectra of C-rich AGB stars. The few C stars which display the PAH features all have a hot companion which emits UV photons \\citep{Speck+Barlow_1997, Boersma+Hony+Tielens_2006}. It has been suggested that PAHs are present in all C stars but they are simply not excited sufficiently to emit at mid-IR due to lack of UV photons. The new visible-IR cross sections for PAH ions \\citet{Mattioda+Hudgins+Allamandola_2005} suggest that PAHs can be excited even by the starlight from a C star, in which case the absence of PAH emission places a limit on the abundance of small PAHs in these outflows.\\footnote{% \\newtext{ \\citet{Jura+Bohac+Sargent_etal_2006} reported detection of PAH emission features in HD\\,233517, an evolved oxygen-rich K2III red giant ($\\Teff \\approx 4390\\K$) with circumstellar dust. \\citet{Jura_2003} argued that the IR excess around HD\\,233517 is unlikely to be produced by a recent outflow in a stellar wind. \\citet{Jura+Bohac+Sargent_etal_2006} hypothesized that there is a passive, flared disk orbiting HD\\,233517 and the PAH molecules in the orbiting disk may be synthesized in situ as well as having been incorporated from the ISM. } }  \\subsection{\\label{sec:6.2um absorption feature} Detectability of 6.2\\um\\ Absorption by Interstellar PAHs}  The dust model appropriate to Milky Way dust is thought to have $\\qpah\\approx4.6\\%$ of the dust mass in PAHs with $N_{\\rm C}<10^3$ C atoms. The carbonaceous grains with $a < 50\\Angstrom$ contain 52 ppm of C per H nucleon, divided approximately equally between neutral and ionized PAHs. For this model, we predict a 6.22\\um\\ absorption feature with integrated strength \\beq \\label{eq:tau62 model} \\frac{1}{N_{\\rm H}}\\int \\Delta \\tau ~d\\lambda^{-1} = 6.7\\times10^{-23} \\left(\\frac{\\rm C_{\\rm PAH}/H}{52~\\rm ppm}\\right)\\cm/{\\rm H} ~~~. \\eeq \\citet{Chiar+Tielens_2001} found an upper limit $\\int\\Delta\\tau ~d\\lambda^{-1} < 0.8\\cm^{-1}$ toward Cyg OB2 No.\\,12. For the estimated $N_{\\rm H}=1.9\\times10^{22}\\cm^{-2}$, this gives an upper limit \\beq \\label{eq:tau62 obs} \\frac{1}{N_{\\rm H}}\\int \\Delta \\tau ~d\\lambda^{-1} < 4.2\\times10^{-23}\\cm/{\\rm H} ~~~, \\eeq somewhat smaller than our estimate (\\ref{eq:tau62 model}).  While the present dust model has 52 ppm C in PAHs with $N_{\\rm C}<5\\times10^4$, the 6.22\\um\\ emission is produced mainly by PAHs with $N_{\\rm C}< 500$, which account for only 35 ppm. The 6.2\\um\\ absorption contributed by PAHs with $N_{\\rm C}< 500$ in our model is consistent with the \\citet{Chiar+Tielens_2001} upper limit.  The absolute PAH band strengths in Table \\ref{tab:drude} are only estimates. Because the 3--25\\um\\ emission spectrum for $U=1$ is almost entirely the result of single-photon heating, the emission spectrum would be almost unaffected if all the PAH band strengths were uniformly reduced by a common factor. However, based on the comparison with theoretical calculations shown in Figure \\ref{fig:Cabs_vs_theory}, the actual band strengths should be similar in magnitude to the values adopted in Table \\ref{tab:drude}. Spectroscopy of heavily extinguished stars using the IRS instrument on Spitzer Space Telescope should be carried out to attempt to detect the predicted 6.22\\um\\ absorption from interstellar PAHs.   \\subsection{Deuterated PAHs\\label{sec:dpah}}  Since the aromatic C--D bond zero-point energy is about 30\\% lower than that of the C--H bond, H loss is favored over D loss. The PAH D/H ratio might therefore be substantially larger than the \\newtext{total} D/H~$\\approx2\\times10^{-5}$ \\citep{Allamandola+Tielens+Barker_1989}. The observed sightline-to-sightline variations of (D/H) in interstellar gas have been interpreted as requiring depletion of D by interstellar dust \\citep{Linsky+Draine+Moos_etal_2006}. \\citet{Draine_2004a,Draine_2006b} has suggested that this depletion of D may take via interactions of PAH ions with the gas, and may imply D/H$\\approx0.3$ in interstellar PAHs. \\citet{Peeters+Allamandola+Bauschlicher_etal_2004} report tentative detection of emission at 4.4 and 4.65\\um\\ from deuterated PAHs from the Orion Bar and M17 PDRs. If the emission is in fact due to C-D stretching modes, \\citet{Peeters+Allamandola+Bauschlicher_etal_2004} estimate ${\\rm D/H}\\approx0.17\\pm0.03$ for the emitting PAHs in the Orion Bar and $0.36\\pm0.08$ in M17.  Our model for the PAH absorption cross section is based on observed emission spectra, and some of the modeled absorption may be associated with C-D bending modes if interstellar PAHs are appreciably deuterated. Because the observations of 4.4 and 4.65\\um\\ emission in the Orion Bar and M17 are tentative, we do not attempt here to include any absorption and emission that might result from C-D stretching modes in the 4.5\\um\\ region.  \\subsection{Ubiquity and Absence of PAHs in Astrophysical Regions: Rationale for Variable $q_{\\rm PAH}$\\label{sec:reason_4_qpah}}  The ISO and Spitzer imaging and spectroscopy have revealed that PAHs are a ubiquitous feature of both the Milky Way and external galaxies. Recent discoveries include the detection of PAH emission in a wide range of systems: distant Luminous Infrared Galaxies (LIRGs) with redshift $z$ ranging from 0.1 to 1.2 \\citep{Elbaz+LeFloch+Dole+Marcillac_2005}; distant Ultraluminous Infrared Galaxies (ULIRGs) with redshift $z\\sim$\\,2 \\citep{Yan+Chary+Armus_etal_2005}; distant luminous submillimeter galaxies at redshift $z\\sim$\\,2.8 \\citep{Lutz+Valiante+Sturm_etal_2005}; elliptical galaxies with a hostile environment (containing hot gas of temperature $\\sim$\\,10$^7\\K$) where PAHs can be easily destroyed through sputtering by plasma ions \\citep{Kaneda+Onaka+Sakon_2005}; faint tidal dwarf galaxies with metallicity $\\sim Z_\\odot/3$ \\citep{Higdon+Higdon+Marshall_2006}; and galaxy halos \\citep{Irwin+Madden_2006, Engelbracht+Kundurthy+Gordon_etal_2006}.  However, the PAH features are weak or even absent in low-metallicity galaxies and AGN: \\begin{itemize} \\item Based on an analysis of the mid-IR spectra of 60 galaxies obtained by ground-based observations, \\citet{Roche+Aitken+Smith+Ward_1991} reported the lack of PAH emission features in AGN. \\item The Small Magellanic Cloud (SMC), an irregular dwarf galaxy with a metallicity just $\\sim$\\,1/10 of solar, exhibits a local minimum at $\\sim$\\,12$\\mum$ in its IR emission spectrum, suggesting the lack of PAHs \\citep{Li+Draine_2002c}. \\item \\citet{Thuan+Sauvage+Madden_1999}, \\newtext{\\citet{Plante+Sauvage_2002},} and \\citet{Houck+Charmandaris+Brandl_etal_2004} have shown that the PAH features are absent in the ISO and Spitzer spectra of SBS\\,0335-052, a metal-poor ($Z\\sim$\\,$Z_{\\odot}/41$) blue compact dwarf galaxy. \\item More recently, \\citet{Hunt+Bianchi+Maiolino_2005}, \\citet{Madden+Galliano+Jones+Sauvage_2006}, and \\citet{Wu+Charmandaris+Hao_etal_2006} performed a more systematic investigation of the mid-IR spectra of a large number of low-metallicity galaxies (with $Z/Z_\\odot$ ranging from 0.02 to 0.6) obtained with ISO and Spitzer. They found the PAH features are substantially suppressed in metal-poor dwarf galaxies. \\item \\newtext{ Using Spitzer IRAC [3.6]--[7.9] colors and {\\it Sloan Digital Sky Survey} (SDSS) [g--r] colors of 313 visible-selected SDSS main sample galaxies, \\citet{Hogg+Tremonti+Blanton_etal_2005} found that low-luminosity galaxies show a deficiency in PAH emission, with weak evidence for a dependence of the PAH-to-stellar radiation ratio on metallicity. \\citet{Rosenberg+Ashby+Salzer+Huang_2006} performed a similar study for a statistically complete sample of 19 star-forming dwarf galaxies, but they found a significant number of low-luminosity galaxies with very red [3.6]--[7.9] colors, indicating the presence of PAHs (and/or hot dust). They also found that the 7.9$\\mum$ emission is more strongly correlated with the star formation rate than it is with the metallicity. } \\item \\citet{Engelbracht+Gordon+Rieke_etal_2005} examined the Spitzer mid-IR colors of 34 galaxies ranging over 2 orders of magnitude in metallicity. They found that the 8\\,$\\mu$m-to-24\\,$\\mu$m color changes abruptly from $F_\\nu(8\\mu {\\rm m})$/$F_\\nu(24\\mu {\\rm m})$$\\sim$\\,0.7 for galaxies with $Z/Z_\\odot$\\,$>$\\,1/3 to $F_\\nu(8\\mu {\\rm m})$/$F_\\nu(24\\mu {\\rm m})$$\\sim$\\,0.08 for galaxies with $Z/Z_\\odot$\\,$<$\\,1/5--1/3. They attributed this color shift to a decrease in the 7.7$\\mum$ PAH feature at low metallicity. \\item  The silicate-graphite-PAH model described in this paper has been employed to quantitatively determine the fraction $\\qpah$ of the dust mass contributed by PAHs for 61 galaxies in the SINGS survey \\citep{Draine+Dale+Bendo_etal_2006}. There appears to be a threshold metallicity: galaxies with [O/H]$\\equiv 12+\\log_{10}({\\rm O/H})< 8.1$ (i.e., $Z < Z_\\odot/3$) have $\\qpah< 1.5\\%$. \\end{itemize}  Knowledge of how $q_{\\rm PAH}$ varies as a function of galaxy parameters (such as metallicity, luminosity, and morphological type) will have a broad impact -- \\begin{itemize} \\item It will be essential for properly interpreting the IR cosmological surveys which now often assume a {\\it constant} PAH strength for galaxies [see, e.g., \\citet{Lagache+Dole+Puget_etal_2004}]. \\item It may provide physical insight into the validity of using the line-to-continuum ratio of the 7.7$\\mum$ PAH feature as a discriminator between starburst and AGN activity in ULIRGs, i.e., whether the dominant luminosity source of ULIRGs is an AGN or a starburst; see \\citet{Genzel+Lutz+Sturm_etal_1998}. \\item As local analogs of early galaxies at high redshift which must have formed at very low metallicity, nearby low metallicity systems can provide insight into the early stages of galaxy evolution. The low PAH abundances observed in these galaxies reflect the balance of PAH formation and destruction, and thereby is a valuable probe of the physical conditions following the onset of star formation in primordial gas. \\item By comparing with the star-formation rates derived from other tracers, it will allow us to quantitatively evaluate the validity of using the IRAC 8$\\mum$ photometry as a reliable tracer for star-formation rates. It is generally believed that the IRAC 8$\\mum$ flux is stronger in regions with stronger star-forming activities; however, observations suggest that PAHs are destroyed in star-forming regions with very strong and hard radiation fields \\citep{Contursi+Lequeux+Cesarsky_etal_2000, ForsterSchreiber+Roussel+Sauvage+Charmandaris_2004, Beirao+Brandl+Devost_etal_2006}.% \\end{itemize}  \\subsection{Diversity of PAH Emission Spectra\\label{sec:diverse}}  The PAH absorption and emission properties presented in this series of papers are for ``astronomical'' PAHs -- we fix the peak wavelength and bandwidth of each PAH band to what is observed in most astronomical regions. The strength of each PAH band (on a per carbon or hydrogen atom basis) is also fixed for a specific charge state to be consistent with what is observed in space, measured in the laboratory, or calculated theoretically.  However, there exist variations in the central wavelength, bandwidth and strength of PAH emission bands among different astronomical objects or among different locations within one object: \\begin{itemize} \\item \\citet{Cohen+Tielens+Bregman_etal_1989} reported that the central wavelength of the 7.7$\\mum$ band varies among different types of objects: it shifts from shorter wavelengths for objects with heavily processed PAHs (e.g. HII regions, reflection nebulae) to longer wavelengths for objects with freshly created PAHs (e.g. planetary nebulae). This wavelength shift has been confirmed by high resolution spectroscopy obtained with the ISO Short Wavelength Spectrometer \\citep{Peeters+Hony+vanKerckhoven_etal_2002}. Similar shifts have also been seen in other bands [e.g. see \\citet{vanDiedenhoven+Peeters+vanKerckhoven_etal_2004}]. More recently, \\citet{Bregman+Temi_2005} observed a progressive blueshift in the central wavelength of the 7.7$\\mum$ band within 3 reflection nebulae from the edge of the nebulae to places closer to the exciting stars. \\item \\citet{Uchida+Sellgren+Werner+Houdashelt_2000} reported a progressive broadening of the FWHM of the 7.7$\\mum$ band moving toward the exciting star in the reflection nebula vdB 17. This broadening has also been reported for the reflection nebula Ced 21 \\citep{Cesarsky+Lequeux+Ryter+Gerin_2000}. But such a broadening has not been seen in other PAH bands. \\item \\citet{Roelfsema+Cox+Tielens_etal_1996} reported that the 8.6$\\mum$ band of the compact HII region IRAS 18434-0242 is almost twice as strong as the 7.7$\\mum$ band, while for most objects the 7.7$\\mum$ band is stronger than the 8.6$\\mum$ band by a factor of $\\sim$\\,5--10. \\item \\newtext{ \\citet{Reach+Boulanger+Contursi+Lequeux_2000} reported that the 11.3$\\mum$/7.7$\\mum$ band ratio of the PAH emission spectrum of the quiescent molecular cloud SMC B1 No.\\,1 in the SMC is much higher than that of any other Galactic or extragalactic objects. (see Figs.\\,16 and 17 of \\citet{Draine+Li_2001}). } \\end{itemize}  While the band-ratio variations can presumably be accounted for by the combined effects of variations in the radiation exciting the PAHs and variations in PAH sizes and ionization state (small PAHs emit more strongly at the 6.2, 7.7, 8.6$\\mum$ bands, while relatively large PAHs emit more strongly at the 11.3$\\mum$ band; neutral PAHs are much stronger 3.3$\\mum$ and 11.3$\\mum$ emitters and weaker 6.2, 7.7 and 8.6$\\mum$ emitters compared to their charged counterparts), we do not intend to confront our PAH model with extreme spectral variations, such as the observed spectrum of IRAS 18434-024 (with the 8.6$\\mum$ band almost twice as strong as the 7.7$\\mum$ band) which, according to \\citet{Roelfsema+Cox+Tielens_etal_1996}, is possibly rich in non-compact PAHs. Because of the way we design the ``astronomical'' PAH model, the present model is not able to account for the central wavelength shifts and the 7.7$\\mum$ band broadening, although experimental investigations have been carried out [e.g. see \\citet{Hudgins+Allamandola_1999}].  \\citet{Hudgins+Allamandola_2005} discuss recent work toward PAH identification, with attention to the effects of N substitution. Future, more realistic modeling will require more extensive knowledge of PAH properties than is now available from laboratory studies [e.g., \\citet{Hudgins+Allamandola_1999, Mattioda+Hudgins+Allamandola_2005, Mattioda+Allamandola+Hudgins_2005}] and theoretical calculations [e.g., \\citet{Bakes+Tielens+Bauschlicher_2001, Bakes+Tielens+Bauschlicher_etal_2001, Malloci+Mulas+Joblin_2004, Mulas+Malloci+Joblin+Toublanc_2006, Malloci+Joblin+Mulas_2006}].  \\subsection{Other Grain Models\\label{sec:others}}  The dust model used in this study is, of course, provisional. The model assumes the amorphous silicate and carbonaceous material to be in physically separate grains.  While this appears to be consistent with observations, including the observed absence of polarization in the 3.4\\um\\ absorption feature \\citep{Chiar+Adamson+Whittet_etal_2006}, it is by no means observationally established.  Our models have also used the dielectric tensor of crystalline graphite to estimate the infrared absorption by carbonaceous grains with $a\\gtsim0.01\\micron$.  Some form of amorphous carbon may provide a better approximation for interstellar carbonaceous material that may be formed from agglomeration of PAHs followed by cross-linking due to ultraviolet photolysis and cosmic ray exposure. The main effect on the present models would be to do away with the broad absorption feature near 30\\um\\ that appears in the calculated $\\Cabs(\\lambda)$ for graphitic material, as seen in Figure \\ref{fig:Cabs for carbonaceous grains}, and the corresponding broad emission feature near 30\\um\\ when the starlight intensity is high enough to heat $a\\gtsim0.01\\micron$ carbonaceous grains to $T\\gtsim10^2\\K$.  In this work we focus on the silicate-graphite-PAH model. For other dust models, if they are tailored to reproduce the Milky Way interstellar extinction% {\\it and} albedo in the optical/UV {\\it and} meanwhile satisfy plausible interstellar abundance constraints, the derived parameters $\\Mdust\\langle U\\rangle$ will be nearly the same as derived from the silicate-graphite-PAH model, simply from energy balance considerations: the energy emitted in the IR must be absorbed by the same dust in the optical/UV. If the dust model reproduces the observed far-infrared emission spectrum of Milky Way dust for $U\\approx 1$, then the temperatures of the larger dust grains will scale with changes in $U$ in a manner similar to the present model, and therefore the inferred properties of the radiation field ($\\Umin$, $\\gamma$) will presumably be similar.  However, the composite model will probably have a lower optical/UV albedo than observed \\citep{Dwek_1997}. Because the extinction per H atom is fixed by observations, the lower albedo will result in more total far-infrared emission per H nucleon. Given the uncertainties regarding both the intensity of interstellar starlight and the far infrared emission from dust, composite grain models must be considered to be viable. Additional observational tests to distinguish between dust models are needed.  }  The principal results of this study are as follows: \\begin{enumerate}  \\item The wavelength dependence of absorption by interstellar PAHs has been revised to better reproduce oberved emission spectra. The adopted $C_{\\rm abs}(\\lambda)$ are generally consistent with measured or calculated absorption cross sections for PAHs (Figure \\ref{fig:Cabs_vs_theory}).  \\item Temperature distribution functions have been calculated for PAHs of different sizes, heated by starlight with the spectrum estimated for the local interstellar radiation field. These temperature distribution functions have been used to compute the time-averaged emission spectra for PAHs of different sizes (Figure \\ref{fig:PAHi and PAHn emission spectra}).  \\item Figures \\ref{fig:band efficiencies} and \\ref{fig:band efficiencies, dism} show the efficiency with which PAHs radiate absorbed energy into different emission features.  The strong 7.7$\\mum$ feature is radiated efficiently only by PAHs with $N_{\\rm C}\\ltsim 10^3$ carbon atoms -- larger PAHs radiate primarily at longer wavelengths.  \\item Emission spectra have been calculated for dust models that reproduce the average Milky Way extinction, using size distributions with different abundances of small PAHs (see Fig.\\ \\ref{fig:variousqpah_spec}).  \\item Emission spectra have been calculated for dust mixtures heated by different starlight intensities (see Fig.\\ \\ref{fig:variousU_spec}).  \\item The dust model used in this paper is found to be in good agreement with observed IRAC~3.6\\um/ IRAC~7.9\\um\\ and IRAC~5.7\\um/ IRAC~7.9\\um\\ band ratios (see Fig.\\ \\ref{fig:irac band ratios}), but underpredicts the observed IRAC~4.5\\um/ IRAC~7.9\\um\\ ratio by about a factor 1.5.  \\item The far-infrared and submm emission spectrum calculated for the present model appears to be in good agreement with the observed emission from dust in the local Milky Way (see Figure \\ref{fig:firspec}).  \\item A prescription (eq.\\ \\ref{fig:TIR estimator}) is found that allows the total IR emission to be estimated from observed fluxes in the IRAC~7.9\\um\\ band and the 3 MIPS bands. For dust heated by starlight intensities $0.1 < U \\ltsim 10^2$, eq.\\ (\\ref{fig:TIR estimator}) is accurate to within $\\sim10\\%$.  \\item Emission spectra have been calculated for dust mixtures heated by power-law distributions of starlight intensities (see Fig.\\ \\ref{fig:spec for Udists}).  \\item A distribution function (eq.\\ \\ref{eq:dM_d/dU}) is proposed for representing the distribution of starlight intensities heating dust in a galaxy, or region within a galaxy.  \\item Graphical procedures are presented that allow dust model parameters to be found using 3 ratios ($P_{7.9}$, $P_{24}$, and $R_{71}$) constructed from fluxes measured in IRAC band 4 and MIPS bands 1--3.  The graphs provided here (Figs. \\ref{fig:IRAC4/MIPS band ratios} and \\ref{fig:MIPS band ratios}) allow model parameters $q_{\\rm PAH}$, $\\Umin$, and $\\gamma$ to be estimated. \\item The fraction $f_{\\rm PDR}$ of the dust infrared emission that is contributed by dust in photodissociation regions with $U > 10^2$ can be estimated from IRAC and MIPS photometry using Figure \\ref{fig:fpdr}. \\item Using a coefficient $\\Psi$ obtained from Figure \\ref{fig:Psi grid}, the total dust mass $\\Mdust$ can be estimated from fluxes in the 3 MIPS bands using eq.\\ (\\ref{eq:compute Mdust}). \\end{enumerate}"
    },
    "0608/astro-ph0608529_arXiv.txt": {
        "abstract": "We present a first detailed spectrum synthesis calculation of a supernova-fallback disk composed of iron. We assume a geometrically thin disk with a radial structure described by the classical $\\alpha$-disk model. The disk is represented by concentric rings radiating as plane-parallel slabs. The vertical structure and emission spectrum of each ring is computed in a fully self-consistent manner by solving the structure equations simultaneously with the radiation transfer equations under non-LTE conditions. We describe the properties of a specific disk model and discuss various effects on the emergent UV/optical spectrum.  We find that strong iron-line blanketing causes broad absorption features over the whole spectral range. Limb darkening changes the spectral distribution up to a factor of four depending on the inclination angle. Consequently, such differences also occur between a blackbody spectrum and our model. The overall spectral shape is independent of the exact chemical composition as long as iron is the dominant species. A pure iron composition cannot be distinguished from silicon-burning ash. Non-LTE effects are small and restricted to few spectral features. ",
        "introduction": "Anomalous X-ray pulsars (AXPs) are slowly rotating ($P_{\\rm rot}=5-12$~s) young ($\\leq 100\\,000$~yr) isolated neutron stars. Their X-ray luminosities ($\\approx 10^{36}$~erg/s) greatly exceed  the rates of rotational energy loss ($\\approx 10^{33}$~erg/s). It is now generally believed that AXPs are magnetars with magnetic field strengths greater than $10^{14}$~G and that their X-ray luminosity is powered by magnetic energy (Woods \\& Thompson 2006). As an alternative explanation the X-ray emission was attributed to accretion from a disk that is made up of supernova-fallback material (van Paradijs \\etal 1995, Chatterjee \\etal 2000, Alpar 2001). The fallback-disk model has difficulties to explain IR/optical emission properties of AXPs.  When compared with disk models, the faint IR/optical flux suggests that any disk around AXPs must be very compact (e.g.\\ Perna \\etal 2000, Israel \\etal 2004).  The discovery of optical pulsations in 4U\\,0142+61 which have the same period like the X-ray pulsations (Kern \\& Martin 2002) appears to be a strong argument against the disk model. It was argued that reprocessing of the pulsed NS X-ray emission in the disk cannot explain the high optical pulsed fraction, because disk radiation would be dominated by viscous dissipation and not by reprocessed NS irradiation (Kern \\& Martin 2002). Ertan \\& Cheng (2004), on the other hand, showed that these optical pulsations can be explained either by the magnetar outer gap model or by the disk-star dynamo model. Therefore, the observation of optical pulsations is not an argument against the disk model. A spectral break in the optical spectrum of 4U\\,0142+61 was discovered by Hulleman \\etal (2004). This was also taken as an argument against the disk model because the authors do not expect such strong features from a thermally emitting disk.  The recent discovery of mid-IR emission from this AXP (Wang \\etal 2006), however, has strongly rekindled the interest in studies of fallback-disk emission properties. While this mid-IR emission was attributed to a cool, passive (X-ray irradiated) dust debris disk by Wang \\etal (2006), it was shown by Ertan on this conference that it can be explained with a model for an active, dissipating gas disk. If true, then the disk emission properties allow to conclude on important quantities, e.g., the magnetic field strength of the neutron star can be derived from the inner disk radius.  Independent hints for the possible existence of fall\\-back disks come from pulsars with particular spin-down properties. For example, the discrepancy between the characteristic age and the supernova age of the pulsar B1757-24  was explained by the combined action of magnetic dipole radiation and accretion torques (Marsden \\etal 2001).  Even more, the presence of jets from pulsars such like the Crab and Vela can possibly be explained by disk-wind outflow interacting with and collimating the pulsar wind (Blackman \\& Perna 2004).  A fallback-disk model was proposed in order to explain the X-ray enhancement following a giant flare of the Soft Gamma Repeater SGR\\,1900+14 (Ertan \\& Alpar 2003). The X-ray light curve is interpreted in terms of the relaxation of a fallback disk that has been pushed back by the gamma-ray flare. This model can also explain the long-term X-ray and IR enhancement light curves of the AXP 1E\\,2259+58 following a major bursting epoch (Ertan \\etal 2006).  The presence of a fallback disk around the stellar remnant of SN~1987A has been invoked in order to explain its observed lightcurve which deviates from the theoretical one for pure radioactive decay (Meyer-Hofmeister 1992). From the non-detection of any UV/optical point source in the supernova remnant, however, tight constraints for the disk extension can be derived (Graves \\etal 2005).  To our best knowledge, the emission from fallback disks in all studies was hitherto modeled with blackbody spectra. In view of the importance of disk models for the quantitative interpretation of observational data it is highly desirable to construct more realistic models by detailed radiation-transfer calculations. ",
        "conclusions": "We have computed a model for a supernova-fallback disk in order to study its structure and optical/UV emission properties. We assumed an $\\alpha$-disk for the radial structure and performed detailed non-LTE radiation transfer calculations for the vertical structure and spectrum synthesis. The input parameters were: \\\\ \\\\ Neutron star mass:  1.4~\\msun \\\\ Inner and outer disk edge radii: $R$=2000 and 200\\,000~km \\\\ Mass-accretion rate: $\\dot{M}=3\\cdot 10^{-9}$~\\msun/yr.  \\\\ \\\\ We have identified that the disk region in the vicinity of $R$=40\\,000~km is the main contributor to the total disk spectrum at UV/optical wavelengths. We therefore investigated in some detail the disk properties at this radius.  We summarize our results as follows:  \\begin{itemize}  \\item The overall disk spectrum is independent of the detailed chemical composition as long as iron is the dominant species. In particular, a pure-iron composition is spectroscopically indistinguishable from a silicon-burning ash composition.  \\item The overall disk spectrum is hardly influenced by non-LTE effects, however, equivalent widths of individual line blends can change by a factor of two.  \\item Limb darkening affects the overall disk spectrum (in addition to the geometric foreshortening factor $\\cos i$). Depending on inclination and spectral band, the disk intensity varies up to a factor of three.  \\item Depending on the inclination, the disk flux can be a factor of two higher or lower compared to a blackbody radiating disk.  \\item Strong iron line blanketing causes broad ($>100$~\\AA) spectral features that could be detectable even from almost edge-on disks. Disks that are cooler than our model (because of a lower mass-accretion rate) could even exhibit a spectral break in the optical band due to massive line blanketing.  \\end{itemize}  We stress that these results hold strictly only for our particular disk model. In order to arrive at more general results a systematic parameter study (disk extent, accretion rate) of the disk emission is necessary. Also, it needs to be investigated in detail how fine the subdivision of the disk in a number of rings is necessary in order to achieve a computed spectrum with a certain accuracy.  In addition, deviations from the $\\alpha$-disk model must be studied. Another important point will be the inclusion of disk irradiation by the X-ray emission from the neutron star. This will reveal the relative importance of viscous dissipation and reprocessed irradiation that is discussed in the context of simultaneous optical and X-ray pulsations in the AXP 4U\\,0142+61. At the moment we do not dare to make any prediction how this affects the results presented here.  The innermost disk ring has a very high effective temperature and its flux distribution peaks in the soft X-ray band. It needs to be investigated systematically under which conditions (inclination, inner disk radius, accretion rate) the innermost disk regions can contribute to the thermal spectrum of the magnetars."
    },
    "0608/astro-ph0608235_arXiv.txt": {
        "abstract": "We explore the infrared spectrum of a three-dimensional dynamical model of planet \\hd~as a function of orbital phase.  The dynamical model predicts day-side atmospheric pressure-temperature profiles that are much more isothermal at pressures less than 1 bar than one-dimensional radiative-convective models have found.  The resulting day-side thermal spectra are very similar to a blackbody, and only weak water absorption features are seen at short wavelengths.  The dayside emission is consequently in significantly better agreement with ground-based and space-based secondary eclipse data than any previous models, which predict strong flux peaks and deep absorption features.  At other orbital phases, absorption due to carbon monoxide and methane is also predicted.  We compute the spectra under two treatments of atmospheric chemistry: one uses the predictions of equilibrium chemistry, and the other uses non-equilibrium chemistry, which ties the timescales of methane and carbon monoxide chemistry to dynamical timescales.  As a function of orbital phase, we predict planet-to-star flux ratios for standard infrared bands and all \\emph{Spitzer Space Telescope} bands.  In \\emph{Spitzer} bands, we predict 2-fold to 15-fold variation in planetary flux as a function of orbital phase with equilibrium chemistry, and 2-fold to 4-fold variation with non-equilibrium chemistry.  Variation is generally more pronounced in bands from 3-10 $\\mu$m than at longer wavelengths.  The orbital phase of maximum thermal emission in infrared bands is 15--45 orbital degrees before the time of secondary eclipse.  Changes in flux as a function of orbital phase for \\hd~should be observable with \\emph{Spitzer}, given the previously acheived observational error bars. ",
        "introduction": "Astronomers and planetary scientists are just beginning to understand the atmospheres of the short period giant planets known as ``hot Jupiters'' or ``Pegasi planets.''  A key subset of these planets are those that transit the disk of their parent stars, which make them well-suited for follow-up studies.  The most well studied of these transiting hot Jupiters is the first to be discovered, \\hd~\\citep{Charb00,Henry00}, a 0.69 \\mj~planet that orbits its Sun-like parent star at a distance of 0.045 AU.  Currently, considerable work on hot Jupiters is occuring on both the observational and theoretical fronts.  In the past few years, several groups have computed dynamical atmosphere models for \\hd~in an effort to understand the structure, winds, and temperature contrasts of the planet's atmosphere \\citep{Showman02,Cho03,CS05,Burkert05,CS06}.  If the planet has been tidally de-spun and become locked to its parent star, dynamical models are surely needed to understand the extent to which absorbed stellar energy is transported onto the planet's permanent night side.  With the launch of the \\emph{Spitzer Space Telescope}, there is now a platform that is well-suited for observations of the thermal emission from hot Jupiter planets.  \\emph{Spitzer} observations spanning the time of the planet's secondary eclipse (when the planet passes behind its parent star) have been published for \\hd~at 24 $\\mu$m \\citep{Deming05b}, \\T~at 4.5 and 8.0 $\\mu$m \\citep{Charb05}, and \\he~at 16$\\mu$m \\citep{Deming06}.  In all cases, the observed quantity is the planet-to-star flux ratio in \\emph{Spitzer} InfaRed Array Camera (IRAC), InfraRed Spectrograph (IRS), or Multiband Imaging Photometer for \\emph{Spitzer} (MIPS) bands.  The dual \\T~observations are especially interesting because they allow for a determination of the planet's mid-infrared spectral slope.  Several efforts have also been made from the ground to observe the secondary eclipse of \\hd.  Although no detections have been made, some important, occasionallly overlooked upper limits at K and L band have been obtained.  These include \\citet{Rich03b} around 2.3 $\\mu$m, \\citet{Snellen05} in K band, and \\citet{Deming05c} at L band.  These ground-based observations constrain the flux emitted by the planet in spectral bands where water vapor opacity is expected to be minimal; therefore, emitted flux should be high.  This influx of data has spurred a new generation of radiative-convective equilibrium models, whose resulting infrared spectra can be compared with data \\citep{Fortney05, Burrows05b, Seager05, Barman05, Fortney06, Burrows06}.  See \\citet{Marley06} and \\citet{Charb06} for reviews.  The majority of these models are one-dimensional. Authors weight the incident stellar flux by 1/4, to simulate planet-wide average conditions, or by 1/2, to simulate day-side average conditions (with a cold night side).  \\citet{Barman05} have investigated two-dimensional models with axial symmetry around the planet's substellar-antistellar axis and computed infrared spectra as a function of orbital phase. \\citet{Iro05} have extended one-dimensional models by adding heat transport due to a simple parametrization of winds to generate longitude-dependent temperature maps, but they did not compute disk averaged spectra for these models.  Very recently \\citet{Burrows06} have also investigated spectra and light curves of planets with various day-night effective temperature differences, assuming 1D profiles for each hemisphere.  These one- and two-dimensional radiative-convective equilibrium models have had some success in matching \\emph{Spitzer} observations, but \\citet{Seager05} and \\citet{Deming05c} have shown that ground-based data for \\hd~do not indicate promiment flux peaks at 2.3 and 3.8 $\\mu$m, which solar composition models predict.  The various dynamical models for \\hd~\\citep{Showman02,Cho03,CS05,Burkert05,CS06} are quite varied in their treatment of the planet's atmosphere.  We will not review them here, as that is not the focus of this paper.  In general, temperature contrasts in the visible atmosphere are expected to be somewhere between 300 and 1000 K.  What has been somewhat lacking for these dynamical models are clear observational signatures, which would in principle be testable with \\emph{Spitzer} or other telescopes.  The purpose of this paper is to remedy that situation.  Here we generate infrared spectra and light curves as a function of orbital phase for the \\citet{CS06} (hereafter: CS06) dynamical simulation. We present the first spectra generated for three-dimensional models of the atmosphere of \\hd. ",
        "conclusions": "\\subsection{Effective Temperature} For the CS06 dynamical model of the atmosphere of \\hd, we find that the apparent $T_{\\rm eff}$ of the visible hemisphere is strongly variable, with a maximum of 1390 K and a minimum of 915 K for equilibrium chemistry and 1025 K for (the probably more realistic case of) disequilibrium chemistry.  This leads to a luminosity of the visible planetary hemisphere that varies by factors of 5.3 and 3.4, respectively for these two cases.  For the one dimensional \\citet{Iro05} planet-wide  profile, we derive a $T_{\\rm eff} = 1325$ K, 1\\% less than found by the authors.  The day-side $T_{\\rm eff}$ for the dynamical model is not significantly larger than this planet-wide $T_{\\rm eff}$.  As an additional comparison, we calculate the mean of the planetary luminosity over the entire orbit and then convert to $T_{\\rm eff}$ to find a mean $T_{\\rm eff}$ for Cases 0 and 1.  This gives $T_{\\rm eff}$s of 1195 K and 1227 K, respectively, showing that the CS06 model, at the pressures levels that radiate to space, is as a whole somewhat colder than the one dimensial profile of \\citet{Iro05}.  This is likely a consequence of the radiative forcing scheme employed in CS06, which will be reinvestigated when models that consistently couple radiative transfer and dynamics are developed.  The $\\sim$400 K $T_{\\rm eff}$ contrasts found are predominantly due to changes in the temperature structure of the hemisphere that is visible as function of orbital phase.  In addition, atmospheric opacity makes an important contribution when the CH$_4$/CO ratio is free to vary with pressure and temperature.  When the CH$_4$ mixing ratio is below that predicted by equilibrium chemistry, this leads to a lower opacity atmosphere, for a given \\emph{P--T} profile.  While the few individual bands of CO are somewhat stronger than those of CH$_4$, CH$_4$ absorption across the planet's broad the 2-10 $\\mu$m flux peak dominates over the two bands of CO.  One is able to see more deeply into a CH$_4$-depleted atmosphere, leading to a higher $T_{\\rm eff}$.  \\subsection{\\hd~Infrared Data} Currently, the only published secondary eclipse datum from \\emph{Spitzer} for \\hd~is the 24 $\\mu$m detection of \\citet{Deming05b}.  The models presented here have a planet-to-star flux ratio during secondary eclipse that is 1.6 sigma larger than this observational data point.  Together with our excellent agreement with the ground-based data at 2.3 $\\mu$m from \\citet{Rich03b} and \\citet{Seager05}, and our 1.1 sigma difference with the 3.8 $\\mu$m data from \\citet{Deming05c}, we regard this as excellent agreement---significantly better than has been previously obtained with one dimensional radiative equilibrium models.  It is interesting to discuss a few issues that arise if the flux ratios are indeed 10-25\\% less than we calculate here, as indicated by the L$^{\\prime}$ and 24 $\\mu$m band data.  For instance, perhaps day-night temperature contrasts in the atmosphere are not as large as predicted by CS06, leading to smaller deviations from a ``planet-wide'' $T_{\\rm eff} \\sim 1200$K. This may involve radiative time constants that are significantly longer than predicted by \\citet{Iro05}, winds faster than predicted by CS06, or both.  It would be worthwhile for other groups that possess hot Jupiter radiative transfer codes to compute radiative time constants at these temperatures and pressures.  This is an area that we will pursue in the near future.  Another possibility for a smaller planet-to-star flux ratio during secondary eclipse would be if the planet had a larger Bond albedo than calculations currently indicate.  This would mean less absorbed stellar flux and correspondingly lower temperatures everywhere on the planet.  Cooler temperatures everywhere on the planet would lead to chemical abundances that differ from our previous cases.  More CH$_4$ would form, at the expense of CO, which would also lead to a slightly higher mixing ratio for H$_2$O, which shares oxygen with CO.  Hot Jupiters are believed to have very low Bond albedos--on the order of 90\\% or more of incident stellar flux is expected to be absorbed.  In \\citet{Fortney05}, we found that our one-dimensional model atmosphere for \\hd~scattered only 8\\% of incident flux.  For \\T, this was 6\\%.  To date, there is at least a hint that the Bond albedo of \\T~may have been underestimated.  \\citet{Charb05}, under the assumption that the planet emits as a blackbody, determined a Bond albedo of $0.31 \\pm 0.14$ from their \\emph{Spitzer} IRAC observations at 4.5 and 8.0 $\\mu$m.  Perhaps hot Jupiters are not quite as hot as had been previously thought.  However, this determination should be considered very preliminary at this time.  \\citet{Richardson06} have recently observed the \\emph{transit} of \\hd~at 24 $\\mu$m, as well.  From an observed transit depth of $0.0149 \\pm 0.0003$, they determined the radius of the planet in this band to be 1.26 $\\pm 0.08$ \\rj, which includes uncertainties in the stellar radius.  Our model predicts a change in the apparent planet-to-star flux ratio of $\\sim$0.00008 during the 20 orbital degrees of the \\citet{Richardson06} transit observations, $\\sim$4 times smaller than their uncertainty, and hence too small to have been detected.  As was seen in \\mbox{Figures~\\ref{figure:spitzer}} and \\ref{figure:mko}, the change in planetary flux as a function of orbital phase is not as pronounced near the time of transit (and secondary eclipse) as it is at other phases.  \\subsection{Temperature Sensitivity} To illustrate the sensitivity of our results to temperature changes, we have computed spectra as a function of orbital phase, using equilibrium (Case 0) chemistry for two additional models.  These are additional dynamical models described in CS06, in which the base \\emph{P--T} profile of \\citet{Iro05} has been increased or decreased by 300 K, with correspondingly warmer or colder night sides.  The full dynamical simulations have been run again with these parameters.  The resulting light curves, in \\emph{Spitzer} bands, are shown in \\mbox{Figures~\\ref{figure:t3}}a, and c.  For the ``cold'' (-300 K) model CH$_4$ is the dominant carbon carrier on the night hemisphere, and CO on the day hemisphere, leading to large flux variation in excess of that found for our nominal CS06 simulation, as shown in \\mbox{Figure~\\ref{figure:t3}}.  For the ``hot'' (+300 K) model, on both the day and night hemispheres, CO is the dominant carbon carrier.  Methane absorption features are very weak on the night side, leading to flux variations in every band no larger than a factor of 2.7 from peak to trough.  This model is somewhat similar to our earlier Case 2, but at warmer temperatures, as at all phases CO is the dominant carbon carrier.  Fluxes are everywhere greater in the hot model than the nominal model, and everywhere less in the cold model than in the nominal model.  This is is due to the differing atmospheric temperatures.  The cold model best fits the \\citet{Deming05b} datum at 24 $\\mu$m.  In addition, the phase of maximum and minimum flux in a given band can change by up to $\\sim$10 orbital degrees between these simulations, due both to differing chemistry and atmospheric dynamics.  As discussed in \\citet{CS05} and CS06, while these models all possess similarly strong east-to-west jets, the dynamical atmospheres differ slightly in detail.  CS06 also examined non-equilibrium CH$_4$/CO chemistry for these models.  For the ``cold'' (-300 K) model, non-equilibrium chemistry was important, and the CH$_4$/CO ratio at $P < 1$ bar became homogenized at 0.20 around the entire planet.  This ratio was 0.014 for the nominal Case 1 decribed earlier.  For the cold model equilibrium chemistry would predict a night side dominated by CH$_4$ and a day side dominated by CO.  For the ``hot'' model (+300 K) both equilibrium and non-equilibrium chemistry predicts that CO is the dominant carbon carrier on both the day and side hemispheres.  \\mbox{Figure~\\ref{figure:t3}}d shows our computed flux ratios for  non-equilibrium chemistry for the cold case.  As was shown previously for our nominal case, flux variation is smaller with non-equilibrium chemistry, because the mixing ratios of CH$_4$, CO, and H$_2$O are the same on the night and day hemispheres.  Again, because of the nearly isothermal temperature structure of the day-side, chemical abundances have little effect on the day-side planet-to-star flux ratios, which are nearly equal for these two chemistry cases.  These additional cases further highlight the importance of CH$_4$/CO chemistry in the computation of infrared light curves.  Infrared fluxes as a function of orbital phase are sensitive to the temperatures of the hemisphere facing the observer, as well as the abundances of important molecular absorbers.  For planetary \\emph{P--T} profiles that cross important CH$_4$/CO chemical boundaries, as most hot Jupiters surely do, whether or not these species are found in their equilibrium mixing ratios has a major impact on resulting infrared flux in \\emph{Spitzer} bands, especially on the night side.  \\subsection{The Future \\& Conclusions} From the size of the error bar from the \\citet{Deming05b} 24 $\\mu$m observations, it is clear that, could this instrument stability be sustained over the course of tens of hours of observations, the change in flux over time presented here could be detected.  One-half of an orbital period for \\hd~is 42 hours.  If flux in the 8.0 $\\mu$m band for \\hd~is higher than models predict, as was the case for \\T~\\citep{Fortney05, Barman05}, then this would be an attractive band as well, as the error bars should be smaller.  What might one hope to see with sustained observations?  If the peak in infrared flux does occur only $\\sim$25-30 orbital degrees before secondary eclipse, this would be difficulty to discern.  However, the detection of any sort of ramp up in flux from the time of transit to secondary eclipse would give us important information on the day-night temperature contrast.  The recently discovered transiting planet \\he~\\citep{Bouchy05} would likely be an even more attractive target, as the planet-star flux ratios are likely twice as large \\citep{Fortney06, Deming06}, and the orbital period is $\\sim$40\\% shorter.  We predict secondary eclipse planet-to-star flux ratios for this system in \\citet{Fortney06}; our calculation at 16 $\\mu$m is a good match to the published observation of \\citet{Deming06}.  A clear prediction from our calculations here is that when one uses realistic non-equilibrium chemistry calculations, the change in planetary flux as a function of orbital phase is greatly reduced, relative to equilibrium chemistry calculations, because the atmosphere's composition is fixed.  For the \\emph{Spitzer} bands, for Case 0, the maximum flux variation is in the 3.6 $\\mu$m band, which varies by a factor of 15 from peak to trough.  The minimum variation is a factor of 2.2, in the 16 $\\mu$m band.  For Case 1, this variation drops significantly, and the maximum variation is a factor of 3.7, in the 5.8 $\\mu$m band, and the minimum is 2.0, again in the 16 $\\mu$m band.  We suggest that the 5.8 and 8.0 $\\mu$m bands may be the best \\emph{Spitzer} bands in which to search for flux variations as a function of orbital phase, as these bands combine a high planet-to-star flux ratio and the sensitivity and stability of the IRAC detectors.  If the day-side thermal emission of hot Jupiters is similar to a blackbody, as we find here, problems arise with the notion that the emission can be used to characterize the atmospheric chemistry from secondary eclipse observations.  Absorption features due to CH$_4$, CO, and H$_2$O would be nonexistent or extremely weak. Observations at other orbital phases would then take on additional importance.  Given the significant spectral differences between our model and radiative-equilibrium models, it is clear that more work in this area is certainly warranted.  We note that blackbody-like hot Jupiter emission can simultaneously explain all observations to date.  First is the secondary eclipse observation at 24 $\\mu$m by \\citet{Deming05b}, a clear detection.  Second is the very low flux ratio upper limit in L$^{\\prime}$ band by \\citet{Deming05c}.  Third is the 2.3 $\\mu$m relative flux observation of \\citet{Rich03b}, which most one-dimensional models cannot fit \\citep{Seager05}.  Fourth is the set of observations of \\T~by \\citet{Charb05}, who found an infrared spectral slope from 4.5 to 8 $\\mu$m that was redder than that found by most one-dimensional models \\citep{Fortney05, Barman05}.  Additional observations of these planets, especially in the \\emph{Spitzer} 3.6 $\\mu$m band, which catches much of the predicted 4 $\\mu$m flux peak, would strengthen or refute this argument and provide a critical test for the CS06 dynamical simulation.  \\citet{Deming05c} posited, perhaps with a wink, that a blackbody-like spectrum would be more consistent with observations to date than any published hot Jupiter model.  Due to the dynamically altered temperature structure of the atmosphere of \\hd, we find that this could be the reality. \\\\ \\\\ We acknowledge support from NASA Postdoctoral Program (NPP) and Spitzer Space Telescope fellowships (J.~J.~F.), NASA GSRP fellowship NGT5-50462 (C.~S.~C.), NSF grant AST-0307664 (A.~P.~S.), and NASA grants NAG2-6007 and NAG5-8919 (M.~S.~M.)."
    },
    "0608/astro-ph0608145_arXiv.txt": {
        "abstract": "The shocked wave created on the accretion disk after different physical phenomena (accretion flows with pressure gradients, star-disk interaction etc.) may be responsible observed Quasi Periodic Oscillations (QPOs) in $X-$ray binaries. We present the set of characteristics frequencies associated with accretion disk around the rotating and non-rotating black holes for one particle case. These persistent frequencies are results of the rotating pattern in an accretion disk. We compare the frequency's from two different numerical results for fluid  flow around the non-rotating black hole with one particle case. The numerical results are taken from our papers Refs.\\refcite{Donmez2} and \\refcite{Donmez3} using fully general relativistic hydrodynamical code with non-selfgravitating disk. While the first numerical result has a relativistic tori around the black hole, the second one includes one-armed spiral shock wave produced from star-disk interaction. Some physical modes presented in the QPOs can be excited in numerical simulation of relativistic tori and spiral waves on the accretion disk. The results of these different dynamical structures on the accretion disk responsible for QPOs are discussed in detail. ",
        "introduction": "\\label{INTRODUCTION}  Observed properties of the high frequency from $\\approx 1$ Hz to $\\approx 1$ kHz quasi periodic oscillations in the X-ray light curve of the black hole and the neutron star binaries have been studied by \\cite{vander1} and \\cite{McCalum1} and references therein.  The studies of high frequency variability in $X-$ray binaries provide a unique opportunity to explore the fundamental physics of space time and matter. The dynamical timescale on the order of several milliseconds is a timescale of the motion of matter through the region located in close to black hole. Thus, Theories of the physics of the inner accretion disk, the orbital motion of matter in relativistic gravity and the relativistic disk oscillations can be probed by high frequency signals from black hole binaries. The observation of high frequency properties have been motivated the theory of relativistic ``diskoseismology'' which  studies the hydrodynamical oscillation modes of geometrically thin accretion disk for one particle case \\cite{Wagoner1}, \\cite{Kato1}.  Different types of shock waves, standing two-armed spiral shock, moving one-armed shock and tori(\\cite{Donmez2},\\cite{Donmez3},\\cite{Donmez5}, \\cite{Donmez4}, \\cite{Donmez1}), with discrete and different range of frequencies can exist in such disk because the radial profile of general relativistic matters oscillation frequencies can create finite region where modes are trapped.  QPOs produced on the accretion disk in the black hole systems are thought to arise from physical processes. Depending on the where the oscillation reside, the various accretion models can be seen. These models are relevant to our understanding of various physical radii in the binary systems that are associated with high frequency variability according to the some theories. Some of these models accommodate movements of the inner accretion disk that are proposed to explained the  change in the frequency of high frequency QPOs.  However, radial pressure gradients themselves can permit the existence of discrete modes because they produce a disk, or torus, of  finite extent. At least in the case of black hole $X-$ray binaries, QPO s are only observed in the power law spectrum states in which the flow is not solely composed of a geometrically thin accretion disk. Therefore, it is believed that the radial pressure gradients may be responsible for trapping different frequencies modes on the accretion disk. Those reported have recently explored low frequency\\cite{Giannios1} and high frequency\\cite{Lee1},\\cite{Kluzniak1} that accretion disk may explain observed QPOs. Geometrically thick tori may form in stellar collapse, and modes of dense tori around black holes have also been suggested as a detectable source of gravitational waves\\cite{Zanotti1}  The first part of the paper provides the theoretical background on the physics of $X-$ray binary system frequency variability which is presented in accretion disks around the black hole for one particle system. This frequencies arise rotating patterns in the accretion disk and shock waves created on it and driven by gravity of black hole. Second part of paper compares the numerical simulation results computed from fully general relativistic hydrodynamics code with one particle case. We also report the different modes on numerical modeled accretion disk around the non-rotating black hole. ",
        "conclusions": "\\label{conclusion}  The oscillations of fluid confined by gravity and centrifugal forces can lead to rich behavior under a variety of astrophysical scenarios, QPOs in X ray binaries. It is plausible that the characteristic frequencies of the collective motion  manifests itself in X-ray luminosity variation. Under quite general considerations we have shown that simple, acoustic modes which have their origin in free particle oscillations are modified by the hydrodynamics and can couple to one another in non  linear fashion. The locking of excited modes could explain the fact that the observed QPO frequencies drift over considerable ranges while still arising from resonant interactions. We believe that their signature is present in the observations of accretion flows in strong gravitational fields, and will allow for the measurement of important parameters related to the black hole. The strong- field effects of general relativity and in particular metric properties of space-time around the black hole make the excitation of a 3:1 or 3:2 anharmonic epicyclic resonance, driven by orbital motion whose orbital frequency may be imprinted on the X-ray flux as a fairly prominent QPOs.  We have shown here that the tori and shock wave created on the accretion disk around the black hole produces oscillatory behavior. These types of structures constitute radial pressure gradients. So the pressure gradients themselves can permit the existence of discrete modes. As a consequence of these dynamics, X-rays can be emitted by black hole binary system. The results of these different dynamical structures on the accretion disk responsible for QPOs are discussed in detail. The observed frequency ratios are consistent with the computed numerical values. On the other hand, our numerical results exhibits different commensurate frequencies in $4:1$, $5:1$..., seen in Figs.\\ref{fig:NA:3:1} and \\ref{fig:Num:3:1}."
    },
    "0608/astro-ph0608690_arXiv.txt": {
        "abstract": "We use a high-resolution $N$-body simulation to study how the formation of cold dark matter (CDM) halos is affected by their environments, and how such environmental effects produce the age-dependence of halo clustering observed in recent $N$-body simulations. We estimate, for each halo selected at redshift $z=0$, an `initial' mass $M_{\\rm i}$ defined to be the mass enclosed by the largest sphere which contains the initial barycenter of the halo particles and within which the mean linear density is equal to the critical value for spherical collapse at $z=0$. For halos of a given final mass, $M_{\\rm h}$, the ratio $M_{\\rm i}/M_{\\rm h}$ has large scatter, and the scatter is larger for halos of lower final masses. Halos that form earlier on average have larger $M_{\\rm i}/M_{\\rm h}$, and so correspond to higher peaks in the initial density field than their final masses imply. Old halos are more strongly clustered than younger ones of the same mass because their initial masses are larger. The age-dependence of clustering for low-mass halos is entirely due to the difference in the initial/final mass ratio. Low-mass old halos are almost always located in the vicinity of big structures, and their old ages are largely due to the fact that their mass accretions are suppressed by the hot environments produced by the tidal fields of the larger structure. The age-dependence of clustering is weaker for more massive halos because the heating by large-scale tidal fields is less important. ",
        "introduction": "In the standard cold  dark matter (CDM) paradigm of structure formation, virialized CDM  halos are both the building blocks of the large-scale structure of the Universe and the hosts within which galaxies are supposed to form. During the last decade, the properties of CDM halos, such as their internal structure, formation histories and clustering properties,  have been studied in great detail using both $N$-body numerical simulations and analytical models. These studies have shown that the halo bias depends strongly on halo mass, in the sense that more massive halos are more strongly clustered (e.g., Mo \\& White 1996; Jing 1998; Seljak \\& Warren 2004). This mass dependence of  the halo bias has played a crucial role in understanding  the correlation function of both dark  matter and  galaxies,  via the halo model (e.g., Mo, Jing \\& B\\\"orner 1997; Ma \\& Fry 2000; Seljak 2000), the halo occupation models (e.g., Jing, Mo \\& B\\\"{o}rner 1998; Peacock \\& Smith 2000; Zheng et al. 2005), and the conditional luminosity function (e.g., Yang, Mo \\& van den Bosch 2003).  Recently Gao et al (2005), using a very large $N$-body simulation of a $\\Lambda$CDM cosmology, reexamined  the dependence of halo bias on the properties of dark matter halos. They found that,  for low-mass halos at redshift $z=0$ with $M_{\\rm h} \\ll M_\\star$ \\footnote{$M_\\star$ is the characteristic mass scale at which the {\\it RMS} of the linear density field is equal to $1.686$ at the present time. For the present simulation $M_\\star\\approx 1.0\\times 10^{13}h^{-1}{\\rm M}_\\odot$.}, the bias depends not only on the mass but also on the formation time of dark matter halos. For a given mass, halos that have earlier formation times are on average more strongly clustered in space. This age-dependence of halo clustering has been confirmed by a number of independent studies (e.g. Zhu et al. 2006; Harker et al. 2006). Since halo formation time is known to be correlated with halo concentration (e.g. Navarro, Frenk \\& White 1997; Jing 2000; Wechsler et al 2002; Zhao et al. 2003a,b; Lu et al. 2006), the age dependence of halo clustering also shows up as a halo-concentration dependence of halo clustering (Wechsler et al. 2005). For massive halos with $M_{\\rm h}\\gg M_\\star$, the age-dependence may be reversed, as shown in Wechsler et al. (2005) and Wetzel et al. (2006).  Although the mass-dependence of halo clustering is well understood in the excursion-set model based on spherical collapse in Gaussian initial density field (Mo \\& White 1996), the formation-time dependence (or age-dependence) is not. In the simplest excursion-set model from which the Press-Schechter formula is derived (Bond et al. 1991), the formation history of a halo is entirely determined by the local density field, independent of the mass distribution on larger scales. Therefore, the large-scale clustering of halos is independent of formation time in this model. An age-dependence of halo clustering is expected in the excursion-set model based on ellipsoidal collapse (Sheth, Mo \\& Tormen 2001), because here the collapse of a halo depends not only on the local density field but also on the tidal field generated by its large-scale environment. Unfortunately, this aspect of the model was not explored in any detail in Sheth et al. (2001). Instead, these authors used a mean relation between halo mass and the shape parameter of the large-scale tidal field to derive an average critical collapse overdensity that depends only on halo mass.  Since the properties of galaxies that form in a halo are expected to depend on the formation history of the halo, understanding the origin of the age-dependence of halo clustering is crucial for a better understanding about how galaxies are related to their large-scale environments. In this paper we use a high-resolution $N$-body simulation to study how the formation of dark matter halos is affected by their environments, and how such environmental effects can be used to understand the age-dependence of halo clustering. The outline of the paper is as follows. In Section \\ref{sec:simulation} we describe briefly the simulation to be used and how dark halos are identified. In Section \\ref{sec:age} we study the age-dependence of halo clustering in terms of halo properties in the initial density field. In Section \\ref{sec:env} we examine the environments of dark halos and how they may affect the formation histories of dark matter halos. We discuss and summarize our results in Section \\ref{sec:discussion}. ",
        "conclusions": "In this paper, we have used a high-resolution $N$-body simulation to study how the formation of CDM halos may be affected by their environments, and how such environmental effects may produce the age-dependence of halo clustering seen in recent large $N$-body simulations. We have shown that old low-mass halos are almost always associated with initial perturbations that correspond to higher peaks than their present masses imply. Consequently, these halos are strongly clustered in space despite of their low masses at the present time. Because of their associations with large density perturbations, old low-mass halos are almost always found in the vicinity of big structures that are dominated by one or more massive halos at the present time. Mass accretions into these low-mass halos are limited at late times by the tidal field of the larger neighboring structures. Such environmental effects are weaker for more massive halos, which explains why the age-dependence of halo clustering is weaker for more massive halos.  Our results demonstrate clearly that environmental effects play an important role during the formation of dark matter halos, especially of low-mass halos. Since the properties of galaxies that form in a halo are expected to depend on the formation history of the halo, such environmental effects may have important implications for galaxy formation in different environments. Galaxies that form in old, low-mass halos are presumably old and faint, and our results suggest that many such galaxies may be located in the vicinity of relatively massive galaxy systems. These galaxies likely  represent an extension of the satellite galaxies observed in galaxy systems such as the local group and other groups/clusters of galaxies. Observationally, it is known that satellite galaxies in galaxy groups are dynamically old systems (dwarf ellipticals and dwarf spheroids) that contain mainly old stellar populations. It is interesting to see whether the stellar populations and kinematics of low-mass galaxies in the neighborhoods of nearby galaxy groups have properties similar to those of the satellite population.  Theoretically, it is very interesting to see if the ellipsoidal collapse model that incorporates large-scale tidal field into the dynamics can indeed explain the environmental effects discussed here and the age-dependence of halo clustering observed in $N$-body simulations. We will come back to this problem in a future work."
    },
    "0608/astro-ph0608373_arXiv.txt": {
        "abstract": " ",
        "introduction": "The Clementine mission to survey the lunar surface, with its high resolution imagery, has re-opened the dormant field of crater counting on the Moon \\citep{pete94,moore96,mcewen97}.  \\citet{morota03} have used these high quality images to count young rayed craters and concluded that a factor of 1.5 enhancement exists in the current cratering record when comparing fitted crater densities precisely at the apex to those directly at the antapex. An excess of leading-hemisphere craters is expected on synchronously-rotating satellites in general; \\citet{hor84} give expressions for the expected asymmetry caused by the leading hemisphere of the satellite ``sweeping up'' more objects.  Based on such expressions and using their previous crater counts, \\citet{morota05} concluded that the measured leading enhancement implies an average impactor encounter speed with the Earth-Moon system of 12-16 km/s, consistent (though on the low end) with previous estimates \\citep{shoe83,chyba94}.  Other synchronously-rotating moons have been studied in great detail using the high-resolution images from Voyager.  By examining previously-obtained crater counts of the Galilean satellites, \\citet{zetal01} reconfirm the prediction of a leading hemisphere excess (which is quite large in the Jovian system due to high satellite orbital speeds) and are able to derive an analytical expression to describe the asymmetry that is consistent with results of \\citet{shoe82} and \\citet{hor84}.  In addition to this leading/trailing asymmetry, some authors believe that the Earth could act as a gravitational lens, focusing incoming projectiles onto the nearside of the Moon, causing an asymmetry between the near and far sides \\citep{turk62,wiesel71}.  This concept has been numerically investigated by \\citet{wood73,fev05} and deemed plausible, though an analytic calculation by \\citet{band73} showed the amplitude of the asymmetry depends on the Earth-Moon distance and the velocity of the incoming projectiles.   Our goal is to determine if the present-day flux of Near-Earth Objects (NEOs) would produce a spatial distribution of young craters on the lunar landscape that agrees with these studies, in terms of both leading/trailing and nearside/farside asymmetries.  Our approach involves many N-body simulations of projectiles drawn from the debiased NEO model of \\citet{bottke02}, which are then injected into the Earth-Moon system.  Since a realistic impactor distribution is not expected to be isotropic, one might find some latitudinal variation in the spatial distribution of deliveries on both the Earth and Moon.  This concept has been studied in the past \\citep{xmas64}, but more recently by \\citet{fev06} who concluded 60\\% and 30\\% latitudinal variations for the Moon and Earth respectively when comparing the impact density at the poles to that of the equator.  We seek to compare our direct N-body simulations to their semi-analytic approach as they also used the \\citet{bottke02} NEO model.  Asymmetries may also be present for the Earth and would be most reliably observed via fireball sightings and meteorite recoveries. Based on observations, many researchers concluded that more daytime fireballs occur in the afternoon hours, causing a morning/afternoon (AM/PM) asymmetry \\citep{wether68,xmas82}.  Several studies \\citep{xmas82,wether85,morglad98} have been done to understand this PM excess.  \\citet{xmas82} produced a PM fireball excess, but drew from a highly-selected group of orbital parameters for the meteoroid source population.  These orbits may not represent the true distribution of objects in near-Earth space. Thus, since we will obtain many simulated Earth impacts in our lunar study, an additional benefit is to reinvestigate the question of the AM/PM asymmetry.  The debiased NEO model we use as our source population should give a better approximation of the objects in near-Earth space than the parameters used by \\citet{xmas82}. Annual effects are also of interest as at different times of the year certain locations on Earth receive an enhancement in the impactor flux due to the location of the Earth's spin axis relative to the incoming flux direction \\citep{xmas82,renkno89}.  In Section~\\ref{sec-theory}, we give a brief overview on the theory of various asymmetries in the Earth-Moon system.  Section~\\ref{sec-MM} describes the model for our simulations as well as the methods of implementation.  Our findings begin in Section~\\ref{sec-ED}, where we first examine deliveries of projectiles to Earth in terms of the latitude distribution and then the AM/PM asymmetry.  Lunar results follow in Section~\\ref{sec-LB} where we examine the latitude distribution and the leading/trailing and nearside/farside hemispheric asymmetries.  Section~\\ref{sec-con} summarizes our key findings and discusses some possible implications. ",
        "conclusions": "\\label{sec-con}  We have used the debiased NEO model of \\citet{bottke02} to examine the bombardment of the Earth-Moon system in terms of various impact and crater asymmetries.  For Earth arrivals we find a $< 1$\\% variation in the ratio between the areal densities within 30\\degrees of the poles and within a 30\\degrees band centered on the equator.  The local time distribution of terrestrial impacts from NEOs is enhanced during the AM hours.  While this fall-time distribution corresponds well to recent radar data, it is in disagreement with the chondritic meteorite data and their derived pre-atmospheric orbital distributions.  This discrepancy thus reinforces the conclusion of \\citet{morglad98} that the large amount of decimeter-scale material being ejected from the main asteroid belt onto Earth-crossing orbits must be collisionally depleted before much of it can evolve to orbits with $a < 1.5$~AU.  A significant result is that we find the average impact speed onto the Moon to be $\\bar{v}_{imp} = 20$~km/s, with a non-negligible higher-speed tail (Fig.~\\ref{fig:velhist}). This combined with quantification of the non-uniform surface cratering has implications for both tracing crater fields back to the size distribution of the impactors and the absolute (or relative) dating of cratered surfaces. First, the higher impact speeds we find mean that lunar impact craters (at least in the post-mare era when we believe the NEO orbital distribution we are using is valid) have been produced by smaller impactors than previously calculated.  This roughly 10\\% higher average impact speed corresponds to lunar impactors which are 10\\% smaller on average than previously estimated; this small correction has ramifications for proposed matches between lunar crater size distributions and impactor populations (\\eg Strom \\etal 2005).   We find two different spatial asymmetries in current crater production due to NEOs.  As expected, due to its smaller mass, the Moon exhibits more latitudinal variation ($\\sim 10$\\%) in our simulations than the Earth. When comparing our simulation results to young rayed craters on the Moon, the surface density variation we predict is completely consistent with available data; we obtain a leading versus trailing asymmetry of $1.29 \\pm 0.01$ (GMAACA value), which corresponds to a 13\\% increase (decrease) in crater density at the apex (antapex) relative to the average.  These results indicate that using a single globally-averaged lunar crater production could give ages in error by up to 10\\% depending on the location of the studied region.  For example, post-mare studies of the Mare Orientale region would overestimate its age by $\\sim10$\\% due to its proximity to the apex, assuming that leading point and poles of the Moon have not changed over the last $\\sim$4~Gyr.  Similarly, the degree of bombardment on Mare Crisium (not far from the antapex) would be lower than the global average.  These effects will only be testable for crater fields with hundreds of counted craters so that the Poisson errors are small compared to the 10\\% variations we find; in most studies of ``young'' ($<4$ Gyr) lunar surfaces the crater statistics are poorer than this (\\eg St\\\"offler and Ryder 2001).  When the orbital distance of the Moon was decreased (as it was in the past because of its tidal evolution), the ratio between simulated craters on the leading hemisphere versus the trailing increased as expected due to the higher orbital speed of the satellite at lower semi-major axes. We find virtually no nearside/farside asymmetry until the Earth-Moon separation is less than 30 Earth-radii, which under currently-accepted lunar orbital evolution models dates to the time before 4~Gyr ago (at which point the current NEO orbital distribution may not be a good model for the impactors). Interior to 30 Earth-radii we find that the Earth serves as a mild shield, reducing nearside crater production by a few percent.   \\bigskip \\centerline{\\bf Acknowledgments} \\medskip  Computations were performed on the LeVerrier Beowulf cluster at UBC, funded by the Canadian Foundation for Innovation, The BC Knowledge Development fund, and the UBC Blusson fund.  BG and JG thank NSERC for research support.  \\singlespace  \\newpage \\newpage"
    },
    "0608/astro-ph0608659_arXiv.txt": {
        "abstract": "{Free-floating planetary-mass objects have masses below the deuterium burning mass limit at about 13 Jupiter masses, and have mostly been found in very young open clusters. Their origin and relationship to stars and brown dwarfs are still a mystery.} {The recent detection by direct imaging of three giant planets at wide separation (50--250 AU) from their primaries has raised the question about the ``true isolation'' of planetary-mass objects in clusters. Our goal was to test the possibility that some free-floating planetary-mass object could in fact be part of wide planetary systems.} {We searched in the literature for stellar and brown-dwarf member candidates of the $\\sigma$ Orionis cluster ($\\sim$ 3\\,Ma, $\\sim$ 360\\,pc) at small angular separations from published candidate planetary-mass objects. We found one candidate planetary system, SE 70,  composed of an X-ray source and a planetary-mass object, namely S\\,Ori 68, separated by only 4.6 arcsec. In order to assess the cluster membership of the X-ray source, we obtained mid-resolution optical spectroscopy using ISIS on the William Herschel Telescope. We also compiled additional data on the target from available astronomical catalogues.} {We have found that SE 70 follows the spectrophotometric sequence of the cluster and displays spectroscopic features of youth, such as lithium in absorption and chromospheric H$\\alpha$ emission. The radial velocity is consistent with cluster membership. Hence, SE 70 is very probably a $\\sigma$ Orionis cluster member. The projected physical separation between SE 70 and S\\,Ori 68 is 1\\,700$\\pm$300\\,AU at the distance of the cluster. If the common proper motion is confirmed in the near future, the system would be composed of an M5--6 brown dwarf with an estimated mass of $\\sim$ 45\\,M$_{\\rm Jup}$ and an L5$\\pm$2 giant planet with an estimated mass of $\\sim$ 5\\,M$_{\\rm Jup}$. It would be the widest and one of the lowest-mass planetary systems known so far.} {} ",
        "introduction": "Planetary-mass objects (PMOs) were {\\em directly} detected for the first time in the $\\sigma$ Orionis cluster and in other very young star-forming regions (Zapatero Osorio et al.\\ 2000; Lucas \\& Roche 2000; Najita et al.\\ 2000). Recent discoveries suggest that PMOs are abundant in clusters and include the detection of objects with masses potentially as low as $\\sim$ 3 Jupiter masses (M$_{\\rm Jup}$) (Zapatero Osorio et al.\\ 2002c; Lucas et al.\\ 2005). Furthermore, observations suggest that many PMOs are surrounded by discs (Luhman et al.\\ 2005). Although cluster PMOs have masses below the deuterium-burning mass limit ($\\sim$ 13\\,M$_{\\rm Jup}$; Chabrier \\& Baraffe 2000), they have not achieved the ``grade'' of exoplanet because they appear to float freely in the intracluster medium, far from the gravitational link to other star cluster members. The direct detection of planetary-mass candidates around a brown dwarf of the TW Hydrae association (Chauvin et al.\\ 2004 -- at a projected physical separation of 55\\,AU) and around the stars \\object{AB Pictoris} and \\object{GQ Lupi} (Chauvin et al.\\ 2005 -- 250\\,AU; Neuh\\\"auser et al.\\ 2005 -- 100\\,AU) presents the opportunity to photometrically and spectroscopically characterize exoplanets at orbital separations much larger than those indirectly detected with radial velocity surveys (the radial velocity exoplanet candidate with the largest semi-major axis is 55 Cnc d, with $a$ = 5.3$\\pm$0.9\\,AU; Marcy et al. 2002). The discussion on how PMOs in wide orbits can be formed has been fed by the recent discovery of the system \\object{Oph 162225--240515} AB (240\\,AU) in the 1\\,Ma-old Ophiuchus region, which may be a very low-mass brown dwarf and a PMO pair (Jayawardhana \\& Ivanov 2006) or a low-mass brown-dwarf binary (Close et al. submitted).  Here, we present the first results of a survey for wide planetary systems in the $\\sigma$ Orionis cluster. We have searched for candidate low-mass stellar and brown-dwarf members in close proximity to known PMOs. % Our knowledge of the stellar and brown-dwarf populations of the cluster has dramatically increased in the last few years (B\\'ejar et al.\\ 1999, 2001, 2004 -- mostly focused on very low mass stars and brown dwarfs; Scholz \\& Eisl\\\"offel 2004; Caballero et al.\\ 2004; Sherry et al.\\ 2004; Kenyon et al.\\ 2005; Burningham et al.\\ 2005; Hern\\'andez et al.\\ 2005). Prior to these studies, only massive and/or notably active stars were known in the cluster (e.g.\\ Haro \\& Moreno 1953; Warren \\& Hesser 1978; Wiramihardja et al.\\ 1989; Brown et al.\\ 1994; Wolk et al.\\ 1996). This cluster has become a favourite hunting ground for brown dwarfs and PMOs because it is nearby (360$^{+70}_{-60}$ pc), young (3$^{+5}_{-2}$ Ma) and free of extinction (e.g.\\ Caballero 2005; Jeffries et al.\\ 2006).   \\subsection{SE 70 and S\\,Ori 68}  The search for cluster candidate members at projected separations less than 15 arcsec (about 5\\,000\\,AU at the cluster distance) from isolated PMO candidates of $\\sigma$ Orionis had been unfruitful until the recent spatial and spectral analysis of a full {\\em XMM-Newton}/EPIC field in the cluster performed by Franciosini, Pallavicini \\& Sanz-Forcada (2006). In that work, out of the 175 X-ray sources detected, 88 were identified with cluster member candidates extracted from the literature. In two cases, the identification was ambiguous since there were two cluster member candidates falling inside their identification radii. One of these systems, associated with the X-ray source NX 67, was formed by an intermediate M-type brown-dwarf candidate proposed by Scholz \\& Eisl\\\"offel (2004), named SE 70, and by a PMO candidate, S\\,Ori 68. NX 67 underwent a flare during the observation, with an increase in the count rate over the quiescent level by a factor of $\\sim$ 4 and a total duration of about 3\\,h. Although Franciosini et al.\\ (2006) did not exclude the presence of a very weak contribution of S\\,Ori 68 to the total X-ray flux, they concluded that the bulk of it was associated with SE 70.  The planetary-mass object candidate S\\,Ori 68 ($\\sim$ 5\\,M$_{\\rm Jup}$), previously classified as ``isolated'', is about 4.6 arcsec southeast of SE 70 ($\\theta \\approx$ 140\\,deg). At the distance of the $\\sigma$ Orionis cluster, the projected physical separation is 1\\,700$\\pm$300\\,AU. S\\,Ori 68 was firstly identified by Zapatero Osorio et al.\\ (2000) as a free-floating planet candidate from an $IJ$ survey with follow-up in the $K$ band. Barrado y Navascu\\'es et al.\\ (2001) classified it as an L5.0$\\pm$2.0-type object from an optical spectrum taken with FORS at the Very Large Telescope. The low resolution and the insufficient signal-to-noise ratio of the data did not allow them to detect spectroscopic features of extreme youth (i.e.\\ Li {\\sc i} $\\lambda$6707.8 in absorption, strong asymmetric H$\\alpha$ emission, anomalous broadening of atomic absorption lines due to low gravity). Although the object follows the spectrophotometric cluster sequence, more data are needed to confirm its membership in the $\\sigma$ Orionis cluster.  Nevertheless, the proximity between SE 70 and S\\,Ori 68 motivated us to characterize in more detail the physical properties of the primary. We have taken intermediate-resolution optical spectroscopy and have compiled astrometric and photometric data of SE 70 to test more rigorously the hypothesis of membership in the $\\sigma$ Orionis cluster. ",
        "conclusions": "Except for the unreliable astrometric measurement, all the studied astrometric parameters indicate that SE 70 is a bona fide non-accreting $\\sigma$ Orionis cluster member: it displays X-ray emission, chromospheric H$\\alpha$ emission, Li {\\sc i} in absorption and radial velocity similar to those of the cluster and of the nebulosity, and lies on the spectrophotometric cluster sequence. The derived spectral type is on the borderline between very low-mass stars and brown dwarfs in $\\sigma$ Orionis. However, the $J$-band magnitude of SE 70, which marks the peak in the spectral energy distribution, is about 1.0\\,mag fainter than those of the most masive brown dwarfs in the cluster (at the hydrogen-burning mass limit, $\\sim$ 75\\,M$_{\\rm Jup}$; B\\'ejar et al.\\ 2001; Caballero et al.\\ 2004). Following our procedure for mass determination in the substellar domain, detailed in Caballero (2006)\\footnote{The mass is derived from the comparison between theoretical luminosities of Chabrier et al.\\ (2000) and the luminosities of the objects. The latter are computed from the $J$-band absolute magnitude and an $I-J$ colour-dependent bolometric correction.}, we estimate SE 70 to be a brown dwarf with a mass of 45\\,M$_{\\rm Jup}$. Hence, S\\,Ori 68 and SE 70 may form a system whose masses resemble those of the Chauvin et al.\\ (2004) \\object{2MASS J1207334--393254} system (about 45 + 5\\,M$_{\\rm Jup}$ in the case of SE 70 + S\\,Ori 68 versus about 25 + 5\\,M$_{\\rm Jup}$ in the case of 2M1207--39 + 2M1207--39b). However, the projected physical separation of the hypothetical $\\sigma$ Orionis planetary system is about 30 times larger than that of 2M1207--39, and about seven times larger than that of AB Pic and Oph 162225--240515. Due to the low mass and the wide separation of the SE 70 + S\\,Ori 68 system, it has the lowest gravitational binding energy among the considered planetary systems. Given the relatively high spatial density of cluster members in $\\sigma$ Orionis and the weak bounding between SE 70 and S\\,Ori 68, it seems dubious that the system can survive the tidal field within the cluster. The existence of a binary system with these characteristics is evidently a challenge for present ultra low-mass star forming scenarios, which predict tighter and more massive binary/planetary systems.  \\begin{figure} \\centering \\includegraphics[width=0.49\\textwidth,angle=0]{se70_f04.eps} \\caption{Angular separation to the nearest neighbour vs.\\ separation to the centre of the $\\sigma$ Orionis cluster. Small dots represent the star and massive brown-dwarf photometric member candidates in the study by Caballero (in prep.), while the large filled stars represent the SE 70 + S\\,Ori 68 system. Low-mass brown-dwarf and IPMO cluster member candidates are not included. Horizontal lines denote the 90 (dotted) and 95\\,\\% (solid) confidence limits below which the probabilities for a pair at 5 to 10\\,arcmin to $\\sigma$ Ori A to be a visual alineation are less than 10 and 5\\,\\%, respectively. There are only two systems below the 95\\,\\% confidence limit: [W96] 4771--899 AB + C and SE 70 + S\\,Ori 68. The results from 0 to 30\\,arcmin to the cluster centre and the description of the spatial distribution study will be given in Caballero (in prep.). } \\label{fe04} \\end{figure}  There is the possibility of the pair being a visual projection of two unrelated cluster members at slightly different heliocentric distances and hence not being a truly binary system. This is in contradiction with the statistical analysis performed by Caballero (in prep.), who has studied the spatial distribution of {\\em all} the stellar population and the highest-mass domain of the brown-dwarf population in the $\\sigma$ Orionis cluster. He has shown that 95\\,\\% of the studied cluster member candidates at angular separations to $\\sigma$ Ori A between 5 and 10 arcmin have their nearest neighbourhood at more than 10.9\\,arcsec (the SE 70 + S\\,Ori 68 system is at about 8\\,arcmin to the cluster centre). Among the more than 400 cluster member candidates studied  from a 2MASS + DENIS correlation, none has a nearest neighbour closer than 7.1\\,arcsec, while the separation between SE 70 and S\\,Ori 68 is only 4.6\\,arcsec. Figure \\ref{fe04} illustrates this discussion, and indicates that the SE 70 + S\\,Ori 68 pair has less than about 1\\,\\% probability to be a chance visual alignment. There are stellar pairs in the $\\sigma$ Orionis cluster with smaller separations between components, but they are spectroscopic binaries (e.g.\\ \\object{OriNTT 429}; Lee et al.\\ 1994) or very tight binaries unresolved by 2MASS or DENIS (e.g.\\ $\\sigma$ Ori A and \\object{$\\sigma$ Ori B} -- 0.250\\,arcsec; $\\sigma$ Ori AB and \\object{$\\sigma$ Ori IRS1} -- 3.32$\\pm$0.15\\,arcsec; \\object{HD 37525 A} and \\object{HD 37525 B} -- 0.47$\\pm$0.04\\,arcsec; \\object{[W96] 4771--899 A} and \\object{[W96] 4771--899 B} -- 0.40$\\pm$0.08\\,arcsec; Caballero 2005)\\footnote{The tight pair [W96] 4771--899 AB and \\object{[W96] 4771--899 C} (aka \\object{S\\,Ori J053847.7--022711}, at $\\rho$ = 7.63$\\pm$0.10\\,arcsec) is in fact a triple stellar system. There is an additional binary system candidate formed by a very low-mass star, S\\,Ori J0539268--026614, and a faint brown dwarf at 4.4\\,arcsec ($J$ = 16.21$\\pm$0.09\\,mag, not found in 2MASS or DENIS catalogues). Both of them show Li {\\sc i} in absorption (Zapatero Osorio, B\\'ejar et al. in prep.).}.  It is not the aim of this paper to prove beyond doubt the cluster membership of the planetary-mass candidate S\\,Ori 68. In order to do that, high-quality mid-resolution near-infrared spectra of S\\,Ori 68 should be taken to detect low-gravity spectral features in the absorption atomic lines of alkalis, such as sodium or potassium. These features have been already found in several spectra of young brown dwarfs (e.g.\\ Gorlova et al.\\ 2003). To date, this study may only be accomplished, and with a time-consuming effort, with instruments such as LRIS at the Keck Observatory. Given the extraordinary faintness of S\\,Ori 68 in the red optical range ($I \\sim$ 24\\,mag, $R >$ 26--27\\,mag; S\\,Ori 68 has the same $I$- and $J$-band magnitudes as \\object{S\\,Ori 66}, an intermediate-L planetary-mass object in $\\sigma$ Orionis with H$\\alpha$ in moderate emission), the detection of lithium in absorption in an optical spectrum is practically out of the question with present instrumentation.  Further ultra-accurate proper motion studies are needed to verify that both S\\,Ori 68 and SE 70 form a common proper motion system. The Ori OB1b Association is in the solar antapex, moving in the opposite direction with respect to the Sun. Hence, the tangential velocities of the stellar members of the association are very low,  less than 5\\,mas\\,a$^{-1}$ according to de Zeeuw et al.\\ (1999). Confirmation of the common proper motion would require an accuracy of a few hundred milliarsec. Such accurate measurements on very faint objects only be will achievable with technology of the near future, such as the Large Binocular Telescope and the LINC-NIRVANA instrument. It could be also done with the {\\em Hubble} Space Telescope if  a large enough time baseline is used. However, it could not be done with the {\\em GAIA} satellite, because of faintness of targets in the optical.  Radial velocity measurements of the pair obtained with current or future near-infrared high-resolution spectrographs attached to $\\ge$10\\,m-class telescopes, like NIRSPEC/Keck or NAHUAL/Gran Telescopio Canarias (Mart\\'{\\i}n et al. 2005), might be able to constrain the binding status of the SE 70 $+$ S\\,Ori 68 system. These measurements must be more precise than 1\\,km\\,s$^{-1}$ to disentangle the common proper motion of the hypothetical binary system and the velocity dispersion of the $\\sigma$ Orionis cluster (that is of the order of 2.4\\,km\\,s$^{-1}$ -- Caballero 2005). To illustrate the capabilities of NIRSPEC/Keck, Mart\\'{\\i}n et al. (2006) reached a precision of 360\\,m\\,s$^{-1}$ in the radial velocity determination of the M9V-type brown dwarf \\object{LP 944--20}."
    },
    "0608/hep-ph0608193_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:introduction}  \\mysect{Intro to SUSY}  Supersymmetry (SUSY), a symmetry between fermions and bosons, allows for the construction of models that link a wide range of physical phenomena. While initially proposed on aesthetic grounds, SUSY also allows for cancellation of the quadratic Higgs divergence and, hence, opens the window for consistent model building up to the grand unification (GUT) or Planck scales.  The extension of supersymmetry to a local gauge theory, supergravity \\cite{sugra}, led to the development of GUT models \\cite{gut,gut_two} giving a description of physics from the GUT scale down to the electroweak scale, in addition to incorporating the successes of the Standard Model (SM).  LEP data confirmed the validity of the idea of SUSY grand unification.  Additionally,  SUSY models with R-parity conservation automatically give rise to a cold dark matter candidate, the lightest neutralino, \\none.  Because these models are consistent to very high energies, they provide descriptions of the early universe that deeply link particle physics and cosmology.  Detailed theoretical calculations \\cite{dark_constraint} confirm that GUT models can account for the experimentally observed dark matter abundance \\cite{wmap}.  \\mysect{Intro to mSUGRA}  There is good reason to believe that SUSY can be discovered at the Large Hadron Collider (LHC) and that it may be possible to determine its parameters with enough precision to distinguish among SUSY models that correctly predict the dark matter content of the universe.  While the ideal is to study the prospects of measuring SUSY in as general a way as possible, we consider mSUGRA \\cite{gut} here, a commonly studied model, for concreteness. In mSUGRA, four parameters and one sign determine all the masses and couplings: $m_0$, the universal scalar soft breaking mass at $M_{GUT}$;  $m_{1/2}$, the universal gaugino mass at $M_{GUT}$; $A_{0}$, the universal cubic soft breaking parameter at $M_{GUT}$; \\tanbeta = $\\langle H_{1} \\rangle / \\langle H_{2} \\rangle$, where $\\langle H_{1(2)} \\rangle $ is the Higgs vacuum expectation value which gives rise to the up (down) quark masses; and the sign of $\\mu$, where $\\mu$ is the coefficient of the quadratic term of the superpotential, $W = \\mu H_{1} H_{2}$.  Dark matter abundance and accelerator data constrain the parameter space to only three regions \\cite{dark_constraint}: (1) the \\mbox{\\stau - \\none} co-annihilation region\\footnote{Note that the notation  \\stau\\ here refers to the lighter of the two supersymmetric tau mass eigenstates.} \\cite{coannihilation} where the \\none\\ is the LSP and \\mbox{\\dm\\ $\\equiv$ $M_{\\stau}-M_{\\none}$}  is small (\\dm\\ $\\sim$ 5-15 GeV),  (2) the focus point region where the \\none\\ has a large Higgsino component, and (3) the scalar Higgs annihilation funnel region.  A  bulk region is also allowed by dark matter constraints, but it is largely excluded by other constraints.  \\mysect{why Co-annihilation?}  A collection of other constraints suggests that, of these regions, the co-annihilation region is particularly important.  They include the following: the light Higgs and lightest chargino mass bounds from LEP \\cite{bound}, the \\mbox{$b\\rightarrow s\\gamma$} branching ratio bound \\cite{rare_decay}, and the 3.3$\\sigma$ deviation from the SM expectation of the anomalous muon magnetic moment from the muon $g-2$ collaboration \\cite{muon}.  If the anomalous muon magnetic moment deviation is proven, mostly the co-annihilation region will be allowed. Dark matter production and annihilation rates in the early universe are believed to control the dark matter relic density we observe today; however, co-annihilation can provide an additional mechanism for the reduction of dark matter abundances in the early universe. In particular,  the \\stau\\ and \\none\\ can co-annihilate sufficiently into SM particles only if \\dm\\ is small.  Since, in principle, $M_{\\stau}$ and $M_{\\none}$ can separately range from 100~GeV into the TeV domain, if this striking smallness of \\dm\\ were found at the LHC, it would be a strong indirect indication that the \\none\\ is the astronomical dark matter particle.  We also note that many other SUGRA models have a co-annihilation region (since co-annihilation can occur with non-universal gaugino masses), so investigating this region has greater applicability than just mSUGRA (though, constraints do not necessarily exist in other SUGRA models to select out the co-annihilation region).   \\mysect{Previous work}  The co-annihilation region is difficult to analyze due to the production of low energy $\\tau$'s from \\mbox{$\\stau\\rightarrow\\tau\\none$} decays.  Some previous work has been done to determine how to measure masses of SUSY particles in the co-annihilation region. One Monte Carlo study  investigated the co-annihilation region by fitting various final state mass distributions in order to determine the SUSY masses \\cite{endpoints}, but this study did not consider any backgrounds.  A second analysis of SUGRA models with large $\\tanbeta$ and large \\dm\\ ($\\sim$50~GeV)  showed that high \\Et\\ hadronic $\\tau$'s can be used to reconstruct final states at the LHC \\cite{large_dm}.  It has also been shown that a small \\dm\\ can be measured at the ILC to 10\\% \\cite{ilc}, but the LHC will be available long before the ILC.  An analysis similar to this one has shown that it is possible to measure a small \\dm\\ at the LHC to 18\\% with a luminosity of 10~fb$^{-1}$, but it was assumed there that an independent gluino mass measurement would exist at the 5\\% level \\cite{two_tau}.  While this may be a good assumption, at high $\\tanbeta$ and small \\dm, low energy final state $\\tau$'s can complicate the interpretation of the typical \\mg\\ measurment and may affect the measurement uncertainty as well as the reliability \\cite{gluino}.  \\mysect{Roadmap to our paper}  In this letter, we describe new techniques and the prospects of measuring \\dm\\ and \\mg\\  simultaneously at the LHC, helping  to establish whether we reside in the co-annihilation region. Following \\cite{ilc,two_tau}, we consider an mSUGRA scenario in the co-annihilation region with $\\tanbeta$~=~40~GeV, $A_{0}$~=~0, and \\mbox{$\\mu>0$}, but we allow $m_{0}$ and $m_{1/2}$ to vary such that \\dm\\ $<$ 20~GeV.  In Sec. 2, we describe the requirements to select inclusive 3$\\tau$+jet+$\\mett$ events and discuss the differences between this 3$\\tau$ analysis and a similar  analysis done with two $\\tau$'s.  In particular, we motivate our selection criteria and  describe the mass and counting observables.  In this section, it will become clear that a primary experimental requirement for this analysis to be realized in practice is the efficiency to identify taus with \\Et\\ as low as 20~GeV.  In Sec. 3, we discuss a method to use both observables to simultaneously determine both \\dm\\ and \\mg.  Conclusions are given in Sec. 4. ",
        "conclusions": "\\label{conclusions}  \\mysect{conclusions}  We have studied the prospects of simultaneously measuring \\dm\\ and \\mg\\ at the LHC in the co-annihilation region using the inclusive 3$\\tau+ \\mathrm{jet}+\\mett$ final state for an mSUGRA scenario. Because this channel is nearly background free after selection cuts, for \\dm\\ $>$ 5 GeV we are able to form two separate observables: the number of signal events (a \\ntwo\\ cross-section measurement equivalent) and \\mtt.  Both vary with \\dm\\ and \\mg\\ allowing the simultaneous measurement of both.  This is particularly important if we are in the co-annihilation region as it is unclear if there will be a high quality \\mg\\ measurement available since current methods for measuring \\mg\\ will likely have to be altered to accommodate low \\Et\\ $\\tau$'s in the final state.  We find that with 30~fb$^{-1}$  we can measure \\dm\\ to 15\\% and \\mg\\ to 6\\% for the example point of \\dm=9~GeV and \\mg=850~GeV.  While our sensitivity to measuring \\dm\\ at the LHC is not as good as that expected at the ILC, it is quite comparable and should be available much sooner.  Further, a 15\\% measurement of \\dm\\ would generally be sufficient to determine if the signal is consistent with co-annihilation, and therefore, with the \\none\\ being the dark matter particle. Since no gluino mass measurement is possible at the ILC unless a very high energy option is available, this may be the only correct measurement of \\mg.  We also note that we have made no attempt to optimize these results, indicating that with actual data from the detector, we will likely be able to optimize our cuts, leading to a more precise measurement or a lower luminosity needed for the same sensitivity.  We have confirmed that efficient $\\tau$ identification down to an \\Et\\ of 20~GeV is crucial for this analysis as in the 2$\\tau$ analysis \\cite{two_tau}.  Further, we expect that this analysis and the 2$\\tau$ analysis could be used to complement each other in the establishment of a co-annihilation signal at the LHC, and perhaps be combined to produce a more accurate  measurement.  Finally, an analysis of this type can by applied to other SUGRA models provided they have a co-annihilation region and do not suppress the production of the \\ntwo\\ particles."
    },
    "0608/astro-ph0608415_arXiv.txt": {
        "abstract": "We present general relativistic solutions for  self-similar  spherical perturbations in an expanding cosmological background of cold pressure-less gas. We focus on solutions having  shock discontinuities propagating  in the surrounding  cold  gas. The pressure, $p$, and energy-density, $\\mu$, in the shock-heated matter are assumed to obey $p=w\\mu$, where $w$ is a positive constant. Consistent solutions are found for shocks propagating from the symmetry center of a region of a positive density excess over the background. In these solutions, shocks  exist outside the radius marking the event horizon of the black hole which would be present in a shock-less collapse. For large jumps in the energy-density at the shock, a black hole is avoided altogether and the solutions are regular at the center. The shock-heated gas does not contain any sonic points, provided the motion of the cold gas ahead of the shock deviates significantly from the Hubble flow. For shocks propagating in the uniform background, sonic points always appear for small jumps in the energy-density. We also discuss self-similar solutions without shocks in fluids with  $w<-1/3$. ",
        "introduction": "Detailed studies of generic self-gravitating  systems are only possible with the aid of complex numerical methods. Part of the complication arises from the long range nature of gravitational interactions. Thus, even equilibrium states, when they exist, do not in general obey standard  statistical mechanics. This  is correct for both Newtonian gravity as well as general relativity (GR), but the latter adds its peculiar technical and conceptual challenges. Analytic or semi-analytic treatment of  systems possessing special symmetries are, therefore, of great interest. Gravity is the dominant force dictating the growth of structure in the universe. Thus, self-gravitating  perturbations in an expanding cosmological has traditionally gained a great deal of attention (cf. Peebles 1990). Here we present a class of solutions which describe the evolution of spherical perturbations in an expanding Einstein-de Sitter cosmological background. We focus on the evolution of self-similar  perturbations in GR. Self-similarity occurs when  dimensionless physical variables can be expressed as functions of some combination of the time and space coordinates. Self-similar systems offer great simplification since they are described by ordinary differential equations. Newtonian treatment of cosmological self-similar perturbations  have been presented  in great detail previously in the literature  (e.g. Fillmore \\& Goldreich 1984; Bertschinger 1985; Chuzhoy \\& Nusser 2000). Nevertheless, it is instructive to review the general features of these solutions. In cosmology, we assume that the density fluctuations are extremely small at time, $t_i$, close to the  the Big Bang singularity. An isolated small density perturbation  having  the form  $\\delta \\mu \\sim r^{-s}$ ($r$ is the distance from the symmetry center) at $t_i$, develops a self-similar behavior  at sufficiently late times, $t\\gg t_i$\\footnote{An additional assumption which is often made is that the initial expansion rate of matter at  $r$ is equal to that of the unperturbed background. This assumption is, however, unnecessary  as long as the deviations from the background expansion rate are not great.}. Consider first a positive density perturbation in  a universe made solely of dust (i.e. collision-less matter). As time goes-by a shell of matter begins to expand at a rate lower than the background until it reaches a maximum distance from the center. The shell then falls inward and ends up oscillating around the center together with shells which  have reached it at earlier times. According to  this picture, the radius, $r_1$, of the shell which is at maximum expansion at time $t$ plays an important role. For $r>r_1$,  matter is still  expanding, while at $r\\ll r_1$ shells of matter are oscillating in the central  potential well generated by their self-gravity. The self-similar variable can then be taken to be the ratio, $r/r_1(t)$. For a universe filled with a collisional fluid (gas), self-similarity develops provided the initial pressure is negligible and that the gas is able to shock and evolve adiabatically. In this case, an inner region of hot shocked gas develops. The  shock bounding this region propagates  away from  the center at a rate $\\dd r_1(t)/\\dd t$.   Self-similarity in GR is obtained for a more restricted class of initial perturbations than in Newtonian gravity. GR admits self-similar solutions only for specific values of the power law index $s$, while Newtonian dynamics does not impose such restrictions. Further, the pressure, $p$, must be related linearly to the energy-density, $\\mu$,  of matter (Cahil \\& Taub 1971). Carr \\& Hawking (1974) and Carr \\& Yahil (hereafter CY90)  have studied   self-similar perturbations in an expanding universe. Their treatment, however, did not include shocks which might form as a result of collapse or energy injection into the system. Here we complement the work of CY90 by exploring solutions which involve shock discontinuities. We assume that the isotropic cosmic background is made of cold pressure-less matter that is able to shock. We will  present solutions in which shocks propagate from a central point into the uniform cosmic background or into the outer regions of a perturbation with  positive excess of energy-density in the cold gas. Shock-heated gas  will be assumed to obey the equation of state $p=w\\mu$ ($w>0$). We also present  solutions with no shocks for $w<-1/3$, a case which has not been discussed at length in the literature before.   The paper is organized as follows. In \\S\\ref{sec:eom} we summarize  GR equations for self-similar perturbations in comoving coordinates. For a detailed  discussion of these equations we refer the reader to  CY90 and Carr \\& Coley (2000) (hereafter CC02). In \\S\\ref{sec:shock} we discuss the formation of shocks and the jump conditions relating the physical variables on both sides of the shock discontinuity. In \\S\\ref{sec:num} we present full numerical solutions. We conclude with a general discussion in \\S\\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  In the standard cosmological model, the universe on large scales is nearly uniform so that super-horizon perturbations  are well described by the general relativistic theory of linear perturbations (e.g. Bardeen 1980). However, in view of resent observational data, it seems that GR might be  needed for a full study of perturbations on small scales where strictly speaking the linearized equations are invalid. These data suggest that the expansion of the universe is accelerating (e.g. Spergel et. al. 2003). It is thought that this acceleration is driven by a {\\it dark-energy} component which, in its simplest incarnation, is described in terms of a fluid with pressure $p$ related to the energy density, $\\mu$,  by $p=w \\mu$ with $w\\approx -1$ at the present cosmic epoch. The continuity equation, $\\dot \\mu=-(1+w) \\mu$ implies that for $w=-1$ the density is constant with time and, therefore, by Lorentz invariance, the fluid must be homogeneous. More generally, dark energy is modeled in terms of a scalar field which, under certain conditions,  could be  represented as a fluid with $w$ varying with time. Such a fluid is not ruled out by observations. If this is the case,  inhomogeneities in the dark energy component  could  develop. Since dark energy is inherently a relativistic entity, GR therefore should  be invoked for a thorough treatment of small scale cosmological perturbations. Seen in this perspective self-similar solutions might be relevant for the description of structure formation in the universe. However, the solutions are quite limited in their applicability in a realistic model for the growth of structure in the presence of dark energy. Putting the subdominant baryonic matter aside, the problem of structure formation requires tracing the evolution of the gravitationally coupled system of collisionless dark matter and dark energy.   Our solutions do not take this coupling into account. It is unclear whether self-similar evolution of cosmological perturbations in a two fluid system is possible at all. Further, we assume a strict equation of state of the form $p=w\\mu$. This leads to an imaginary ``sound\" speed, $(\\delta p/\\delta \\mu)^{1/2}$ for $w<0$, while  a more realistic model for dark energy in terms of  scalar fields leads in general to a non-imaginary sound speed. Verertheless, because of the paucity of analytic solutions for time-dependent systems in GR,  a treatment of systems with restricted symmetry is a worthwhile effort. Analytic study of these systems could also  be  useful in testing and calibrating numerical techniques for solving the GR equations."
    },
    "0608/astro-ph0608309_arXiv.txt": {
        "abstract": "We present ground-based infrared $K$- (2--2.5 $\\mu$m) and $L$-band (2.8--4.1 $\\mu$m) spectroscopy, as well as interferometric observations at $\\lambda \\sim$ 3mm, for the individual merging components (A, B, and C) of the luminous infrared galaxy Arp 299. We investigate the presence and location of the putative buried active galactic nucleus (AGN) inferred from previous X-ray observations at E $>$ 10 keV. Our sub-arcsec-resolution infrared spectra clearly reveal that the putative buried AGN resides in the nucleus B1 (a subcomponent of B), based on a very low equivalent width of 3.3 $\\mu$m polycyclic aromatic hydrocarbon emission, a weak 2.3 $\\mu$m CO absorption feature, and a large time variation of the $K$- and $L$-band continuum fluxes. In component C, we find strong 3.1 $\\mu$m ice absorption at $L$ and weak 2.3 $\\mu$m CO absorption at $K$, as expected in a buried AGN; however, a centrally concentrated young super star cluster is an alternative possibility because of the modest infrared luminosity and non-galaxy-nucleus nature of this component. The infrared $K$- and $L$-band spectra of the infrared brightest nucleus, A, are typical of a normal starburst with no explicit AGN signatures. Our interferometric observations simultaneously obtain HCN (J = 1--0) and HCO$^{+}$ (J = 1--0) emission lines with $\\sim$4 arcsec resolution, and we find the HCN to HCO$^{+}$ brightness-temperature ratios to be as low as those found in starburst nuclei in all the major merging components of Arp 299. The low ratio even in the AGN-hosting nucleus B may be due to the presence of a large amount of high-density molecular gas whose chemistry is dominated by coexisting starbursts and/or shocks, rather than by the central strong X-ray-emitting AGN. ",
        "introduction": "Luminous infrared galaxies (LIRGs; L$_{\\rm IR}$ $\\gtrsim$ 10$^{11}$L$_{\\odot}$), discovered by the {\\it IRAS} all sky survey, radiate the bulk of their luminosities as dust emission in the infrared \\citep{sam96}. Powerful energy sources, starbursts and/or active galactic nuclei (AGNs), must therefore be present, but hidden behind dust. Since these LIRGs dominate the cosmic infrared background emission, which is the sum of dust-obscured activity in the universe \\citep{dol06}, distinguishing their energy sources is closely related to clarifying the connections between obscured AGN and starbursts.  The difficulty in disentangling an AGN from a starburst as the primary energy source of a LIRG comes from the fact that highly concentrated molecular gas and dust in a LIRG's nucleus \\citep{sam96} can easily {\\it bury} an AGN in all directions. A luminous AGN surrounded by torus-shaped dust is distinguishable from a normal starburst galaxy relatively easily through optical and infrared high-resolution spectroscopy \\citep{vei87,gen98} because emission from the so-called narrow line regions (NLRs), which are well developed along the torus axis above a torus scale height and are photoionized by the central AGN's radiation, is visible from all directions \\citep{ant93}. However, the presence of a {\\it buried} AGN in a LIRG's core is difficult to unveil with this conventional type of spectroscopy because no significant NLRs are expected to develop.  Even if a luminous buried AGN is present in a LIRG's core, starbursts are likely to coexist, surrounding the central AGN. When such a starburst/buried-AGN composite LIRG is observed at wavelengths of large dust extinction, the observed flux is dominated by the less obscured starbursts, with a small contribution from the more highly obscured buried AGN. It is therefore essential to observe at wavelengths of low dust extinction if we are to properly investigate the presence of a buried AGN in a LIRG's core. Several useful methods for this purpose currently exist. The first is to observe in the hard X-ray band at E $>$ 2 keV. Since an AGN is intrinsically a much stronger X-ray emitter than a starburst \\citep{elv94,ran03}, the detection of strong but highly-absorbed X-ray emission can provide evidence of a luminous buried AGN in a LIRG's core. The second is millimeter interferometric searches for the chemical effects of the strong X-ray-emitting AGN on the surrounding interstellar medium. Around a buried AGN, X-ray dissociation regions (XDRs; Maloney et al. 1996) should develop, rather than the photodissociation regions (PDRs) usually found in starburst regions. Since the emission mechanism is different, XDRs and PDRs can be distinguished if they show different flux ratios of emission lines. The feasibility of carrying out this distinction has been demonstrated using HCN (J = 1--0) ($\\lambda$ = 3.3848 mm) and HCO$^{+}$ (J = 1--0) ($\\lambda$ = 3.3637 mm) emission lines in the millimeter wavelength range \\citep{koh05}. The third method is infrared spectroscopy. While a normal starburst should always show large-equivalent-width polycyclic aromatic hydrocarbon (PAH) emission features in the spectra at $\\lambda_{\\rm rest}$ = 3--25 $\\mu$m in the rest-frame, an AGN displays a PAH-free continuum \\citep{gen98,imd00}. Infrared $K$-band (2.0--2.5 $\\mu$m) spectroscopy can also be used to separate starburst and AGN emission, because a normal starburst displays CO absorption features at $\\lambda_{\\rm rest}$ = 2.3--2.4 $\\mu$m, whereas an AGN does not \\citep{iva00,ia04}. Since a buried AGN is generally more difficult to find with any method than the surrounding less obscured starbursts, using a combination of these potentially powerful approaches can increase the likelihood of detecting the elusive buried AGN and help to draw a quantitatively consistent picture about the energetic importance of the buried AGN.  Arp 299 ($z$ = 0.010) is a nearby well studied LIRG. It consists of three main merging components, the eastern component IC 694 (Arp 299 A) and the western components making up NGC 3690 (Arp 299 B and C) \\citep{geh83,wyn91}. The Arp 299 B nucleus is divided into subcomponents (B1 and B2) with a separation of $\\sim$2.5 arcsec in the infrared \\citep{wyn91,zho93,mil96,ket97,lai99,soi01}. The Arp 299 C region is taken to be a ``non-nucleus''\\citep{dud93,cha02} and consists of the C and C' components with a separation of $\\sim$5 arcsec in the infrared \\citep{zho93,mil96,ket97,lai99,soi01}. A component D is detected in the radio at 6 cm (5 GHz) \\citep{nef04}, but is not found in the deepest infrared image \\citep{soi01}. The whole merging system has an infrared luminosity of L$_{\\rm IR}$ $\\sim$ 5 $\\times$ 10$^{11}$L$_{\\odot}$ (= 10$^{45.3}$ ergs s$^{-1}$; see Table 1), approximately 60\\% of which comes from the A (IC 694) nucleus, based on observations using the {\\it Kuiper Airborne Observatory} \\citep{joy89,cha02}. Although no explicit AGN signatures have been found in any nucleus through optical spectroscopy \\citep{bus84,kee85,arm89,coz98}, highly-absorbed (Compton-thick; N$_{\\rm H}$ $>$ 10$^{24}$ cm$^{-2}$) but intrinsically strong X-ray emission has been detected at E $>$ 10 keV \\citep{del02}, suggesting the presence of a luminous buried AGN. Unfortunately, because of the lack of spatial resolution of the X-ray data at E $>$ 10 keV, it is unclear which nucleus (or nuclei) possesses the buried AGN. While the equivalent width of the 6.4-keV iron K$\\alpha$ line (EW$_{K \\alpha}$) in higher spatial-resolution X-ray data at E $<$ 10 keV was used to argue that a putative AGN resides in nucleus B1, at least \\citep{zez03,bal04}, it is unclear whether an AGN is present in any of the other nuclei because the 6.4-keV emission can be weak in a Compton-thick {\\it buried} AGN \\citep{fab02}. Further observational data are required to better understand the role of the putative buried AGN in Arp 299.  To spatially separate the subcomponents of the Arp 299 merging system, observations with spatial resolution better than a few to 5 arcsec are desirable. Ground-based infrared $K$- (2.0--2.5 $\\mu$m) and $L$-band (2.8--4.1 $\\mu$m) spectroscopy using a large ($>$3m) telescope at a good observing site can achieve sub-arcsec spatial resolution and so is suitable for this purpose. At millimeter wavelengths, interferometry is better suited than a single-dish radio telescope to obtain morphological information, by resolving the subcomponents, and to investigate the properties of molecular gas at nuclear regions where AGNs are thought to exist, by reducing the contamination from extended emission. In this paper, we report on the results of sub-arcsec-resolution infrared $K$- and $L$-band spectroscopy and millimeter interferometric maps with $\\sim$4 arcsec resolution, of the HCN (J = 1--0) and HCO$^{+}$ (J = 1--0) emission lines for the individual main merging components of Arp 299. The aim is to identify the location(s) of the putative buried AGN. Throughout this paper, $H_{0}$ $ = $ 75 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M}$ = 0.3, and $\\Omega_{\\rm \\Lambda}$ = 0.7 are adopted, so that the distance to Arp 299 is 40.3 Mpc and 1 arcsec corresponds to 190 pc. ",
        "conclusions": "\\subsection{Comparison to Previously Obtained Data}  \\subsubsection{Infrared $L$-band spectra}  The 3.3 $\\mu$m PAH emission from Arp 299 has been detected previously by several groups \\citep{miz94,den90,sat99,imd00}. The equivalent widths of the 3.3 $\\mu$m PAH emission (EW$_{\\rm 3.3PAH}$) measured from these older data are tabulated in Table 3 and are compared to our data.  For the Arp 299 A nucleus, the EW$_{\\rm 3.3PAH}$ values are similar for all the recently obtained data (this work; Satyapal et al. 1999; Imanishi \\& Dudley 2000). Older data presented by \\citet{den90} and \\citet{miz94} show smaller values, which could be attributed to larger scatter and narrower wavelength coverage for their spectra compared to the more recent observations.  For Arp 299 C, the EW$_{\\rm 3.3PAH}$ values reported here and by \\citet{miz94} are a factor of 3 lower than that derived by \\citet{sat99}. This difference can naturally be explained by the strong 3.1 $\\mu$m ice-covered dust absorption feature detected in our $L$-band spectrum. Because of the flux depression caused by this absorption feature, the actual continuum level is better derived from data at $\\lambda$ $>$ 3.4 $\\mu$m. The use of a data point at $\\lambda$ = 3.159 $\\mu$m \\citep{sat99} can lead to a major overestimate of the 3.3 $\\mu$m PAH emission strength.  For Arp 299 B1, our EW$_{\\rm 3.3PAH}$ value is substantially lower than those derived by \\citet{sat99} and by \\citet{miz94}. Unlike Arp 299 C, no strong 3.1 $\\mu$m absorption feature is detected in B1. We believe that this EW$_{\\rm 3.3PAH}$ difference is likely to be caused by an increased continuum flux in our data compared to the older data, for the following reasons. First, the $L$-band continuum flux level is more than 1 mag brighter than the value from the literature \\citep{zho93}. Although our slit spectra have some flux ambiguity ($\\S$2.1), this difference in B1 is difficult to account for solely by this ambiguity because the $L$-band continuum flux levels for A and C are similar to those from the same study \\citep{zho93}. Next, the B1 and B2 nuclei were simultaneously observed in our spectra, and yet the $L$-band continuum level in B1 relative to B2 is more than 1 mag brighter than expected \\citep{zho93}. Unless the B2 flux has decreased substantially, a flux increase in B1 is strongly suggested.  In Table 4, the 3.3 $\\mu$m PAH emission fluxes measured in previously obtained data are shown and compared with our data. We mainly use the highest quality recent data taken by \\citet{sat99} for comparison with our data. In the A nucleus, both datasets provide similar fluxes. For the B nucleus, the bulk of the 3.3 $\\mu$m PAH emission comes from B1 \\citep{sat99}, and our estimate (B1) and that by Satyapal et al. (B1+B2) are similar to each other. In the C region, the PAH emission is extremely spatially extended \\citep{sat99}, and the lower flux in our dataset can be explained by our smaller aperture, in addition to the above-mentioned PAH flux overestimate from their data. Overall, the comparison of PAH flux measurements suggests that our absolute flux calibration is reasonable, and thus the $L$-band continuum flux increase in B1 as implied above should be real.  \\subsubsection{Infrared $K$-band spectra}  Some parts of the $K$-band spectra of the individual merging components of Arp 299 have been previously obtained \\citep{nak89,rid94,shi94,sug99,alo00}. For the A and C components, the CO$_{\\rm spec}$ values, as well as the emission properties of the H$_{2}$ S(1) and Br$\\gamma$ (Table 5), are similar for our results and older spectra \\citep{nak89,rid94,shi94,sug99}. For the B1 nucleus, however, the rest-frame equivalent widths of both H$_{2}$ S(1) and Br$\\gamma$ emission lines in our data (Table 5) are a factor of $\\sim$3 smaller than those estimated from older data \\citep{alo00}. Our estimate of CO$_{\\rm spec}$ ($<$ 0.03) is also substantially lower than previous estimates (0.09--0.14; Ridgway et al. 1994; Alonso-Herrero et al. 2000). Finally, while the $K$-band continuum flux level for C is similar for our data and older spectra \\citep{alo00}, that for B1 is brighter in our data. All of these results for B1 support the argument that the flux of the featureless $K$-band continuum has increased markedly and dilutes the emission and absorption features in the $K$-band spectrum of B1.  \\subsubsection{Millimeter interferometric data}  For the continuum and HCN emission, our results can be compared to previously obtained interferometric maps \\citep{aal97,cas99} and to single-dish data \\citep{sol92}.  The continuum flux level was measured by \\citet{cas99} at $\\nu$ = 87.7 GHz to be 23.5 (A), 6 (B), 3 (C), and 4 mJy (C'). Our measurements at $\\nu \\sim$ 88 GHz (Table 6) are similar to these for all components.  The HCN emission peaks at the A nucleus in our map, corroborating previous interferometric maps \\citep{aal97,cas99}. The integrated HCN flux at the peak of the A nucleus (Table 8) and an integrated area surrounding it (Table 9) are similar within $\\sim$30\\% to the estimate by Aalto et al. (6.4 and 7 Jy km s$^{-1}$, respectively). In the map of \\citet{cas99}, clear HCN emission is detected in B, C, and C', while it is not in our map. In the HCN integrated intensity map of \\citet{cas99}, the noise level is lower than ours, and the highest contour at the B1 nucleus is 1.2 Jy km s$^{-1}$ beam$^{-1}$, which is lower than the 3$\\sigma$ level of our map. For the HCN emission from C and C', the overall spatial distribution is similar for our maps and those of \\citet{cas99}. In summary, the distribution of HCN emission from the whole Arp 299 merging system is consistent for all the available interferometric maps.  The linewidths are 310 km s$^{-1}$ and 280 km s$^{-1}$ in FWHM for HCN and HCO$^{+}$ emission, respectively, in the A nucleus. These values agree to within 20\\% with the $^{12}$CO linewidth (350 km s$^{-1}$; Aalto et al. 1997). In many LIRGs, the line profiles of HCN and $^{12}$CO have been found to be similar \\citep{gao04,ink06}, and our data show this to be the case for the Arp 299 A nucleus. For the Arp 299 B and C components, the FWHM value is better determined from HCO$^{+}$ than HCN in our data because of the much higher S/N ratios in the former. The HCO$^{+}$ linewidths in FWHM are 230 km s$^{-1}$ (B) and 130 km s$^{-1}$ (C). For comparison, the $^{12}$CO linewidths in FWHM are 260 km s$^{-1}$ (B) and 60--80 km s$^{-1}$ (C) \\citep{aal97}.  Finally, our new HCO$^{+}$ interferometric map, the first to be made for Arp 299, is compared to previous $^{12}$CO maps \\citep{sar91,aal97,cas99}. At the A nucleus, an extended tail toward the southeast direction is recognizable in both HCO$^{+}$ and $^{12}$CO maps. For the C' component, an extended HCO$^{+}$ emission component is detected on the eastern side of the C' peak, as was found in $^{12}$CO and $^{13}$CO maps \\citep{sar91,aal97,cas99}. However, while the $^{12}$CO emission peaks at C and C', the HCO$^{+}$ emission, as well as the HCN emission \\citep{cas99}, has a peak between the C and C' components, suggesting that high-density (n$_{\\rm H2}$ $>$ 10$^{4}$ cm$^{-2}$) molecular gas, as probed by HCN and HCO$^{+}$, is more concentrated in the in-between region than the diffuse (n$_{\\rm H2}$ $\\sim$ 10$^{2-3}$ cm$^{-2}$) gas probed by $^{12}$CO.  \\subsection{The Location of the Buried AGN--Implications from Infrared Spectra}  The infrared luminosity of Arp 299 is 2 $\\times$ 10$^{45}$ ergs s$^{-1}$ (Table 1). The measured 3.3 $\\mu$m PAH emission flux from the whole Arp 299 merging system (A+B+C) is $\\sim$1.8 $\\times$ 10$^{-12}$ ergs s$^{-1}$ cm$^{-2}$ \\citep{sat99}, which corresponds to a 3.3 $\\mu$m PAH luminosity of 3.5 $\\times$ 10$^{41}$ ergs s$^{-1}$, with the caveat that the actual value is slightly smaller because of the overestimate in C ($\\S$4.1.1). If we adopt a 3.3 $\\mu$m PAH to infrared luminosity ratio of $\\sim$10$^{-3}$, typical for starbursts \\citep{mou90,ima02}, then the luminosity of detected starbursts becomes 3.5 $\\times$ 10$^{44}$ ergs s$^{-1}$, which is a factor of 5 below the measured infrared luminosity. Although the extinction correction of the 3.3 $\\mu$m PAH emission can increase the intrinsic starburst luminosity, the correction factor is only a factor of 2 to 3 for dust extinction with A$_{\\rm V}$ $<$ 20 mag measured in Arp 299 \\citep{nak89,sat99,alo00}, if the Galactic dust extinction curve is assumed \\citep{rie85}. Taken at face value, it is quite possible that AGN activity can play an important role energetically. In this section, we investigate in detail the nuclei containing buried AGN.  \\subsubsection{Arp 299 B1}  The infrared $L$- and $K$-band spectra of the B1 nucleus favor the presence of a buried AGN there, based on the following three main arguments. First, the equivalent width of the 3.3 $\\mu$m PAH emission in the $L$-band spectrum is only EW$_{\\rm 3.3PAH}$ = 7 nm, more than an order of magnitude below typical values found in starburst galaxies (EW$_{\\rm 3.3PAH}$ $\\sim$ 100 nm: Moorwood 1986; Imanishi \\& Dudley 2000). Second, the CO index in the $K$-band spectrum is CO$_{\\rm spec}$ $<$ 0.03, much lower than typical values found in starburst galaxies (CO$_{\\rm spec}$ $>$ 0.15; Goldader et al. 1997a,b; Ivanov et al. 2000). A strong contribution from a featureless continuum, most likely originating from hot dust heated by an AGN, dilutes both the 3.3 $\\mu$m PAH emission and 2.3 $\\mu$m CO absorption features. Third, an increase in the $L$- and $K$-band continua is clearly recognizable, compared to old data taken about 5--10 years ago. Since no major flux variation is expected in a starburst over this time period, an AGN is the natural explanation. We conclude that the B1 nucleus contains a buried AGN, as was previously argued based on observations of 2--10 keV X-ray emission \\citep{zez03,bal04} and of the infrared 5--16 $\\mu$m band \\citep{gal04}. Assuming a Galactic extinction curve of $\\tau_{3.4}$/A$_{\\rm V}$ = 0.004--0.007 \\citep{pen94,ima96} and A$_{\\rm L}$ $\\sim$ 0.06 $\\times$ A$_{\\rm V}$ \\citep{rie85}, we can roughly estimate the dereddened luminosity of AGN-heated dust in the B1 nucleus from the observed $L$-band spectrum (optical depth of dust absorption features and observed continuum flux level at 3 to 4 $\\mu$m) \\citep{idm01,im03,idm06}, which is 3--5 $\\times$ 10$^{43}$ ergs s$^{-1}$.  \\subsubsection{Arp 299 C}  Our infrared $L$-band spectrum of C displays a strong 3.1 $\\mu$m ice absorption feature with $\\tau_{3.1}$ $\\sim$ 1.4. For a buried AGN in which the energy source is more centrally concentrated than dust, a foreground-screen dust-absorption geometry is applicable, while a mixed dust/source geometry is at work for a normal starburst in which stellar energy sources and dust are spatially well mixed \\citep{pux91,mcl93,for01}. In a normal starburst with mixed dust/source geometry, whose radiation density, and therefore the fraction of dust covered with an ice mantle, is similar to the well studied starburst galaxy M82, the $\\tau_{3.1}$ value is $<$0.3 for any amount of dust extinction of the starburst \\citep{im03}. Even if all dust is covered with an ice mantle ($f$ = 1 in Imanishi \\& Maloney 2003), the mixed dust/source geometry can only provide a maximum of $\\tau_{3.1}$ $\\sim$ 0.7, assuming the dust properties are similar to the Galactic interstellar medium. This is because the observed flux is dominated by foreground, weakly obscured, less attenuated emission, with a small contribution from obscured, highly attenuated emission at the far side, in this mixed dust/source geometry. The large $\\tau_{3.1}$ value in the Arp 299 C component ($\\sim$1.4) would require a foreground-screen dust geometry. A centrally concentrated energy source, such as a buried AGN, may be present, in addition to a normal, PAH-emitting starburst (Figure 1). Alternatively, if a large amount of foreground-screen dust exists exterior to the normal starburst core (e.g., in an edge-on host galaxy), then the large $\\tau_{3.1}$ can be explained. However, the EW$_{\\rm 3.3PAH}$ value is $\\sim$40 nm in C. In this geometry, the EW$_{\\rm 3.3PAH}$ value should have the value usually found for normal starbursts ($\\sim$100 nm) because both PAH and continuum emission are attenuated in a similar manner. Dilution of the PAH emission by PAH-free continuum from an energy source other than a normal starburst is indicated from the $L$-band spectrum of C. In the C component, the observed CO$_{\\rm spec}$ value ($\\sim$0.1) is also smaller than expected for a normal starburst ($>$0.15). Dilution by the featureless $K$-band continuum, possibly from AGN-heated hot dust or stars younger than $\\sim$10$^{6}$ yr \\citep{lei99}, is required to explain this result.  Unlike the A and B nuclei, Arp 299 C+C$'$ is thought to be a non-galaxy nucleus \\citep{dud93,cha02}. Under the assumption of an Eddington mass accretion rate, a more massive black hole is necessary to produce a larger AGN luminosity, and supermassive black holes are generally present at galaxy nuclei. Whether the C component possesses a sufficiently large supermassive black hole for AGN activity is less clear than for the other Arp 299 nuclei. \\citet{alo00} noticed a weaker CO absorption feature in the $K$-band spectrum of C, but attributed its weakness to the youth of the starbursts. If a very active super star cluster with an infrared luminosity of L$_{\\rm IR}$ $\\sim$ 10$^{9}$L$_{\\odot}$ \\citep{gor01} is formed in the innermost regions of dense molecular gas and the radiation density there is so high as to destroy PAH molecules \\citep{sel81}, then a large $\\tau_{3.1}$ value and PAH-free featureless $L$-band continuum \\citep{alo04} may be produced. For the nuclei of ultraluminous infrared galaxies with L$_{\\rm IR}$ $>$ 10$^{12}$L$_{\\odot}$, energy sources that are more centrally concentrated than dust would require both large luminosity and high surface brightnesses, for which AGNs are the natural explanation \\citep{soi00,soi03}. However, the infrared luminosity of Arp 299 C is at the level of L$_{\\rm IR}$ $\\sim$10$^{10}$L$_{\\odot}$ \\citep{cha02}. In Arp 299 C, the scenario of multiple young super star clusters embedded in dust and molecular gas may hold in term of absolute luminosity.  \\subsubsection{Arp 299 A}  Among the merging components of Arp 299, the A nucleus (IC 694) has long been suggested to be the most probable AGN-harboring nucleus, based on near-infrared color \\citep{nak89}, compact high surface-brightness radio emission \\citep{lon92,lon93}, and starburst modeling \\citep{geh83,shi96}. However, our infrared $L$- and $K$-band spectra reveal no explicit AGN evidence in this nucleus: its EW$_{\\rm 3.3PAH}$ and CO$_{\\rm spec}$ values are typical of starburst galaxies \\citep{moo86,gol97a,gol97b,iva00}. Recent high-quality infrared data \\citep{imd00,cha02,gal04} have also failed to provide clear AGN signatures from A. Although \\citet{bal04} inferred the presence of an AGN in the A nucleus, based on the presence of 6.7 keV ionized iron emission, their argument is weaker than that based on the 6.4-keV iron K$\\alpha$ line for the B1 nucleus. If a luminous buried AGN is present in the Arp 299 A nucleus, the AGN must be more elusive than in B1.  \\subsection{Interpretation of the HCN/HCO$^{+}$ Ratios}  \\citet{koh01} and \\citet{koh05} proposed that AGN-dominated galaxy nuclei (XDRs) tend to show higher HCN (J = 1--0) to HCO$^{+}$ (J = 1--0) and HCN (J = 1--0) to CO (J = 1--0) brightness-temperature ratios than starburst-dominated galaxy nuclei (PDRs), and thus that these ratios can be used to distinguish the two types of galaxies. In particular, both the HCN and HCO$^{+}$ molecules have similarly large dipole moments ($\\mu$ $\\sim$ 3--4 debye; Millar et al. 1997; Botschwina et al. 1993) and so probe similarly high-density (n$_{\\rm H2}$ $>$ 10$^{4}$ cm$^{-2}$) molecular gas. Hence, the HCN/HCO$^{+}$ ratios are very useful for distinguishing between the two galaxy types, being unaffected by different dense-to-diffuse molecular gas fractions in different galaxies. Imanishi et al. (2004; 2006b) applied this method to a few LIRG nuclei, and successfully found high HCN/HCO$^{+}$ brightness-temperature ratios in LIRGs known to possess luminous buried AGNs based on previously obtained data at other wavelengths. An HCN {\\it abundance} enhancement in XDR chemistry of dense molecular gas around a strong X-ray illuminating source (Meijerink \\& Spaans 2005; their Fig.10, model 4) can naturally explain the high HCN/HCO$^{+}$ brightness-temperature ratios in luminous AGNs, whether the HCN and HCO$^{+}$ emission is optically thin or thick \\citep{ink06}.  We use the HCN and HCO$^{+}$ flux measurements obtained with the NMA, to apply this energy diagnostic method to the individual merging components of Arp 299. For Arp 299, HCN and HCO$^{+}$ measurements using a single-dish radio telescope are available \\citep{gra06}. However, our interferometric data have the following advantages over the single-dish measurements: Although single-dish radio data measure the HCN and HCO$^{+}$ emission from the combined B+C regions, our interferometric data can clearly resolve B and C because of better spatial resolution. In addition, we can probe the dense molecular gas for regions confined to the nuclei where putative AGNs are expected to reside, by reducing the contamination from extended starburst emission. The XDR volume can be large in a {\\it pure} buried AGN with no starburst activity. However, in a starburst/buried-AGN composite galaxy, such as the majority of LIRGs, strong XDR signatures are expected only in the nuclear region in the close vicinity of a luminous X-ray-emitting AGN. Spatially extended molecular gas can be affected by the surrounding starbursts and thus show properties expected from PDRs. For the purpose of finding XDR signatures in starburst/AGN composite LIRGs, such as Arp 299, investigating the HCN/HCO$^{+}$ ratios in regions confined to the nuclei, based on high-spatial-resolution interferometric data, is superior to using single-dish data. Furthermore, in the old single-dish data, HCN and HCO$^{+}$ observations were made at different times, whereas for our NMA interferometric data, the HCN and HCO$^{+}$ lines were simultaneously obtained, so that possible uncertainties due to time-variable observing conditions and intercalibration are minimized, making the discussions of the HCN/HCO$^{+}$ brightness-temperature {\\it ratios} reliable.  Figure 8 shows the HCN (J = 1--0) to CO (J = 1--0) (abscissa) and HCN (J = 1--0) to HCO$^{+}$ (J = 1--0) (ordinate) ratios in brightness temperature of Arp 299 and other galaxies. For the Arp 299 merging components, the HCN/HCO$^{+}$ brightness-temperature ratios are directly derived from our NMA data, while the HCN/CO ratios are taken from the older interferometric data of \\citet{cas99}. The Arp 299 A, B, and C components show low HCN/HCO$^{+}$ brightness-temperature ratios, as observed in starburst-dominated galaxy nuclei.  Table 10 summarizes the detection or non-detection of AGN signatures based on various methods. Since a luminous buried AGN is almost certainly present in the B1 nucleus, we must consider why no XDR signature is found in B based on the HCN/HCO$^{+}$ ratio. Several explanations are possible, but whichever scenario is correct, we have to conclude that the HCN/HCO$^{+}$ method may miss some types of buried AGN despite their detection at other wavelengths.  First, the dereddened AGN luminosity at the B1 nucleus derived from our infrared $L$-band spectrum is L$_{\\rm IR}$(AGN) $\\sim$ 3--5 $\\times$ 10$^{43}$ ergs s$^{-1}$ ($\\S$4.2.1). The absorption-corrected 2--10 keV X-ray luminosity is $\\sim$5 $\\times$ 10$^{42}$ ergs s$^{-1}$ \\citep{del02}. If we apply 2--10 keV to the bolometric luminosity correction of $\\sim$10--20 \\citep{elv94}, the estimated AGN luminosity becomes roughly similar, within a factor of 2--3, to the value quoted above based on our infrared $L$-band spectrum. However, the detected starburst luminosity in the B1 nucleus is estimated to be L$_{\\rm IR}$(SB) $\\sim$ 4--5 $\\times$ 10$^{43}$ ergs s$^{-1}$ based on the observed 3.3 $\\mu$m PAH emission luminosity (This work; Satyapal et al. 1999) and on the ratio of L$_{\\rm 3.3 PAH}$/L$_{\\rm IR}$ $\\sim$ 10$^{-3}$ found in starbursts \\citep{mou90,ima02}. If we assume dust extinction of A$_{\\rm V}$ = 6--15 mag for the detected starbursts in B \\citep{geh83,bec86,sat99}, the extinction-corrected starburst luminosity can increase to L$_{\\rm IR}$(SB) $\\sim$ 7--10 $\\times$ 10$^{43}$ ergs s$^{-1}$. Thus, in the B nucleus, the starburst luminosity is comparable to the AGN luminosity. The presence of high-density molecular gas strongly affected by this luminous starburst component may hide the signatures of XDRs traced by our HCN/HCO$^{+}$ method.  Second, for the HCN/HCO$^{+}$ method to work in finding XDR signatures, the high-density molecular gas, as probed by HCN and HCO$^{+}$, must be sufficiently close to be affected by the central AGN's X-ray radiation. In the double nuclei merging LIRG NGC 6240, the HCN and HCO$^{+}$ emission peaks are not at individual nuclei, but in the internuclear region where shocks are important \\citep{nak05}. In Arp 299, although the HCO$^{+}$ emission peak from B (B1+B2) appears to coincide with the B1 nucleus in our interferometric map with 4$\\farcs$2 $\\times$ 3$\\farcs$8 spatial resolution (Figure 4), it is possible that the actual distribution is slightly displaced from the exact position of the B1 nucleus to the B2 side, and shocks may play a major role. It is known that shock regions show enhanced HCO$^{+}$ abundance \\citep{dic80}, which can decrease the HCN/HCO$^{+}$ brightness-temperature ratio, compared to a pure buried AGN. In fact, shocks driven by starburst superwinds are found to be prominent in the whole area of the Arp 299 merging system \\citep{hec99}. These spatially extended shocks at the surface of the buried AGN may decrease the HCN/HCO$^{+}$ brightness-temperature ratio measured toward the B1 nucleus.  Third, an increase in AGN luminosity compared to older data is indicated in our infrared spectra of B1 ($\\S$4.2.1). Some time lag may perhaps occur between the onset of an active AGN phase and the stage when the X-ray-emitting AGN can chemically affect the surrounding high-density molecular gas. If the very active AGN phase has just started, then the low HCN/HCO$^{+}$ brightness-temperature ratio may be explained if the surrounding molecular gas has not been subjected to sufficient chemical effects.  Finally, the HCN/HCO$^{+}$ abundance ratio is highly dependent on the density and distribution of molecular gas surrounding a central X-ray-emitting AGN \\citep{mei05}. The Arp 299 merging components may have different molecular gas distributions (e.g., density and geometry) from other LIRGs for which our HCN/HCO$^{+}$ method succeeded. Further detailed chemical calculations are needed to interpret more unambiguously the observed HCN/HCO$^{+}$ brightness-temperature ratios in galaxies."
    },
    "0608/astro-ph0608623_arXiv.txt": {
        "abstract": "We use a purely data-driven rectified factor analysis to identify early-type galaxies with recent star formation in Data Release 4 of the Sloan Digital Sky Survey Spectroscopic Catalogue. We compare the spectra and environment of these galaxies with those of `normal' early-type galaxies, and a sample of independently selected \\ea\\ galaxies. We calculate the projected local galaxy surface density from the nearest five and ten neighbours ($\\Sigma_{5}$ and $\\Sigma_{10}$) for each galaxy in our sample, and find that the dependence on projected local density, of the properties of \\ea\\ galaxies, is not significantly different from that of early-type galaxies with young stellar populations, dropping off rapidly towards denser environments, and flattening off at densities \\ls\\ 0.1$-$0.3 Mpc$^{-2}$. The dearth of \\ea\\ galaxies in dense environments confirms that \\ea\\ galaxies are most likely the products of galaxy-galaxy merging or interactions, rather than star-forming galaxies whose star formation has been quenched by processes unique to dense environments, such as ram-pressure stripping or galaxy harassment. We see a tentative peak in the number of \\ea\\ galaxies at $\\Sigma_{10} \\sim 0.1-0.3$ Mpc$^{-2}$, which may represent the local galaxy density at which the rate of galaxy-galaxy merging or interaction rate peaks. Analysis of the spectra of our early-type galaxies with young stellar populations suggests that they have a stellar component dominated by F stars, $\\sim$ 1$-$4 Gyr old, together with a mature, metal-rich population characteristic of `typical' early-type galaxies. The young stars represent \\gs 10\\% of the stellar mass in these galaxies. This, together with the similarity of the environments in which this `E$+$F' population and the \\ea\\ galaxy sample are found, suggests that E+F galaxies used to be \\ea\\ galaxies, but have evolved by a further $\\sim$ one to a few Gyr. Our rectified factor analysis is sensitive enough to identify this hidden population, which allows us to study the global and intrinsic properties of early-type galaxies created in major mergers or interactions, and compare them with those early-types which have had the bulk of their stars in place since a much earlier epoch. ",
        "introduction": "There is still much uncertainty about the relative epochs of mass assembly and star formation in early-type galaxies (e.g. de Lucia et al. 2006; Cimatti, Daddi \\& Renzini 2006). The most massive galaxies contain the oldest stellar populations, but hierarchical structure formation requires that their present-day mass was assembled at a recent epoch. On the other hand, galaxy `downsizing' (e.g. Cowie et al. 1996; Treu et al. 2005) means that less massive galaxies have more recent star formation.  One would expect that different formation histories would give rise to different observed characteristics in early-type galaxies. Several studies (e.g. Faber et al. 1997; Emsellem et al. 2004; Khochfar \\& Burkert 2005) have in fact found that various sub-classes of early-type galaxy exist (e.g. `core' and `cuspy' galaxies, fast and slow rotators). We would like to know how these sub-classes are related to the various modes of formation of early-type galaxies, and to the effects of their immediate environments.  The stellar populations of early-type galaxies hold the key to understanding the difference between early-types created via different evolutionary routes. If we can identify those galaxies which have recently undergone a strong starburst, and we associate that recent starburst with a major galaxy-galaxy merger or interaction, then we can identify the products of recent merger or interaction, map their evolution with environment and redshift, and study their intrinsic properties in comparison with more massive `primordial' early-types.  In this work, we are particularly interested in the relationship between \\ea\\ galaxies and `typical' early-types. \\ea\\ galaxies are defined as having strong hydrogen Balmer absorption lines, but no corresponding [OII] or \\ha\\ emission lines. Their spectra, therefore, resemble the linear superposition of a characteristic early-type galaxy spectrum and that of an A star. This suggests that these galaxies have undergone a burst of star formation in the last Gyr or so, but that there is no on-going star formation. It is hypothesised that these galaxies will eventually settle down as `normal' passively-evolving early-types.  \\ea\\ galaxies may therefore represent an intermediate stage between gas-rich, disc-dominated, rotationally-supported, star-forming systems on the one hand, and bulge-dominated, pressure-supported, passive systems on the other. Hence, understanding the physical processes by which the star formation is triggered and subsequently truncated, is a vital stage in understanding the evolution of early-type galaxies.  There are three possible scenarios which could explain the existence of \\ea\\ galaxies. The first is that they are in fact dust-obscured star-forming galaxies. This can be verified by observing at radio wavelengths, where star formation can be detected, but dust obscuration is not significant. Miller \\& Owen (2001) and Goto (2004) observed between them a total of 51 \\ea\\ galaxies. Moderate star formation was detected in only two of these galaxies, so it is unlikely that the majority of \\ea s are dusty starbursts.  More plausibly, then, \\ea\\ galaxies may be the product of disk-disk merging or interaction, at a stage when the starburst has burnt itself out, or they may be galaxies whose star formation has been initiated and then quenched by, for example, ram-pressure stripping or tidal stripping. The environment of these galaxies therefore holds the key to understanding them. Although, at intermediate redshift (0.33 $<$ z $<$ 0.83), Tran et al. (2003) find that \\ea s make up a non-negligible fraction ($\\sim$7--13\\%) of cluster members, suggesting a cluster-related mechanism for the truncation of star formation, such as ram-pressure or tidal stripping, at lower redshift (z $\\sim$ 0.1), \\ea s appear to occur in lower-density environments, indicative of merging or interaction processes (Zabludoff et al. 1996; Blake et al. 2004). In addition, both Zabludoff et al. (1996) and Blake et al. (2004) found a significant number of their \\ea\\ galaxies displayed tidal features and disturbed morphologies, typical of merger remnants. More recently, Goto (2005) found that \\ea s selected from the Sloan Digital Sky Survey (SDSS) Data Release~2 have an excess of local galaxy density only at scales of $<$ 100 kpc, and not at the scale of either clusters or larger-scale structures. This is also highly suggestive of a merger or interaction origin for \\ea s.  Conventional methods of identifying \\ea\\ galaxies use measurements of the equivalent width of hydrogen Balmer absorption lines (typically \\ha\\ and H$\\delta$) and [OII] emission. However, this can suffer from dust obscuration of the emission lines, and emission-filling of H Balmer lines, and it requires a high signal-to-noise ($\\ge$ 10 per pixel). In addition, there has been some debate over the criteria for selecting \\ea\\ galaxies, with some authors choosing \\ha , H$\\delta$ and [OII], with differing line strength criteria (e.g. Goto 2004, 2005), and some including H$\\gamma$ as well as \\ha, H$\\delta$ and [OII] (e.g. Zabludoff et al. 1996; Blake et al. 2004).  Here, we use a purely data-driven rectified factor analysis data model, developed in Nolan et al. (2006), on the spectra of early-type galaxies in the SDSS Data Release 4 (DR4). This was specifically designed to exploit large data sets in a model-independent way. It uses all the data present in the spectra, i.e. the spectral features and the overall shape of the continuum, to quantify the relative strength of recent star formation in early-type galaxies extracted from the database, in a rapid, objective and robust manner. Our method allows us to reconstruct even partial spectra, without the need to measure individual lines accurately. The advent of large surveys, such as the Sloan Digital Sky Survey allows us to study large samples of these rare objects ($<$ 1 $\\%$ of the overall zero-redshift galaxy population, Zabludoff et al., 1996), in a way that has not been previously possible.  It has recently been shown that \\ea\\ galaxies occur more frequently at high redshift (1 $\\leq$ z $\\leq$ 2) compared with the local universe, at least amongst the massive red galaxy population (Doherty et al. 2005, Le Borgne et al. 2006). They may, therefore, represent a critical phase in the evolution of early-type galaxies. However, it is difficult to constrain the properties of these objects at high redshift, due to the difficulty of measuring absorption lines accurately at that redshift, and in obtaining statistically significant spectroscopic samples. Studying large samples at low redshift is therefore vital for the understanding of these objects.   In this work, we use our data modelling to identify early-type galaxies (ETGs) which have a young stellar population, suggesting that they have undergone a recent burst of star formation. We compare these galaxies with \\ea\\ galaxies, selected in a completely independent manner, and investigate the environment of our galaxy samples, in order to understand the physical mechanisms leading to the cessation of star formation.  In Section \\ref{sample}, we define the selection of the early-type sample, and the \\ea\\ subsample for comparison. We discuss the factor analysis of the early-type spectra that is used to reveal young stellar populations in Section \\ref{model}. In Section \\ref{anal}, we compare our results from the analysis of early-type galaxies with the sample of \\ea\\ galaxies. In  Section \\ref{environ}, we investigate the relationship between the galaxies in our various samples and their local environment. Our conclusions are discussed in Section \\ref{conc}.  We use $H_{0} = 70$ kms$^{-1}$Mpc$^{-1}$, $\\Omega_{m} =$0.3 and $\\Omega_{\\Lambda} =$0.7 throughout. ",
        "conclusions": "\\label{conc}   We have successfully used our data-driven rectified factor analysis to identify early-type galaxies with young stellar populations, without having to apply the far more computationally expensive method of fitting detailed stellar population synthesis models to a sample of $>$ 13,000 ETGs in the redshift range 0.06 $<$ z $<$ 0.14. We have used the sample generated as a result to map the relationship between early-type galaxies, \\ea\\ galaxies and early-type galaxies with young stellar populations. Our analysis recovers two components from the set of early-type galaxy spectra, one of which represents the contribution to the flux from mature, metal-rich stars, and the other which represents the contribution from F stars, $\\sim$1--4~Gyr old. Our method allows us to reconstruct spectra even when the observed spectra are incomplete. We do not need high signal-to-noise measurements of individual lines. We identify those early-types which have undergone recent ($\\sim$1--4~Gyr) star formation, with the young stars accounting for at least 10\\% of their stellar mass (YSPETs). At redshifts 0.06 $<$ z $<$ 0.14, $\\sim$17\\% of ETGs are classified as YSPETs. We propose that these YSPETs are `E$+$F' galaxies, whose spectra are a linear combination of a typical early-type galaxy spectrum and that of an F star, analogous to \\ea\\ spectra.  We also select a sample of \\ea\\ galaxies in a completely independent way, without using  the morphological criteria we used to identify early-type galaxies, but instead relying on the emission and absorption-line criteria of Goto (2005). The characteristic `\\ea' spectra of the Goto-selected sample, and the fact that many of these \\ea s do not fulfil the colour and concentration criteria for typical early-type galaxies, suggests that these galaxies are in an early ($<1$ Gyr) post galaxy-galaxy merger or interaction phase. Goto (2005) notes that most of the \\ea s in his sample have concentrations as high as typical ellipticals ($C_{in}\\! <\\! 0.4$), so it is likely that many of these galaxies are not included in our early-type sample because they are too blue, indicating that massive (A star and earlier) stars are still burning. If \\ea s are recently quenched star-forming galaxies, we would expect them to be located preferentially in dense cluster regions, whereas we see the majority residing in field or group environments, and the number falls off towards denser environments, in agreement with Goto's result. Their high concentrations and preferentially low-density environments argue strongly that these are the products of galaxy-galaxy mergers or interactions, which will eventually settle down to become typical early-type galaxies.  We note that there is in fact a slight peak in the local galaxy surface density distribution of \\ea\\ galaxies at 0.1--0.3 Mpc$^{-2}$. This may represent the optimum local galaxy surface density for galaxy-galaxy mergers and interactions. Realistic cosmological simulations should be able to confirm or reject this hypothesis.  Given that the distribution of local galaxy surface density of the G\\ea\\ galaxies and that of the independently determined YSPET populations are not significantly different, one might expect a similar history for both sets of galaxies. We hypothesise that the YSPETs, with typical ETG colours and concentrations, and `E$+$F' spectra, represent an intermediate evolutionary stage between \\ea s and early-types. As this is a longer-lived stage than \\ea\\ galaxies (F stars persist on the main sequence for longer than A stars, and dominate the integrated flux at ages \\ls 5 Gyr), we are able to identify a larger sample of these galaxies than is possible with the very rare \\ea s. We therefore have identified a powerful tool for studying in detail, both globally and individually, the evolution of early-type galaxies.  Our future work will compare the dependence on local density of \\ea\\ and E$+$F galaxies that we have recovered here, with those recovered from simulations. We intend to perform the same analysis on 2dFGRS galaxies and a sample of high-redshift galaxies for comparison. We also intend to investigate the relationship between \\ea\\ and E$+$F galaxies by studying their kinematics in detail using integral field spectroscopy."
    },
    "0608/astro-ph0608553_arXiv.txt": {
        "abstract": "Several kinds of measurements are combined in an attempt to obtain a consistent estimate of the spectrum and composition of the primary cosmic radiation through the knee region. Assuming that the knee is a signal of the high-energy end of a galactic cosmic-ray population, I discuss possible signatures of a transition to an extra-galactic population and how they might be detected. ",
        "introduction": "\\label{sec1}  \\begin{figure} \\begin{center} \\includegraphics[width=13cm]{forPO+Auger2.ps}  \\end{center} \\caption{Spectrum of cosmic rays} \\label{fig1} \\end{figure}   The cosmic-ray spectrum extends from the sub-GeV region to at least $10^{11}$~GeV (Fig.~\\ref{fig1}.)  Up to $100$~GeV and somewhat higher, measurements with magnetic spectrometers flown above most of the atmosphere provide good momentum resolution along with identification of the charge and mass of individual primaries.   Measurements with calorimeters continue to identify individual primaries to beyond $100$~TeV, but with larger systematic uncertainties in the energy assignment. In the PeV region and beyond, the cosmic-ray intensity is too low for direct measurements; only indirect measurements of air showers from the ground are possible.  Since the particles are not identified on an event-by-event basis, the energy spectrum derived from measurements of air showers is given as an ``all-particle\" spectrum, in terms of energy per particle rather than energy per nucleon.  In the air-shower regime, identification of the primary mass is made in one of several indirect ways on a statistical basis, complicating the search for features in spectra of individual elements.  The gyroradius of a proton in a typical galactic magnetic field of 3~$\\mu$Gauss is about half a parsec at one PeV.   The parsec scale is also typical of the size of structures in the interstellar medium driven by supernova explosions.  For particles with higher energy and larger gyroradii, diffusion in the interstellar medium may become less efficient.  Estimates of the maximum energy of particles accelerated at supernova shocks are around the same energy. It is therefore natural to guess that the knee of the cosmic-ray spectrum around 3 PeV has something to do with the end (or at least the beginning of the end) of the galactic cosmic-ray spectrum.  Peters~\\cite{Peters} described the consequences for energy dependence of the primary composition if the spectrum is characterized by a maximum rigidity, $R_c$, which could be associated either with propagation or with acceleration (or both).  The relation between rigidity and total energy is \\begin{equation} R\\;=\\;{P\\,c\\over Z\\,e}, \\end{equation} where $P\\,c\\approx E_{\\rm tot}$ is the total energy of a nucleus of charge $Z$ and mass $A$.  If cosmic rays are classified by energy per particle, as is the case for air shower measurements, then the spectrum should steepen first for protons, then for helium, then for the the CNO group etc. Elsewhere~\\cite{erix} I have called this sequence the \"Peters cycle\".  Several air shower measurements show some evidence that the spectrum becomes progressively enriched in heavy nuclei through the knee region.  The clearest evidence for the Peters sequence for several groups of nuclei comes from analysis of the KASCADE experiment~\\cite{KASCADE}. ",
        "conclusions": ""
    },
    "0608/hep-ph0608299_arXiv.txt": {
        "abstract": "The parameter space for $A$-term inflation is explored with $W=\\lambda_p \\phi^p/p\\mpl^{p-3}$. With  $p=6$ and $\\lambda_p\\sim 1$, the  observed spectrum and spectral tilt can be obtained with soft mass of order $10^2\\GeV$ but not with a much higher mass. The case $p=3$ requires $\\lambda_p\\sim 10^{-9}$ to $10^{-12}$. The ratio $m/A$ requires fine-tuning,  which may be justified on environmental grounds. An extension of the MSSM to include non-renormalizable terms and/or Dirac neutrino masses might support either $A$-term inflation or modular inflation. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/hep-th0608200_arXiv.txt": {
        "abstract": "Recently a non-inflationary mechanism of generation of scale-free cosmological perturbations of metric was proposed by Brandenberger, Nayeri,  and Vafa in the context of the string gas cosmology.  We discuss various problems of their model and argue that the cosmological perturbations of metric produced in this model have blue spectrum with a spectral index $n = 5$, which strongly disagrees with observations. We conclude that this model in its present form is not a viable alternative to inflationary cosmology. ",
        "introduction": "At present, inflation is the only well established mechanism which solves the homogeneity, isotropy, flatness and horizon problems, and explains why the universe has such a large mass and entropy. In addition, it provides a simple mechanism of generation of metric perturbations of with a flat spectrum, as required for the formation of large-scale structure of the universe. For a review of inflationary theory and the theory of perturbations of metric there see e.g.  \\cite{book,mukhbook}.  There were many attempts to suggest an alternative solution to the major cosmological problems, but so far the progress in this direction was quite limited. For example, the pre-big bang scenario \\cite{PBB} and the ekpyrotic/cyclic scenario \\cite{cyclic} are based on the assumption that eventually one will find a solution of the cosmological singularity problem and learn how one could transfer small perturbations of metric through the singularity. This problem, as well as several other problems discussed in \\cite{nogo,KKL,Linde:2002ws}, still remain unsolved. An alternative explanation of the origin of cosmological perturbations was advocated in \\cite{Hollands:2002yb}. The mechanism was based on the alternation of the QFT rules. However this mechanism does not address any other cosmological problems, and it relies on speculations about dynamics of perturbations  at an epoch when the energy density was $10^{95}$ times greater than the Planck density \\cite{Kofman:2002cj}. In comparison, inflation solves the major cosmological problems and produces metric perturbations with a flat spectrum practically independently of the speculative processes at extremely high energies and large curvatures, or before the big bang.  In a recent series of papers \\cite{Nayeri:2005ck,Brandenberger:2006xi,Nayeri,Brandenberger:2006vv} it was claimed that after a certain modification of the string theory inspired cosmological model proposed in \\cite{Brandenberger:1988aj,Tseytlin:1991xk} one can obtain scale free cosmological perturbations of metric due to the physical processes {\\it after} the big bang. This could be  advantageous over the pre-big bang and the ekpyrotic/cyclic scenario, because in those models the cosmological perturbations are produced {\\it before} the big bang, making the corresponding conclusions unreliable. The authors of \\cite{Nayeri:2005ck,Brandenberger:2006xi,Nayeri,Brandenberger:2006vv} admitted that their model does not solve the flatness problem and has several other unexplained features. Thus it is not a real alternative to inflation. Nevertheless it would be quite interesting to have a   non-inflationary mechanism of production of perturbations of metric with a flat spectrum. Our goal here is to examine this issue.  In Section 2 we will go over some of the properties and unsolved problems of the basic cosmological model of \\cite{Nayeri:2005ck,Brandenberger:2006xi,Nayeri,Brandenberger:2006vv}. In Section 3 we will argue that the spectrum of scalar perturbations produced in this model is dramatically tilted. Instead of the scale-free spectrum with the spectral index $n = 1$, the perturbations produced in the model of \\cite{Nayeri:2005ck,Brandenberger:2006xi,Nayeri,Brandenberger:2006vv} have blue spectrum with the spectral index $n = 5$, disagreeing with the cosmological data. The main problem with the proposal of \\cite{Nayeri:2005ck,Brandenberger:2006xi,Nayeri,Brandenberger:2006vv} stems from the fact that they worked in the string frame, but directly employed the standard description of metric perturbations \\cite{mukhbook} tailored for the Einstein frame. Meanwhile, due to a nontrivial dilaton dynamics these two frames strongly differ from each other at the epoch when the perturbations of metric were produced. In Section 4 we will show that a similar conclusion is  valid for the tensor perturbations of metric as well (though for a different reason): the spectrum of the gravitational waves produced in their scenario is far from being flat. ",
        "conclusions": "Our investigation shows once again that it is very difficult to offer a non-inflationary solutions of the major cosmological problems. It remains quite instructive, however, to learn why is it so difficult to find an alternative to inflation. In this particular case, the difficulties start from the very beginning, since even the existence of the basic cosmological solution is rather problematic, see Section \\ref{basic}. Nevertheless, one could still nurse a hope that this situation could somehow be improved in future development of the model, by adding new ingredients, finding a better effective action, etc.  However, we do not think that this could help. Our main concern stems from the fact that one of the ingredients of production of the cosmological perturbations in the scenario proposed in \\cite{Nayeri:2005ck,Nayeri,Brandenberger:2006vv} is purely kinematic: Perturbations should freeze when they cross the horizon. This, of course, is the main idea of the inflationary mechanism of generation of perturbations of metric: during inflation the wavelength of the inflaton perturbations grow, and when the wavelength becomes greater than the size of the horizon, the perturbations freeze. The new idea of Ref. \\cite{Nayeri:2005ck,Nayeri,Brandenberger:2006vv} is that in their model the wavelengths of the perturbations grow very slowly, but the quantity $H_{s}^{{-1}}$ should rapidly shrink. Then, they posit, perturbations freeze when $H_{s}^{{-1}}$ becomes smaller than the wavelengths of the perturbations. However, this was just a conjecture, based on an analogy with the results obtained in inflationary cosmology.  The easiest way to test this conjecture in the model proposed in \\cite{Nayeri:2005ck,Brandenberger:2006xi,Nayeri,Brandenberger:2006vv} is to study the evolution of the perturbations in the Einstein frame. In this frame, the perturbations shown in Fig. 1 from the very beginning have the wavelength much greater than $H_{E}^{{-1}}$, so the analogy with the inflationary mechanism breaks down. This is one of the reasons why the spectral index of the cosmological perturbations generated in this scenario is unacceptably large, $n = 5$.  Of course, it might happen that eventually one will find a version of this model where the size of the horizon $H_{E}^{{-1}}$ evolves extremely slowly, so that for a long time it remains nearly constant. Then the cosmological evolution will cause the wavelengths of perturbations to cross the horizon, which, under certain conditions, may lead to their freezing. But the slow evolution of the Hubble constant $H_{E}$ is a definitive feature of inflationary cosmology. Thus, at the moment we do not see any reason to expect that this model or its possible generalizations can provide a viable alternative to inflation.  {\\bf Note Added:} The main results of this paper were reported and debated with the authors of \\cite{Nayeri:2005ck,Brandenberger:2006xi,Nayeri,Brandenberger:2006vv} during the conference ``Inflation+25'', Paris, June 2006. Recently Brandenberger, Nayeri, Patil and collaborators issued a paper \\cite{Brandenberger:2006pr} where they confirmed our statement that the perturbations produced in the original scenario of \\cite{Nayeri:2005ck,Brandenberger:2006xi,Nayeri,Brandenberger:2006vv} have an unacceptably large spectral index $n = 5$. They argued that this scenario can be salvaged if the dilaton in the early universe were fixed  by some unspecified strong coupling effects. We discuss this speculative possibility in Sect.  \\ref{para}, and conclude that it violates the null energy condition. Even if such regime were possible, we show that the universe in this regime would accelerate both in the string frame and in the Einstein frame. In other words, the universe in this speculative regime is inflationary, so it does not provide a true alternative to inflation."
    },
    "0608/astro-ph0608547_arXiv.txt": {
        "abstract": "By combining photometry from the 2MASS Point Source Catalog and the USNO-B1.0 Catalog with optical and infrared spectroscopy, I have performed a search for young brown dwarfs in an area of 225~deg$^2$ encompassing all of the Taurus star-forming region ($\\tau\\sim1$~Myr). From this work, I have discovered 22 new members of Taurus, five of which were independently found by Guieu and coworkers. Sixteen of these new members have spectral types later than M6 and thus are likely to be brown dwarfs according to the theoretical evolutionary models of Chabrier and Baraffe. After adding these new members to the previously known members of Taurus, I have compared the spatial distributions of stars and brown dwarfs across the entire region. I find no statistically significant difference between these two distributions. Taurus does not contain the large, extended population of brown dwarfs that has been predicted by some embryo ejection models for the formation of brown dwarfs. However, these results are consistent with other ejection models, as well as models in which stars and brown dwarfs share a common formation mechanism. ",
        "introduction": "\\label{sec:intro}    Measuring the spatial distribution of a population of newly-formed stars can offer insight into the star formation process. Because of its youth and low stellar density, the stellar content within the Taurus-Auriga molecular cloud probably has undergone the least dynamical evolution of any nearby star-forming region, and therefore is the best available site for measuring a primordial spatial distribution of young stars. Several studies over the last decade have taken advantage of this fact to address various aspects of star formation. For instance, measurements of the distribution of stars in Taurus have revealed effects of both binary formation and the Jeans condition on clustering properties \\citep{gom93,lar95,sim97,nak98,bat98} while a comparison of the distributions of stars and molecular material in Taurus has constrained the velocity dispersion with which the stars are born and the nature of cloud fragmentation \\citep{har02}.  Measuring the spatial distribution of Taurus members as a function of mass is another potentially rewarding experiment. Several studies have suggested that brown dwarfs might form as protostellar sources whose accretion is prematurely halted by ejection from multiple systems \\citep{rc01,bos01,bat02,del03,umb05}. Some of these ejection models have predicted that newborn brown dwarfs could have higher velocity dispersions than their stellar counterparts \\citep{rc01,kb03}. If so, brown dwarfs would be more widely distributed than stars in star-forming regions. For instance, \\citet{kb03} favored an ejection model with a one-dimensional velocity dispersion of $\\sim2$~km~s$^{-1}$, which corresponds to an angular distance of $0\\fdg8$ for an object traveling for 1~Myr at the distance of Taurus. Meanwhile, other models of ejection predict that stars and brown dwarfs should have similar spatial and velocity distributions \\citep{bat03}.  Over time, surveys for substellar members of Taurus have encompassed larger areas surrounding the stellar aggregates. Some of these data have exhibited no statistically significant differences in the spatial distributions of the high- and low-mass members of Taurus \\citep{bri02,luh04tau} while other data have suggested a possible difference \\citep{gui06}. To search for an extended population of brown dwarfs beyond the previous survey fields and to measure the distribution of brown dwarfs in Taurus on the largest size scales, I have performed a survey for brown dwarfs across an area of 225~deg$^2$ containing the entire Taurus star-forming region. In this paper, I describe the selection of candidate substellar members of Taurus (\\S~\\ref{sec:cand}) and the spectroscopy and classification of these candidates (\\S~\\ref{sec:spec}), evaluate the completeness of this survey and the recent one by \\citet{gui06} (\\S~\\ref{sec:complete}), and use the resulting updated census of Taurus to perform a definitive comparison of the distributions of stars and brown dwarfs in this region (\\S~\\ref{sec:spatial}). ",
        "conclusions": "To measure the spatial distribution of substellar members of Taurus, I have performed a search for brown dwarfs across the entire star-forming region. The results of this work can be summarized as follows:  \\begin{enumerate}   \\item Through analysis of photometry from the 2MASS and USNO-B1.0 catalogs for an area of 225~deg$^2$ encompassing Taurus, I have identified 83 potential substellar members of the region. Spectroscopy was performed on 15 of these candidates by \\citet{luh04tau}, who classified seven of them as new members. Among the 68 remaining candidates observed in this work, I have identified 22 new members, five of which were independently discovered by G06.  Sixteen of these 22 objects have spectral types later than M6 and thus are likely to be brown dwarfs according to the theoretical evolutionary models of \\citet{bar98} and \\citet{cha00} and the temperature scale of \\citet{luh03b}.  \\item Among the candidates classified as non-members, three objects have late spectral types (M8-L1) and spectral features that are suggestive of surface gravities between those of Taurus members and typical field dwarfs. I speculate that these objects could be young members of the field ($\\tau\\sim100$~Myr).  \\item Using the low-mass Taurus members found in previous surveys, I estimate that my survey has a completeness of $\\sim75$\\% for members with spectral types of $>$M6 to M9 and extinctions of $A_V\\leq4$. The same type of analysis demonstrates that the completeness of the recent survey by G06 is $<63$\\% for this range of types.    \\item G06 concluded that the abundance of brown dwarfs relative to stars varies with distance from the stellar aggregates in Taurus. I find that this apparent variation is a reflection of differences in the spectral classification systems used for low-mass members inside and outside of the aggregates. No variation is evident when spectral types for all objects are adopted from the same system. Furthermore, a reliable measurement of the relative numbers of stars and brown dwarfs as a function of position in Taurus is currently not possible because the available census of Taurus members has an unknown level of completeness for low-mass stars (M2-M6, 0.1-0.6~$M_{\\odot}$) for most of the region.  \\item After updating the census of Taurus members with the new objects discovered in this work, I find that the spatial distribution of brown dwarfs closely follows that of the stars. An extended population of brown dwarfs outside of the stellar aggregates is not present. These results are consistent with a common formation mechanism for stars and brown dwarfs \\citep[e.g.,][]{pn04} and with some models for embryo ejection \\citep{bat03}, but not others \\citep{kb03}.  \\end{enumerate}"
    },
    "0608/astro-ph0608637_arXiv.txt": {
        "abstract": "We derive the density profile of cold dark matter halos using a self-similar accretion model. We show that if the clumpiness of the infalling matter is taken into account, then the inner density slope, $\\delta =d\\log{\\rho}/d\\log{r}$, is close to $-1$. Compared with the density profiles predicted by different numerical simulations, we find that outside $\\sim 0.1\\%$ of the virial radius, our solutions agree best with the fitting formula proposed by Navarro et al. (2004), $d\\ln{\\rho}/d\\ln{r}=-2(r/r_{-2})^\\alpha$, $\\alpha\\sim 0.17$, where $r_{-2}$ is a characteristic radius, inside which the density profile becomes shallower than isothermal ($\\delta<-2)$. ",
        "introduction": "Introduction} In the standard $\\Lambda$CDM paradigm dark matter halos evolve out of small fluctuations in the initial density field.  Since the structure of virialized objects may depend on their formation history, many attempts have been made to link the observed density distribution of dark matter halos with the initial spectrum of density fluctuation. High resolution N-body simulations indicate that density profiles of dark matter halos have a universal shape \\begin{eqnarray} \\label{NFW} \\rho(r)=\\frac{\\rho_s}{(cr/r_{vir})(1+cr/r_{vir})^2} \\end{eqnarray} whose dependence on the initial conditions is apparent only from the value of the concentration parameter $c$ (Navarro, Frenk \\& White 1997; henceforth NFW). While the NFW model has been usually successful in reproducing the observed mass profiles of large galaxies and galaxy clusters, it seems to be inconsistent with the rotation curves of the dwarf galaxies. In addition, several other N-body simulations \\cite{Mo,FM,Ri}, while generally obtaining similar results at large radii, found different density slopes in the inner regions of halos, thus raising some doubts about the accuracy of the NFW model.  The alternative approach to N-body simulations is provided by analytic modeling and in particular by self-similar accretion models \\cite{FG,Bert}. The choice of a particular set of initial conditions (the profile of an initial density perturbation is assumed to be a power law, $\\Delta M/M\\propto r^{-3\\epsilon}$) allows to recover an exact analytic solution for the final density distribution, that can be compared to the N-body results. Further, it can be shown that the appropriate value of $\\epsilon$ (and therefore the final density profile) is fixed by the spectrum of the initial density fluctuations, $\\epsilon=(n+3)/6$, where $n(k)=d\\log P(k)/dk$ \\cite{P}. Unfortunately, the very simplicity that is a strong point of the analytic models, can also lead to their downfall, if an important physical process has been left out.  Purely radial infall that results from the evolution of a spherically symmetric density perturbation, produces an extremely steep inner density profile, $\\delta=d\\log \\rho(r)/d\\log r\\leq -2$. As neither the observations no any of the  N-body simulations produced a similarly steep density distribution, attempts have been made to fix the self-similar models by giving collapsing halos some initial angular momentum \\cite{WZ,STW,Sub,Nus}. Including angular momentum allows to reduce the steepness of the inner slope, down to $-9\\epsilon/(3\\epsilon +1)=-(3n+9)/(n+5)$.  However, in general the solutions, being dependent on the amount of added angular momentum, still did not reproduce the universal shape given by equation (\\ref{NFW}).  In this paper we show that the discrepancy between the self-similar models and N-body simulations is naturally resolved by including the process of dynamical relaxation into the former. We prove that even tiny deviations from spherical symmetry are sufficient to make the dark matter particles velocity distribution isotropic in the inner region, provided the density slope is steeper than $-1$. Furthermore, we show that for $n<-1$ ($\\epsilon<1/6$) isotropic velocity distribution leads to the inner density slope $\\delta\\sim -1$. Since the spectral index, $n(k)$, spans the range from $-3$ (for smallest halos) to $-1$ (for galaxy clusters), our results imply that if the dark matter gravity is the only force to be considered,  then $\\delta\\sim -1$ in all existing virialized halos.  The paper is organized as follows. In \\S 2, we show that density profiles, which are steeper than $\\rho\\propto r^{-1}$, are inconsistent with unisotropic velocities at the central region. In \\S 3, we derive a density profile for a self-similar accretion model that includes dynamical relaxation process. In \\S 4, we compare our results with the N-body codes. We summarize our conclusions in \\S 5. ",
        "conclusions": "We have shown that the results of high-resolution N-body simulations of dark matter infall can be quite accurately reproduced by a simple self-similar accretion model. Moreover, the self-similar accretion model, being unlimited by computational constraints, may have an advantage over the N-body codes in reconstructing the density profile in the central region of dark matter halos. Naturally, it should be kept in mind that close to the center of a halo, a stellar disk or a supermassive black hole may strongly affect the density profile, making inadequate any model or code that rely exclusively on the dark matter gravity.  It is interesting to note that the two threshold values of $n$, which separate different types of self-similar solutions, have a special significance in the real Universe. Thus $n=-2$, which separates the solutions with finite and infinite central energy density, is also the largest possible value of $n$ on scales of existing virialized halos. The $n=1$ threshold, which separates the solutions with finite and infinite gravitational potential in the center, seems also to be the largest value of $n$ on all scales \\cite{Sp}.   \\begin{figure}[t] \\resizebox{\\columnwidth}{!} {\\includegraphics{f3.eps}} \\caption{\\label{fig3} The dark matter density versus radius. The dotted,  dashed-dotted and dashed lines correspond, respectively, to $n=-2$, $-2.5$ and $-2.75$. The solid line corresponds to the fitting function (Eq. \\ref{pow}) with $\\alpha=0.17$.} \\end{figure}"
    },
    "0608/astro-ph0608367_arXiv.txt": {
        "abstract": " ",
        "introduction": "The number of known planets has increased by more than an order of magnitude within the last decade - for the current list of extrasolar planets see e.g.\\ \\cite{SchneiderExoEnc}. Those additional planets created a diversity which made the task of defining \\emph{what is a planet} all the more difficult. Three distinct groups of properties are relevant when discussing the planethood criteria: physical characteristics, planetary dynamics, and cosmogony \\citep[for discussion see e.g.\\ ][]{Basri_PlanetDef}. \\paragraph{Orbital dynamics} While the majority of the community agrees that dynamically dominant bodies bound to stars or stellar remnants should be counted as planets (provided that they're not massive enough to support fusion), presently there's no consensus on how to designate potential unbound planetary-mass objects. Disagreement partly stems from a lack of understanding of the formation process(es) that could create unbound planetary-mass objects. The current guideline of the International Astronomical Union's (IAU) Working Group on Extrasolar Planets \\citep[][]{wgesp_definition} classifies such bodies not as `planets', but as `sub-brown dwarfs'. \\paragraph{Cosmogony} The present knowledge on the cosmogony of planets in general \\citep[e.g.][ for an overview]{ppiv_book} is far from being complete. On the other hand, for a foreseeable time there will be no way to observationally infer the history of exoplanets in detail. It is difficult to use cosmogony for any general planethood criterion which would be determinable from easily observed characteristics \\citep[e.g.\\ ][]{Basri_PlanetDef, SternLevinson}. Hence, WGESP guideline \\citep[][]{wgesp_definition} does not include cosmogony in the planet definition. \\paragraph{Physical properties} The `planetary mass' has its upper boundary around 13 M$_{\\mathrm\\jupiter}$ (M$_{\\mathrm\\jupiter}$ being Jupiter mass) as a deuterium burning limit \\citep[e.g][]{wgesp_definition, Basri_PlanetDef, SternLevinson}, but also as an effective upper end of the extrasolar planet mass-distribution \\citep[][]{marcy_butler_00, SchneiderDB}. The observational limit suggests that 13 M$_{\\mathrm\\jupiter}$ effective boundary is a consequence of an as-of-yet not fully understood formation process, and that the deuterium burning limit is either a side effect or a numerical coincidence. The emerging 1 M$_{\\mathrm\\jupiter}$ characteristic mass \\citep{SchneiderDB} also questions the importance of the thermonuclear reactions as the planetary delimiter. We should note that the recent trend of discovering Neptune-class planets \\citep[e.g.][]{Butler04, Bonfils05} could further reduce the observed characteristic mass.  Various concepts have been proposed for the lower planetary mass limit, but the official consensus is not yet reached. Concepts are based either on a particular physical property of the body in question, or on setting the arbitrary value for the required mass or radius of the body. An arbitrary boundary is based either on a historical heritage (e.\\,g.\\ discovery of Pluto) or on some numerical peculiarity of the Solar System (e.\\,g.\\ size of Ceres). In the past decade, the diversity of members of planetary family has been shown to grossly exceed examples from the Solar System. Hence, we feel that the general definition of a planet cannot include arbitrary values introduced from our own planetary neighborhood.\\\\* Continuing on the physical concepts for the low-mass planetary boundary, a few similar ideas are based on roundness of the body, i.\\,e. a body must be massive enough for its self-gravity to dominate over any material forces that might produce asymmetric shapes \\citep[e.g.][]{Basri_PlanetDef, SternLevinson}. The minimum size of such body somewhat depends on its composition, but one can calculate it to be a bit less than 500 km in diameter \\citep[][]{Basri_PlanetDef}. However, much smaller bodies can also achieve round shape via melting, or other form of asteroid differentiation; for review see \\cite{McCoy04, Taylor93}. \\cite{McCoy04} write that already a 10 km iron meteorite could reach 90\\% partial melting early in the formation process, and `Given that iron meteorite parent bodies were in the range of 10-100 km in radius, high degrees of partial melting must have been reasonably widespread, given the abundance of iron meteorite groups and ungrouped meteorites.' Additionally, from the round shape of km-sized rubble pile asteroids \\citep[e.g.\\ Dactyl, moon of Ida, c.f.][]{Davis96} we know that roundness can be acquired via violent kinetic events.\\\\* These facts hint that `roundness' could be connected to the cosmogony of the object rather than to its physical properties only, and as such would not be an ideally suited observable. Hence, roundness is not a proper concept at the bottom of the planetary mass scale.  The presence of other physical properties, such as a magnetic field, an atmosphere, or the existence of a satellite, has been proposed as the planethood criterion. \\citet{SternLevinson} show that any single such criterion is unable even to include all of the Solar system planets.  While the concept for the lower planetary mass limit is not yet a big problem for the extrasolar-planet surveys, it is the subject of a heated debate concerning the Solar system: Pluto's double classification (both as a minor planet and a planet) and the upcoming issue of what to do with its potentially even larger cousins from the trans-neptunian-object family \\citep[e.g.][]{Brown05}. It has been suggested \\citep[e.g.][]{marcy_butler_00} -- completely arbitrary and mostly for historic reasons -- to accept Pluto's mass as the lower mass delimiter. Others have argued the same for the mass of Ceres \\citep[e.g.][]{Basri_PlanetDef}. IAU has no official position yet, with WGESP opting for `The minimum mass/size required for an extrasolar object to be considered a planet should be the same as that used in our Solar System.'  In the next section we present our version of the low-mass planetary boundary. The third section summarizes the model of \\cite{PecnikWuchterl_05}, which we use in the fourth section to apply our planethood criterion to the large bodies in the Solar system. The last section summarizes the paper with conclusions and remarks. ",
        "conclusions": "\\paragraph{Summary} We provide a new lower planet boundary concept, requiring that a planet must be able to keep an atmosphere in vacuum. The previous attempts to use the atmosphere as a planetary delimiter had problems differentiating a dilute atmosphere from a vacuum. With the manifold paradigm \\citep[][]{PecnikWuchterl_05} we completely avoid the issue of weak atmospheres, and use the critical core mass, a crucial quantity of the planet formation, to discriminate a planet from a minor body. Because of such a general framework, the criterion stands a good chance to be relevant both for planets within and outside of the Solar system.  The proposed planethood concept is not sensitive to individual cosmogony, and requires only observables which are available from a transit event (mass, density, orbital distance) or can be extracted from the observations of the primary (primary mass, temperature). Only rough planetary classification can be performed with the current isothermal model. However, the proposed planethood concept will be quantitatively more accurate with a better physical description of a planet \\citep[for an example of a manifold with better constitutive relations see][]{Broeg05}.  \\paragraph{Outlook} The extrasolar planet surveys are presently yielding hot Neptunes - 10-15 $M_\\mathrm{\\earth}$ close-in planets. The advances in the observational technology and methods will further increase the sensitivity, and the trend of discovering ever-smaller planets \\citep[see e.g. Fig.\\ 1 in][]{santos05} is likely to continue. The upcoming wide-field transit surveys (e.g. COROT, Kepler) will extend the range of the known extrasolar planets to the Earth-mass regime. A clear, easily determinable, and quantitative concept for the lower planetary mass is urgently needed.  \\paragraph{Arbitrary planethood} Before the intricate process(es) of the planet formation and evolution are sufficiently understood, \\emph{every planethood concept}, including the one we champion, \\emph{will be completely arbitrary}. While being arbitrary, it should reflect in the best possible way what the astronomical community or maybe even humanity as a whole considers 'a planet'. We think that the ability to retain an atmosphere could be such a criterion. At the same time this criterion is broad enough that we believe it to be applicable to all the expected and unexpected planet discoveries in the future.  We discussed present ideas on the lower mass limit. We showed that roundness of the celestial body can be interwoven with cosmogony, and as such is not a good observable. Additionally, if the cosmogony would be accepted as the part of the planet definition, it would be impossible (in the foreseeable future) to decide if most of the exoplanetary candidates are indeed planets. This would be quite unpractical. Once we can infer details of the history of exoplanets through observation, \\emph{and} once we have in-depth knowledge of the general theory of planet formation, cosmogony \\emph{might} become useful component of the planet definition. On the other hand, an arbitrary fixed value for the object's mass as the limit (e.g. Pluto) might be irrelevant outside of the Solar system.  We base our \\emph{arbitrary} concept on a physical property (as opposed to a historical paradigm, or to a fixed number coming from Sol system), and, using the Occam's razor, apply the same reasoning that appears to work well for giant planets.  With the growing knowledge of the (exo)planetary properties, along with a healthy debate in peer-refereed publications, we think that the astronomical community stands a good chance of agreeing on a proper planethood criteria in the forseable future.  \\textbf{Acknowledgement} \\\\*We thank G\\\"{u}nther Wuchterl for the valuable discussions. His comments greatly improved this paper. We also thank anonymous referees whose comments helped to make this manuscript clearer. This research was in part supported by Max-Planck-Institut f\\\"{u}r extraterrestrische Physik, by National Foundation for Science, Higher Education and Technological Development of the Republic of Croatia, and by DLR project number 50-OW-0501.   \\def\\baselinestretch{1.0} \\begin{table}[p] \\label{Tab_Sol_Sys_CCM} $$ \\begin{tabular}{l|c|c|c|c|c|c|c} \\hline \\noalign{\\smallskip} object & $M_\\mathrm{primary}$ & orb. dist.& $\\varrho_\\mathrm{core}$& $T_\\mathrm{cs}$/[K]& $M_\\mathrm{crit}$/[kg]&$M_\\mathrm{obj}/M_\\mathrm{crit}$&TPC\\\\ &                      &   $a$/[AU]&          /[kg m$^{-3}$]&                    &                                   &                                &satisfied     \\\\ \\noalign{\\smallskip} \\hline \\hline \\noalign{\\smallskip} Mercury & $M_\\mathrm{\\odot}$ & 0.387 & 5427 & 440 & \\e{2.51}{22} & 13.2  & \\checkmark ?\\\\ Venus   & $M_\\mathrm{\\odot}$ & 0.723 & 5240 & 737 & \\e{7.08}{22} & 68.7  & \\checkmark \\\\ Earth   & $M_\\mathrm{\\odot}$ & 1.0   & 5515 & 287 & \\e{1.91}{22} & 312.8 & \\checkmark \\\\ Mars    & $M_\\mathrm{\\odot}$ & 1.523 & 3934 & 210 & \\e{1.58}{22} & 40.6  & \\checkmark \\\\ Jupiter$^{\\dagger}$ & $M_\\mathrm{\\odot}$ & 5.203 & 5515 & 153 & \\e{1.12}{22} & - & - \\\\ Saturn$^{\\dagger}$  & $M_\\mathrm{\\odot}$ & 9.537 & 5515 & 143 & \\e{1.17}{22} & - & - \\\\ Uranus$^{\\dagger}$  & $M_\\mathrm{\\odot}$ & 19.19 & 5515 &  68 & \\e{4.47}{21} & - & - \\\\ Neptune$^{\\dagger}$ & $M_\\mathrm{\\odot}$ & 30.07 & 5515 &  53 & \\e{3.16}{21} & - & - \\\\ Pluto   & $M_\\mathrm{\\odot}$ & 39.48 & 2000 &  44 & \\e{4.05}{21} & 2.5 & \\checkmark ? \\\\ Ceres   & $M_\\mathrm{\\odot}$ & 2.767 & 2050 & 167 & \\e{1.66}{22} & 0.057 & $\\times$ \\\\ Moon    & $M_\\mathrm{\\oplus}$&0.00254& 3344 & 250 & \\e{1.0}{22} & 7.35 & \\checkmark ?\\\\ Io & $M_\\mathrm{\\jupiter}$& 0.0028& 3528 & 130 & \\e{1.0}{21}  & 89.32  & \\checkmark \\\\ Europa&$M_\\mathrm{\\jupiter}$&0.004486&3013& 103& \\e{1.12}{21} & 42.78 & \\checkmark \\\\ Ganymede&$M_\\mathrm{\\jupiter}$&0.00715&1942&115& \\e{2.24}{21} & 66.14  & \\checkmark \\\\ Callisto&$M_\\mathrm{\\jupiter}$&0.01259&1834&115& \\e{3.16}{21} & 34.05 & \\checkmark ?\\\\ Mimas&$M_\\mathrm{\\saturn}$&0.00124& 1152 & 70  & \\e{1.1}{20} & 0.34 & $\\times$ \\\\ Enceladus&$M_\\mathrm{\\saturn}$&0.00159&1606&70 & \\e{2.5}{20} & 0.28 & $\\times$ \\\\ Tethys&$M_\\mathrm{\\saturn}$&0.00197& 956 & 86  & \\e{4.47}{20} & 1.40 & \\checkmark ? \\\\ Dione &$M_\\mathrm{\\saturn}$&0.00252& 1500 & 87  & \\e{7.08}{20} & 1.55 & \\checkmark ? \\\\ Rhea  &$M_\\mathrm{\\saturn}$&0.00352& 1240 & 76  & \\e{1.0}{21}  & 2.32 & \\checkmark ? \\\\ Titan &$M_\\mathrm{\\saturn}$&0.00817& 1880 & 94  & \\e{2.24}{21} & 60.1 & \\checkmark \\\\ Iapetus&$M_\\mathrm{\\saturn}$&0.02381& 1088 & 76  & \\e{3.55}{21} & 0.45 & $\\times$ ?\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} $$ \\def\\baselinestretch{1} \\caption[]{Terrestrial Planethood Criterion (TPC) for Solar System objects calculated with \\citet{PecnikWuchterl_05} planetary model. Celestial bodies are assumed to fulfill TPC if they can keep an equilibrium nitrogen atmosphere of temperature $T_\\mathrm{cs}$, in a vacuum, orbiting at a distance $a$ from a parent star of mass $M_\\mathrm{primary}$. Solar system giant planets are supercritical even for a solar composition gas (hence also for nitrogen), thus we only show the value for the nitrogen-supercritical object, at the present-day locations of the giant planets. See text for discussion. Symbols in the last column: `\\checkmark' - TPC fulfilled, `\\checkmark ?' - TPC fulfilled for the isothermal \\citep{PecnikWuchterl_05} model, unknown if TPC is satisfied for the model with more detailed constitutive relations, `$\\times$ ?' - TPC most likely not fulfilled even for the isothermal manifold, `$\\times$' - TPC not fulfilled. \\\\* $^{\\mathrm{\\dagger}}$ At the locations of gas giants we show a critical core mass for Earth's density\\\\* } \\end{table} \\def\\baselinestretch{1.66}"
    },
    "0608/astro-ph0608684_arXiv.txt": {
        "abstract": "We have analyzed a sample of nearby cool and warm infrared (IR) galaxies using photometric and structural parameters. The set of measures include far-infrared color ($C = \\log_{10}[S_{60\\mu m}/S_{100\\mu m}]$), total IR luminosity ($L_{TIR}$), radio surface brightness as well as radio, near-infrared, and optical sizes. In a given luminosity range cool and warm galaxies are considered as those sources that are found approximately $1 \\sigma$ below and above the mean color in the far-infrared $C - L_{TIR}$ diagram. We find that galaxy radio surface brightness is well correlated with color whereas size is less well correlated with color. Our analysis indicates that IR galaxies that are dominated by cool dust are large, massive spirals that are not strongly interacting or merging and presumably the ones with the least active star formation. Dust in these cool objects is less centrally concentrated than in the more typical luminous and ultra-luminous IR galaxies that are dominated by warm dust. Our study also shows that low luminosity early type unbarred and transitional spirals are responsible for the large scatter in the $C - L_{TIR}$ diagram. Among highly luminous galaxies, late type unbarred spirals are predominately warm, and early type unbarred and barred are systematically cooler. We highlight the significance of $C - L_{TIR}$ diagram in terms of local and high redshifts sub-millimeter galaxies. ",
        "introduction": "Observations by the {\\it Infrared Astronomical Satellite} ({\\it IRAS}) led to the identification  of very luminous and ultra-luminous infrared galaxies, commonly known as the LIRGs and ULIRGs, respectively. These galaxies have enormous far-infrared (FIR) luminosities, $10^{11} < L_{FIR}/L_{\\odot} < 10^{12}$ for LIRGs and $L_{FIR} > 10^{12} \\ L_{\\odot}$ for ULIRGs, and emit the bulk of their energy at the infrared wavelengths (Soifer et al. 1984, 1986, 1987; Sanders et al. 1988). Subsequent multi-wavelengths studies reveal that in the local universe ($z < 0.1$) ULIRGs have comparable total luminosities ($L_{TIR}$) and higher space densities than those of optical quasars (Sanders et al. 1988; Kim \\& Sanders 1998).  Understanding the nature and origin of energy sources in the LIRGs and ULIRGs has been the subject of much debate. Studies indicate that in the vast majority of these objects power comes from dust which is heated to various temperatures by various thermal and non-thermal processes such as ongoing (steady-state) star formation, intense starburst phase, and/or synchrotron radiation from the supernovae explosions (Lawrence et al. 1986; Genzel et al. 1998; Lutz et al. 1998). These systems may also be powered by active galactic nuclei (AGN; Lonsdale, Smith, \\& Lonsdale 1993). The energy distributions suggest that most of the observed FIR emission from luminous disk galaxies is due to the thermal radiation from warm dust heated by hot stars embedded in HII regions and molecular clouds heated directly by young OB stars; and cool dust from the general interstellar radiation field (Helou 1986; Presson \\& Helou 1987; Barvainis, Antonucci, \\& Coleman 1992). Dust could also be heated by shocks in the interstellar medium during the collision or interaction of galaxies (Harwit et al. 1987).  Dust heating mechanisms can be traced indirectly as each of these should correspond to particular spatial distributions of the FIR light. For example, the FIR emission originated by galaxies dominated by an AGN would appear compact and unresolved. On the other hand, starburst heated dust should have about the same scale size as the burst itself because young stars are well mixed with the gas from which stars are forming. Dust heated  by the non-ionizing photons from cold, population I stars, might be expected to follow the smoothed distribution of older giants in the galaxy (Zink et al. 2000). A determination of the total energy output such as luminosity, the amount of dust present in the interstellar medium and its spatial distribution within the galaxies as well as its relationship to other basic components, such as the atomic and molecular gas, the stars, etc, is essential to studies connecting galactic structure and nature of the IR energy sources (Carico et al. 1990; Sopp \\& Alexander 1991, 1992; Andreani \\& Franceschini 1992, 1996; Alton et al. 1998; Domingue et al. 1999; Siebenmorgen, Kr\\\"{u}gel \\& Chini 1999; Haas et al. 2000; Stickel et al. 2000; Trewhella et al. 2000; Zink et al. 2000).  The FIR color $C$, defined by the {\\it IRAS} $60$ -to- $100$ $\\mu m$ flux density ratio, an indicator of characteristic dust temperature, is a diagnostic of the typical heating conditions in the interstellar medium of a galaxy (Bothun, Lonsdale, \\& Rice 1989; Soifer \\& Neugebauer 1991).  It has been shown that low-redshift {\\it IRAS} galaxies exhibit positive correlation of $C$ with $L_{FIR}$ such that more luminous galaxies tend to be warmer as compared to less luminous galaxies (Soifer et al. 1987; Dale et al. 2001; see Fig.1 in this study). The color changes systematically over roughly three orders of magnitude in luminosity with large scatter. As a result there is a substantial number of highly luminous but cool galaxies as well as low luminosity yet warm/hot galaxies (Dale et al. 2001; Chapman et al. 2003). In this study we attempt to understand cool and warm IR luminous sources in terms of the geometric distribution of dust and galaxy morphology. In a given FIR luminosity range, these galaxies are considered in the context where cool galaxies are $1 \\sigma$ below and warm galaxies are $1 \\sigma$ above the mean color at a given luminosity. We are interested in differences between cool and warm galaxies in the same luminosity range that might suggest differences in the nature of the heating sources.  In an interesting study Young (1999) showed that star formation efficiency (SFE), defined by the ratio of $L_{TIR}$ to molecular hydrogen mass $M_{H_2}$, anti-correlated with disk (optical) sizes. The trend is prevalent in galaxies of various Hubble types and environments. It has been attributed to the shear present in disk galaxies. The molecular clouds in larger disks would experience increased turbulence which would reduce the efficiency of star formation (Young 1999). Galaxies in Young's sample span a broad range in FIR color comprising normal star forming galaxies (SFGs, $L_{FIR}/L_{\\odot} < 10^{11}$) as well as LIRGs and ULIRGs. In Young's study, however, it remains unexplored how cool and warm systems behave in the visual as well as in the longer wavelengths. Since the IR luminous galaxies contain very large amount of dust, the interpretation of any observation at optical or shorter infrared wavelengths gets complicated by extinction. Therefore, to probe the origin of large IR luminosity it is necessary to have the knowledge of the spatial extents of the emitting regions in observations unaffected by dust extinction.  The paper is organized as follows: we describe the data in $\\S2$ and parameter error estimation in $\\S$3. We present our results in $\\S$4 and discuss the implications of our results in $\\S$5. The conclusions are given in $\\S$6. ",
        "conclusions": "We have analyzed a sample of 213 nearby IR galaxies to study the correlations of galaxy FIR color vs. several photometric and physical parameters. The set of measures include total IR luminosity, radio surface brightness as well as radio, NIR, and optical sizes. Our objective is to understand cool and warm IR sources using various correlations. We find that galaxy radio surface brightness is well correlated with FIR color whereas size is less well correlated with color. A weak color-size correlation signals a significant uncertainty associated with the distance measurements of sample galaxies. It may also reflect methodological differences in estimating angular size in different wavelengths. We also find that late type galaxies (from Sc and beyond) of all morphology classes show less scattering than early types in the color-luminosity diagram. Our study shows that dust in the cool IR sources is probably less centrally concentrated than in the more typical luminous and ultra-luminous IR galaxies that are dominated by more centrally concentrated warm dust. We believe this result has significant implication in terms of FIR color-luminosity diagram: in absence of any structural information of high-$z$ luminous sources one can get rough estimate of the extent of the disks if one simply knows the FIR colors and luminosities of galaxies."
    },
    "0608/astro-ph0608151_arXiv.txt": {
        "abstract": "We present a wide-field (4.5 deg$^2$) photometric and spectroscopic survey of the Leo~I dwarf spheroidal (dSph) galaxy to explore its extended morphology and dynamics.  As in previous papers in this series, we take advantage of photometry in the $M$, $T_{2}$, and $DDO51$ filter system to select Leo~I red giant branch star candidates, and, so far, this selection technique has proven 100\\% reliable in selecting actual Leo~I members among more than 100 $M < 21.5$ Leo I giant candidates having previous or new Keck DEIMOS spectroscopy to a radius $>1.3$ times the limiting radius of the fitted, central King profile.  The two-dimensional distribution of all similarly-selected Leo~I giant candidates is well fitted by a central single-component King profile of limiting radius 13.3 arcmin, but many giant stars are found outside this newly derived King limiting radius.  The density profile thus shows a break at a major axis radial distance of $\\sim 10$ arcmin produced by an excess of stars at and beyond the King limiting radius (spectroscopically confirmed to be made of true Leo~ I members), and primarily along the major axis of the main body of the rather elongated satellite.  This spatial configuration, a rather flat velocity dispersion profile and an asymmetric radial velocity (RV) distribution among the Leo~I members at large radii together support a picture where Leo~I has been tidally disrupted on at least one, but at most two, perigalactic passages of a massive Local Group member. We demonstrate this hypothesis using mass-follows-light, $N$-body simulations of satellites in a Milky Way-like potential that reproduce the observed structural and dynamical properties of Leo~I remarkably well. These models include $\\sim3\\times10^{7}$ solar mass, tidally disrupting dSph systems on bound orbits with rather high eccentricity (0.93--0.96) and small perigalactica (10-15 kpc).  The simulations allow the first observationally constrained orbit for Leo~I {\\it without the measurement of its proper motion} and show that the observed RV distribution is more consistent with a two Milky Way orbit history for the satellite while ruling out a Leo~I orbit that includes a previous association with M31 within the last 10 Gyr. Given the overall success of tidally disrupting mass-follows-light satellite models to account for the observed properties of Leo~I, we conclude that there is no need to invoke an extended dark matter halo around the satellite (e.g., as one explanation of the velocity dispersion and radial profiles at large radii), and that an overall modest $M/L$ for the satellite is consistent with the available data. That a satellite on such a large (apogalacticon of $\\sim450$ kpc), long period ($P \\sim 6$ Gyr) orbit as Leo~I can experience tidal disruption suggests that similarly structured satellites with even smaller (eccentric) orbits will experience even greater tidally-induced mass loss rates. ",
        "introduction": "\\subsection{Motivations for a New Study of the Leo~I System}  Modern theories of the evolution of structure in the Universe that include cold dark matter (CDM) grow galaxies and clusters of galaxies and their dark halos through the accumulation of smaller subunits \\citep[e.g.,][]{White78,Davis85,Navarro96,Navarro97,Moore99}. But the notion that the Milky Way (MW) halo was built up by protracted infall of ``protogalactic fragments'' after an initial central collapse had already been established by \\citet{Searle78} based purely on stellar population arguments.  In CDM scenarios dwarf satellite galaxies represent the visible parts of (predominantly dark) subhalos; however, the number of dwarf galaxies discovered so far is several orders of magnitude less than the predicted number of subhalos made by CDM simulations \\citep{Kauffman93,Klypin99,Moore99}. This may be an indication that the majority of (especially smaller) subhalos have either not formed stars \\citep{Bullock01} or have been destroyed \\citep{Hayashi03}, and that the visible satellites of today represent the high mass end of the mass spectrum of DM subhalos \\citep{Stoehr02,Hayashi03} or those that were able to accrete substantial amounts of gas before reionization \\citep*{Bullock2000, Benson2002, Taylor2004}.  On the other hand, dwarf galaxies exhibit some properties that may be inconsistent with the expected properties of CDM subhalos. For example, the flat central density profiles of dwarf spheroidal (dSph) galaxies are at odds with the cuspy interiors predicted by CDM \\citep*[e.g.,][]{Navarro96, Navarro97, Moore1998}, although this might be reconciled by appealing to triaxial halos \\citep{Lokas2002, Navarro2004}.  Alternatively, warm dark matter (WDM) allows smaller central phase-space densities, and studies of dSph systems have provided important constraints on the properties of WDM species \\citep{Lin1983, Gerhard1992, Goerdt2006, Strigari2006}. In addition to the problem of the central concentrations of satellites, their apparent alignments around parent halos --- as has been observed in the Milky Way and argued to be related to ``dynamical families'' of satellites   \\citep*{Kunkel1979, Lynden-Bell1982, Majewski1994, FusiPecci1995, Lynden-Bell1995, Palma02} --- have also been used to question the viability of dwarf galaxies as putative subhalos \\citep{Kroupa05, Kang2005}.  On the other hand, such alignments have also been reported in high resolution N-body, hierarchical structure simulations, where subhalos accrete along filaments \\citep*{Libeskind2005, Zentner2005, Wang2005, Libeskind2007}.  Despite uncertainties over the precise connection of observed satellite galaxies to the prevailing theoretical CDM models, it has become observationally clear that satellite galaxies \\citep[e.g.,][] {Ibata95,Newberg02,Majewski03b,Yanny03,Crane03,Rocha-Pinto03} and even star clusters \\citep{Odenkirchen01,Odenkirchen03,Rockosi02, Lee04,Grillmair2006a,Belokurov2006c,Grillmair2006c} could be significant contributors to the {\\it luminous} halo of the MW and other galaxies. N-body simulations  \\citep{Oh95,Piatek95,Johnston95,Johnston99, Johnston02,Munoz2007} show how a dwarf galaxy can experience tidal disruption in its outer parts during close encounters with the central potential well of MW-like, parent systems.  Stars that escape the system will form tidal tails, like those observed in the Sagittarius (Sgr) system \\citep[e.g.,][]{Ibata95,Newberg02,Majewski03b,Belokurov2006d}. Such extratidal stars can create  ``breaks'' in the projected radial star-count profiles of the satellites as the density law transitions from a steeply-declining central density law for bound stars to a much more gradual decline at radii where unbound stars start to contribute significantly. Such ``breaks'' have been observed in several MW dwarf spheroidal (dSph) satellites \\citep{Eskridge88a,Eskridge88b,Irwin95,Majewski00a, Majewski03b,Palma03,Wilkinson04,Westfall06,Siegel05,Munoz06}. For example, the radial density profiles of nearly every one of the then-known dSphs of the Milky Way in the thorough \\citet[][hereafter IH95]{Irwin95} study (see their Figure 2) show an excess of stars with respect to the outer parts of their best-fitting model King profiles. In the particular case of Leo~I, because of its extreme distance, the number of dSph stars relative to the background level in the outermost regions was too low to judge conclusively the significance of its apparent break and limited the ability of IH95 to explore the Leo~I profile significantly past the King limiting radius with their photographic star counts.  Undoubtedly, a radial surface density profile with better signal-to-noise and spatial coverage will help verify whether this extremely distant Galactic companion shows the break profile trait seen in other MW dSphs.  Determining the structure of Leo I to large radii is a primary goal of the present study.  A second goal is to increase our understanding of the internal dynamics of Leo I.  If dSph galaxies are the visible counterparts of the largest DM lumps \\citep[e.g.][]{Stoehr02}, one might expect the process of tidal disruption to be inhibited or lessened in these massive subhalos. On the other hand, if dSphs can be proven to have tidal tails, the latter would place constraints on the dark matter components within these objects \\citep{Moore1996}.  The Sgr dSph provides a vivid example of a tidally disrupting system for which the degree of disruption, as measured by the nature of the tails, has been used to constrain its dark matter content \\citep{Law05}; on the other hand, the Sgr system is sometimes considered to be an exception to the norm for dSph galaxies \\citep{Mateo98, MateoARAA98}. That Sgr may {\\it not} be an exceptional case is suggested by studies of possible tidal disruption in other MW dSphs \\citep{Gould92,Kuhn96, Majewski00a,Majewski06,Gomez-Flechoso03,Palma03,Walcher03,Westfall06, Munoz05,Munoz06,Munoz2007}.  The effects of tidal disruption can be inferred not only by the existence of ``break'' in the density profiles of these other MW dSphs, but also, more recently, by the velocity characteristics of the dSphs at large radii. The past decade has seen substantial progress in measuring radial velocities for large numbers of stars in at least the more nearby MW dSph galaxies \\citep{Armandroff1995, Tolstoy2004, Wilkinson04, Munoz05, Munoz06, Walker2006}. These studies not only imply large central mass-to-light ratios in many of the satellites, but generally reveal relatively flat velocity dispersion profiles reaching into the break profile regions \\citep{Munoz05,Munoz06}, a phenomenon reproducible in N-body simulations of disrupting, {\\it mass-follows-light} satellites \\citep{Munoz07}.  However, a tidal disruption interpretation of these velocity dispersion profiles to large radii is not unique based on the extant data, because flat velocity dispersion profiles  may also be accommodated to arbitrary radii by extending the DM halos in which the luminous dSphs are embedded.  These ``equilibrium models'' imply mass-to-light ratios that increase with radius and boost the global ratio --- in some cases quite substantially \\citep[e.g.,][]{Kleyna02}.  Knowing whether mass follows light in dSphs or if their luminous components lie within extended dark halos is critical to establishing the regulatory mechanisms that have inhibited the formation of galaxies in all subhalos and the extent to which luminous satellites are vulnerable to disruption and populate the stellar halos of $L*$ galaxies.  Leo~I provides an interesting contrasting case to other dSphs in that (1) this particular satellite clearly has an unusual orbit that has protected it from but a few potential tidal encounters, and (2) unlike the apparently high $M/L$ systems (like Ursa Minor, Draco and Carina), the previously measured mass-to-light ratio of Leo~I is only $\\sim 6 M_{\\sun}/L_{\\sun, V}$, i.e. at the low end of the dSph $M/L$ scale.  Of course, it is of interest to know whether and how these two traits of this satellite may relate, and whether the properties of Leo~I, if explored more extensively, may lend new insights into the issues raised above. For example, as we shall show in this paper, new signatures found in the velocity distribution of Leo~I stars at large radii (namely, an asymmetry) may provide a way to break the degeneracy in interpretations of the velocity dispersion profile for at least some dSphs.  The heart of our study is a new photometric survey covering 4.5 deg$^{2}$ of the sky centered on Leo~I, and new Keck spectroscopy of 105 more Leo~I giant stars than available in the literature.  As with other contributions in this series, we adopt a technique based on multicolor Washington filter imaging, including the use of the $DDO51$ filter centered on the gravity-sensitive MgH+Mgb triplet spectral feature near 5150 \\AA\\ to identify giant stars associated with the Leo~I dSph.  We use bright K giants of Leo~I as a tracer population of the structure of the dSph because (1) they are the most easily detected type of stars over large areas with the use of a mid-size telescope, (2) the Washington+$DDO51$ separation technique has already been proven successful in the discrimination of metal-poor dSph giant stars from foreground metal-rich field dwarf stars and thereby to significantly increase the signal-to-noise of the faint, diffuse outer parts of Local Group dSphs \\citep[e.g.,][]{Majewski00a,Palma03,Munoz05,Munoz06}, and (3) these stars are accessible to spectroscopic membership and dynamical follow-up with the currently largest telescopes.  Because we are able to explore the Leo~I dSph to large radii with much better signal-to-noise than previous studies, we can rederive the structural parameters for Leo~I (\\S 4.1), which are important to assessing the $M/L$ using the traditional, \\citet{King62} methodology based on the internal dynamics of the system.  A newly derived $M/L$ for Leo~I in this way is discussed in \\S 4.2. We also explore the possible mass distribution for Leo I within the context of a tidal disruption scenario using $N$-body simulations of Leo I--like satellites orbiting a MW--like galaxy in \\S 6.2.  The unusually high radial velocity of Leo~I at its extreme distance naturally leads to interesting questions about its specific orbit.  There have been two studies that are in disagreement about potential orbits for Leo~I.  On one hand, \\citet{Byrd94} conclude that Leo~I was once loosely bound to M31 and now is in an unbound, hyperbolic orbit about the MW.  They estimate Leo~I's MW perigalactic distance to be $\\sim 70$ kpc and for this nearest approach to have occurred 2-4 Gyr ago. On the other hand, in an earlier study, Z89 concluded that Leo~I probably did not originate in the M31 system and that {\\it the most reasonable assumption is that Leo~I is bound to the MW}.  An unexpected benefit of our study is that we have been able to derive new constraints on the Leo~I orbit from detailed study of its structure and dynamics (\\S6.2), and we find reasonable agreement of these results to inferences about Leo~I's  orbit based on its star formation history.  Dwarf spheroidal galaxies, with other halo objects such as globular clusters and field stars, are also useful test particles for probing the large scale mass distribution and total mass of our Galaxy. Despite much work, the spatial extent and total mass of the MW remain among the more poorly established Galactic parameters. While traditionally thought to be a ``lesser sibling'' to the Andromeda galaxy (M31) in terms of mass, some recent work \\citep{Cote00, Evans00, Evans00b, Geehan06, Seigar07} suggests that, in fact, the MW edges out M31 as the most massive galaxy in the Local Group.  With the traditional analyses, the mass estimate is sensitive to the inclusion or exclusion of one particular object, Leo~I because of its combined unusually large Galactocentric radial velocity ($+177 \\pm 3$ km/s; \\citealp[hereafter Z89] {Zaritsky89}) and great distance ($257\\pm 8$ kpc, see \\S 1.2). Estimates of the Galactic mass by Z89 varied by a factor of 3--4 depending on whether Leo~I is considered as bound to the MW or not. While recent studies by \\citet{Wilkinson99}, \\citet{Sakamoto03} have decreased the dependency of MW mass estimates on including Leo~I by using Bayesian likelihood methods and larger samples of halo objects, Leo~I is still considered to be a determining factor for fine tuning the results.  \\subsection{Previous Photometric Studies of Leo~I}  The Leo~I dSph was discovered more than half a century ago by \\citet{Harrington50} during the first Palomar sky survey. Due to its extreme distance and its angular proximity to the 1st-magnitude foreground star Regulus, photometric studies of Leo~I have been difficult.  Not until CCD arrays were developed were the first color-magnitude diagrams (CMD) of Leo~I constructed. \\citet{Fox87}, \\citet{Reid91}, \\citet{Lee93}, and \\citet{Demers94} presented early ground-based CMDs.  Later, \\citet{Caputo99} and \\citet{Gallart99a} presented Leo~I CMDs based on data taken with the {\\it Hubble Space Telescope} ({\\it HST}) and they reach the main- sequence turnoff (MSTO) of the oldest ($> 10$ Gyr) Leo~I populations. These deep CMDs allowed detailed studies of the multiple stellar populations and the complex star formation history of Leo~I \\citep{Gallart99b}.  Leo~I was thought to be unique among the MW satellite dSphs for not having a conspicuous horizontal branch (HB) population until a 12\\arcmin $\\times$ 12\\arcmin\\ ground-based survey on Leo~I by \\citet{Held00} revealed an extended HB structure in its CMD.  More recently, \\citet{Held01} discovered more than 70 RR Lyrae variables with pulsational properties suggesting an intermediate Oosterhoff type similar to other dwarf galaxies in the Local Group \\citep{Siegel00,Cseresnjes01,Pritzl02}.  The existence of both extended blue HB population and RR Lyrae stars suggest that Leo~I is in fact similar to other local dSph galaxies in having a $> 10$ Gyr population, likely formed in the initial collapse of the system.  The extended star formation history of Leo~I, which includes this initial starburst, followed by a quiescent phase and a new burst of star formation activity starting $\\sim 7$ Gyr ago \\citep{Gallart99b} may be intimately tied to its orbital dynamics, since close interactions between galaxies are known to be triggers of star formation.  The distance to Leo~I has been derived using various photometric methods.  \\citet{Lee93} used the tip of red giant branch (TRGB) method to derive the Leo~I distance modulus $(m-M)_{0} = 22.18\\pm 0.11$, while \\citet{Demers94} used both the apparent magnitudes of red clump and of the carbon stars to estimate $(m-M)_{0} = 21.56\\pm 0.25$.  \\citet{Held01} used the mean magnitude of RR Lyrae variables in Leo~I and derived $(m-M)_{0} = 22.04\\pm 0.14$, and more recently, \\citet{Bellazzini04} provided a new estimate of $(m-M)_{0} = 22.02 \\pm 0.13$ using the TRGB method.  A weighted average of the distance modulus derived in the four studies gives $(m-M)_{0} = 22.05\\pm 0.07$, which converts to a distance of $257\\pm 8$ kpc.  Throughout this study, we adopt these values.  The following sections include a presentation of the data from our photometric survey of Leo~I (\\S 2), and a description of the photometric identification of the Leo~I giant star candidates (\\S 3). The two-dimensional distribution of Leo~I giant star candidates is discussed in \\S 4, and new morphological parameters for the dSph are derived. In \\S 5 we present new Keck Observatory spectroscopy of a subsample of our Leo~I giant candidates and in \\S 6 we discuss the implications of our results, making use of new $N$-body simulations of a tidally disrupting dSph satellite galaxy that appear to generate similar Leo~I properties to those we have observed.  Finally, a summary of our work and conclusions are in \\S 7. ",
        "conclusions": "We have photometrically surveyed a 4.5 degree$^2$ region centered on Leo~I in the $M$, $T_{2}$, $DDO51$ filter system in order to explore the extended morphology of this dwarf spheroidal galaxy, which is currently regarded as the most distant of the known Galactic satellites (unless the more distant Phoenix system is bound to the MW).  The photometric data were used to select Leo~I giant candidates based on the two criteria: (1) the gravity sensitive ($M-DDO51$, $M-T_2$) diagram, which separates distant giant stars from contaminating, foreground, metal-rich disk dwarfs, and (2) CMD positions commensurate with the temperature--apparent magnitude combination of stars on the Leo~I giant branch.  The background level of our ``Leo~I giant star'' sample is determined to be small, and a 100\\% reliability in the identification of bona fide Leo I giants is found via testing with a total of 133 stars in the Leo I field with previously published or new Keck spectroscopy.  We derive a new set of Leo~I structural parameters by fitting a single-component King profile to the Leo~I giant candidates. Coupling this to the central velocity dispersion we have measured (\\S 5.5), we use core-fitting techniques to derive a mass for the Leo I system of $3.5 \\times 10^{7} M_{\\sun}$ and a total $(M/L)_{V}$ of 4.6$M_{\\sun}/L_{\\sun ,V}$,\\footnote{These values are 14\\% smaller than those in \\S4.2, where the M98 central velocity dispersion was used to derive the mass.} values not too dissimilar from previous study using the same technique (M98).  The two-dimensional distribution of Leo~I giant candidates shows many giant stars outside the derived King limiting radius.  These are primarily along the major axis and spectroscopy of a subsample of these ``extratidal'' stars shows they are actually associated with the dSph. This population of stars shows up as a break at a major axis radial distance of $\\sim 10$ arcmin in the radial density profile. Our new Keck spectroscopy confirms this excess of stars beyond the King limiting radius to be made of true Leo~I members.  Our Leo~I velocity dispersion profile shows a flat and then rising trend at large radii.  We also find that the Leo~I RV distribution, particularly for stars at large angular separations (which in our data are primarily to the west side of Leo I), includes a population of stars with a broad and skewed, asymmetrical distribution toward positive RVs overlapping a second, very much colder population at the nominal Leo~I mean velocity.  We interpret these features as support for a picture wherein Leo~I has been tidally disrupted on at least one, but at most two, perigalactic passages by a massive Local Group member (most likely the MW for both), and find these phenomena naturally produced in mass-follows-light $N$-body simulations of a tidally disrupting Leo~I analog in a Milky Way-like potential.  The best-matching $N$-body simulations to both the observed structural morphology and velocity distribution of Leo I are those where the tidally disrupting satellite is on a rather high eccentricity (0.93--0.96), small perigalacticon (10-15 kpc), bound orbit around the Milky Way, and has a present total $M/L \\sim 3-4.5 (M/L)_{\\sun}$.  These best fitting model masses are 58--85\\% the $M/L$ derived from the central velocity dispersion and core fitting, but it is not unreasonable to presume that the latter method yields somewhat inflated masses because of systematic increases in the true dispersion due to astrophysical processes such as the presence of binaries and the atmospheric jitter common to giant star atmospheres.  Thus, the likely $M/L$ of the satellite is rather modest and not unlike those of other elliptical systems of approximately similar mass scale that are typically regarded as low in, or devoid of, DM (dE galaxies and globular clusters).  Because of the rather close match between our model results and observations, and because disrupting satellites on highly radial orbits appear to yield great discriminatory power in this regard, we have been able to constrain the likely orbit of Leo I {\\it without the measurement of its proper motion}.  The orientation of the current satellite orbital plane can be fixed by matching model tidal tails to the predominant direction of the observed break population (i.e., more or less along the major axis), whereas we find that the {\\it direction of angular momentum} in this orbital plane is well constrained by the sense of the RV asymmetry we have observed to the west side of the satellite.  Our models demonstrate that a positive RV asymmetry in the models is produced by {\\it trailing} tidal debris for the receding satellite.  Thus, we predict that an opposite RV asymmetry will be found on the east side of Leo~I from leading tidal debris.  Such a result would provide an important verification of the tidal disruption model we have put forward in this paper.  Our observed Leo~I RV distribution is most consistent with a two MW orbit history for Leo~I, with both orbits around the MW, however we cannot yet definitively rule out a one orbit scenario.  However, whether Leo~I is bound to the MW or on an unbound, hyperbolic orbit around the MW, our results seem to rule out a Leo~I orbit that includes a previous association with M31 within the last 10 Gyr, in contradistinction to the \\citet{Byrd94} scenario of a relatively recent origin of Leo~I from M31.  Leo I has long played a ``spoiler'' role in setting the mass of the MW because of its huge $rv^2$ lever arm in Jean's equation-based determinations using tracer particles of unknown proper motion (i.e. unknown orbit size and shape.  The large implied escape velocity at Leo~I's distance implies large MW masses if Leo~I is bound to our galaxy (Z89, \\citealt{Wilkinson99}), although its influence on the mass determination is lessened when more complete samples of objects with complete phase space data are employed \\citep[e.g.,][]{Sakamoto03}.  Our observed RV distribution for Leo~I slightly favors a two perigalacticon pass, bound MW orbit for Leo~I, thereby suggesting that higher mass estimates for the MW may be more correct.  In general, our tidally disrupting mass-follows-light satellite models provide a quite satisfactory match to the observed properties of Leo~I, but a few details --- namely the apparent double break in the density profile and the density at which the inner of these breaks occur --- we have yet to account for in these initial modeling efforts.  Nevertheless, we contend that the scenario of a tidally disrupting, low $M/L$ system on a highly radial orbit provides a rather complete explanation for the observed properties of Leo~I.  While some properties of Leo~I and other dSphs (e.g., flat velocity dispersion profiles and break populations) have also been explained by postulating that these systems are embedded in extended DM halos, such an explanation in the case of Leo~I appears less compelling in that it cannot account for as many of the observed properties of the system (e.g., the asymmetry in the RV distribution and increasing ellipticity with radius).  In contrast, tidal tails by now provide a well-established observational paradigm for dSph satellites, with the Sagittarius dSph the most vivid example. We contend that the tidal disruption of Sagittarius is not unique and that Leo~I may be another example of the phenomenon, albeit at a much lower mass loss rate commensurate with a satellite on such an extreme orbit.  Tidal disruption observed among both the closest and farthest of the Milky Way satellites suggests that this process may be ubiquitous, and that similarly structured satellites with elliptical orbits and distances between these two examples might also be expected to be experiencing tidally induced mass loss, likely with an intermediary range of mass loss rates."
    },
    "0608/astro-ph0608198_arXiv.txt": {
        "abstract": "Nulling interferometry is an astronomical technique that combines equal wavefronts to achieve a deep rejection ratio of an on-axis star, and that could permit to detect Earth-like planets in the mid-infrared band [5 -- 20 $\\mu$m]. Similarly to what is done in the near-infrared, high frequencies spatial filtering of the incoming beams can be achieved using single-mode waveguides operating in the mid-infrared. An appreciable reduction of the instrumental complexity is also possible using integrated optics (IO) devices in this spectral range. The relative lack of single-mode guided optics in the mid-infrared has motivated the present technological study to demonstrate the feasibility of dielectric waveguides functioning at longer wavelengths. We propose to use selenide and telluride components to pursue the development of more complex IO functions. ",
        "introduction": "In the last ten years, the photonics field has provided a number of valuable technological solutions to astronomers to improve the performance of the observatories instruments. In particular, optical fibers and integrated optics (IO) devices have already been part \\cite{Foresto, Berger}, or will be part in the future \\cite{Malbet, Perrin} of interferometric facilities. As a result, this context has raised the question of extending the use of photonics devices to the Darwin Infrared Space Interferometer of the {\\it European Space Agency} (ESA) \\cite{Fridlund}.\\\\ \\indent This mission, based on the implementation of nulling interferometry \\cite{Bracewell}, is devoted to the search and characterization of terrestrial exoplanets in the mid-infrared range thanks to its combined high contrast % and high spatial resolution (long-baseline interferometry) features \\cite{Angel}. The principle of a two aperture nulling interferometer is to recombine artificially $\\pi$ phase-shifted wavefronts from an on-axis star in order to achieve their mutual cancellation. The signal from an off-axis object -- like a planet orbiting the star -- is transmitted through the instrument by tuning the interferometer baseline so that a total phase shift of 2$\\pi$ is achieved. A key requirement of nulling interferometry is % equal wavefronts. Reducing the wavefront errors (e.g. pointing errors, optical defects due to polishing and coatings) can be obtained with single-mode waveguides as modal filters \\cite{Mennesson}. Using single-mode Integrated Optics (IO) components merges the functions of modal filtering and beam combination, providing a stable and compact instrument with regards to external mechanical and thermal constraints. To date, the implementation of modal filtering with single-mode waveguides has become the major option for space-based nulling interferometry.\\\\ \\indent The lack of single-mode waveguides for the mid-infrared in the photonics field has pushed % advanced technological studies to develop such devices \\cite{Borde, Wehmeier, Wallner}. In this context, we have initiated, under ESA contract, the {\\it IODA} program ({\\it I}ntegrated {\\it O}ptics for {\\it DA}rwin) to investigate the feasibility of a single-mode Integrated Optics concept operating over the full 5 -- 20 $\\mu$m spectral range. Our program has focused on the feasibility of an IO beam combiner based on either conductive waveguides \\cite{Labadie} or on dielectric materials.\\\\ \\indent In this paper, we present results obtained using selenide and telluride materials to manufacture planar waveguides for the mid-infrared range, the first step toward further efforts for the development of channel waveguides. We show that the modal behavior of the various structures can be experimentally investigated and assessed with the {\\it m}-lines method, known in the visible and near-infrared range, and which we have successfully transferred  to the mid-infrared range. To our knowledge, this paper presents the first study of the modal properties of selenide and telluride thick films at longer wavelengths. ",
        "conclusions": "\\noindent This paper has discussed the optical characterization of chalcogenide-glass-based thick films manufactured for the purpose of future mid-IR integrated optics development. We have fully developed an {\\it m}-lines setup operating at $\\lambda$=10.6 $\\mu$m based on prism coupling theory. This setup was used to characterize the modal behavior of the first mid-infrared planar waveguides by observing mode lines. In this study, two different compositions were successfully tested at 10 $\\mu$m with the first observation of three mode lines for TE and TM polarizations. A simple numerical analysis based on a best fit between the experimental and theoretical mode indices distributions has permitted us to derive the refractive index and thickness of the different samples. Our results have underlined the possibility of producing full-selenide multimode planar waveguides with a sufficient film thickness. The near and mid-IR results are consistent, since similar thicknesses are obtained in the two spectral ranges. Extending the transmission range of the components to 20 $\\mu$m may be possible with the telluride solution.\\\\ \\indent Following these results, a complementary technological study has been conducted to investigate the feasibility of a first full-telluride single-mode planar waveguide. This study has produced positive results. As a consequence, the question of etching of a channel waveguide to be used as a modal filter for nulling interferometry has been investigated straightforward and the results recently obtained will be part of a forthcoming paper."
    },
    "0608/astro-ph0608221_arXiv.txt": {
        "abstract": "{Extreme ultraviolet resonance lines of neutral and ionised helium observed in prominences are difficult to interpret as the prominence plasma is optically thick at these wavelengths. If mass motions are taking place, as is the case in active and eruptive prominences, the diagnostic is even more complex.} {We aim at {studying the effect of radial motions on the spectrum emitted by moving prominences in the helium resonance lines and at facilitating the interpretation of observations}, in order to improve our understanding of these dynamic structures.} {We develop our non-local thermodynamic equilibrium radiative transfer code formerly used for the study of quiescent prominences. The new numerical code is now able to solve the statistical equilibrium and radiative transfer equations in the non-static case by using velocity-dependent boundary conditions for the solution of the radiative transfer problem. This first study investigates the effects of different physical conditions (temperature, pressure, geometrical thickness) on the emergent helium radiation.} {The motion of the prominence plasma induces a Doppler dimming effect on the resonance lines of \\ion{He}{i} and \\ion{He}{ii}. The velocity effects are particularly important for the \\ion{He}{ii} $\\lambda$\\,304~\\AA\\ line as it is mostly formed by resonant diffusion of incident radiation under prominence conditions. The \\ion{He}{i} resonance lines at 584 and 537~\\AA\\ also show some sensitivity to the motion of the plasma, all the more when thermal emission is not too important in these lines. We also show that it is necessary to consider partial redistribution in frequency for the scattering of the incident radiation.} {This set of helium lines offers strong diagnostic possibilities that can be exploited with the SOHO spectrometers and with the EIS spectrometer on board the Hinode satellite. The addition of other helium lines and of lines from other elements (in particular hydrogen) in the diagnostics will further enhance the strength of the method.} ",
        "introduction": "It is still not clear how solar prominences reach a state of equilibrium. However it is known that during the formation stage, significant plasma motions are taking place. Similarly, mass motions are important during the disappearance of prominences. Between these two stages of the life of prominences, one may also observe periods of activity with internal plasma motions \\citep[see for example][]{ofmanetal98,ktp03}. In order to make an accurate assessment of the energy budget in solar prominences, detailed observations of a number of lines formed at different temperatures are crucial. {Among them, the helium lines are important, as they are strong and complement the hydrogen lines to reveal the physical conditions of the prominence plasma in different regions of the structure.} However, extreme ultraviolet resonance lines of neutral and ionised helium are difficult to interpret as the prominence plasma is optically thick at these wavelengths. If mass motions are taking place, the diagnostic is even more complex. {Furthermore , h}elium is the most abundant element after hydrogen, and therefore significantly contributes to the energy budget, {for instance by absorbing coronal radiation at wavelength below the \\ion{He}{i} ionisation threshold ($\\lambda < 504$~\\AA) -- see, e.g., \\citet{vaetal99,ah05}. It does not, however, significantly contribute to the radiation losses \\citep{kp91}. In all cases, it appears that in order to build realistic prominence models, detailed non-LTE radiative transfer modelling is necessary.}  We aim at facilitating the interpretation of observations in the helium resonance lines in moving prominences, in order to improve our understanding of these dynamic structures. {This paper is thus devoted to the study of the effect of radial motions in prominences on the helium resonance lines.} We present the first computations of the helium line profiles emitted by a moving prominence. The prominence is modelled as a plane-parallel slab standing vertically above the solar surface and moving upward as a solid body. The helium spectrum is computed with a non local thermodynamic equilibrium radiative transfer code. The modelling aspects are presented in \\S~\\ref{modelling}. The effect of Doppler dimming is investigated in \\S~\\ref{results} for the EUV resonance lines of \\ion{He}{i} at 584~\\AA\\ and 537~\\AA\\ and \\ion{He}{ii} at 304~\\AA. We focus on the line profile properties and the resulting integrated intensities {normalised to their respective values when the prominence is at rest}. We also study the effect of frequency redistribution in the line formation mechanisms. We discuss the results in \\S~\\ref{discuss}. ",
        "conclusions": "\\label{discuss}  {In this paper we have studied the effect of radial motions in prominences on the helium resonance lines \\ion{He}{i} 584~\\AA, \\ion{He}{i} 537~\\AA, and \\ion{He}{ii} 304~\\AA.}  The consideration of partial redistribution in frequency (PRD) is necessary to compute the profiles of the helium resonance lines emitted by moving material in prominences. This moving material could correspond to upwards/downwards mass flows in active prominences or ejected plasma in erupting prominences.  {The study of the variation of the relative intensities of the helium resonance lines with the radial velocity of the plasma indicates that th}e Doppler dimming effects are essentially present when the relative contribution of thermal emission compared to the scattering of incident radiation in the lines studied is low. At velocities larger than 150~km\\,s$^{-1}$, most prominences show a strong sensitivity to the motion of the plasma in the \\ion{He}{ii} 304~\\AA\\ line, and to a lesser extent, in the \\ion{He}{i} 584~\\AA\\ and 537~\\AA\\ lines. At lower velocities, all \\ion{He}{i} and \\ion{He}{ii} EUV resonance lines are sensitive to Doppler dimming.  It has been shown in this paper that most of the information on the radial velocity can be obtained through the variation of the {relative} integrated intensity, \\textit{ie.} by measuring the amount of Doppler dimming on the EUV resonance lines of helium. {Although this measurement is difficult to perform (mainly because one has to be observing at the right place and at the right time), the current space-based instruments such as EIT, SUMER, and CDS on SOHO, as well as TRACE and Hinode, are able to obtain a temporal series of intensity measurements covering an entire prominence eruption, starting before the radial motions take place. In this case it is possible to measure the dimming in intensity, and to compare it with our predicted relative intensity curves.}  {In a future study, a} more comprehensive diagnostic of the prominence plasma will be done through a comparison between the observed and our computed line profiles. Detectors with a spectral pixel size around 0.05 \\AA\\ or less (e.g. the SUMER instrument on SOHO) are ideal as they allow line asymmetries to be observed. A lower spectral resolution still allows good comparison with theoretical line profiles, however it makes the task of finding a unique model that fits the observed optically thick profiles more difficult. The line intensity of \\ion{He}{ii} 304 may be used without knowing {the line} profile for measuring the radial speed of the erupting prominence. In combination with the study of the apparent motion of the prominence material brought by {imagers on SOHO, Hinode, and STEREO,} the full velocity vector will be inferred.  We have developed a powerful diagnostic tool that will significantly improve our understanding of the thermodynamical parameters of active and eruptive prominences. Using this set of three helium lines (\\ion{He}{i} $\\lambda\\lambda$\\,584 and 537~\\AA\\ and \\ion{He}{ii} $\\lambda$\\,304~\\AA) allows to constrain the models in order to reproduce the observations. Adding more lines such as hydrogen lines will yield even better constraints. We will better understand the differences between \\textit{typical} erupting prominences showing the same morphology in \\ion{He}{ii} and in H$\\alpha$ \\citep{wfb86}, and \\textit{abnormal} erupting prominences showing extensive emission in \\ion{He}{ii} but nothing in H$\\alpha$ \\citep{wangetal98}. This will be the subject of a future work."
    },
    "0608/astro-ph0608017_arXiv.txt": {
        "abstract": "} \\nc{\\eab}{  Recently, two issues concerning the three-year Wilkinson Microwave Anisotropy Probe (WMAP) likelihood code were pointed out. On large angular scales ($l \\lesssim 30$), a sub-optimal likelihood approximation resulted in a small power excess. On small angular scales ($l \\gtrsim 300$), over-subtraction of unresolved point sources produced a small power deficit.  For a minimal six-parameter cosmological model, these two effects conspired to decrease the value of $n_s$ by $\\sim 0.7 \\sigma$.  In this paper, we study the change in preferred parameter ranges for extended cosmological models, including running of $n_s$, massive neutrinos, curvature, and the equation of state for dark energy. We also include large-scale structure and supernova data in our analysis. We find that the parameter ranges for $\\alpha_s$, $\\Omega_k$ and $w$ are not much altered by the modified analysis. For massive neutrinos the upper limit on the sum of the neutrino masses decreases from $M_\\nu < 1.90$eV to $M_\\nu < 1.57$eV when using the modified WMAP code and WMAP data only. We also find that the shift of $n_s$ to higher values is quite robust to these extensions of the minimal cosmological model. ",
        "introduction": "The temperature fluctuations in the cosmic microwave background (CMB) radiation have proved to be the single most important cosmological observable we have today, and the high-precision full-sky maps provided by the Wilkinson Microwave Anisotropy Probe (WMAP) play a very important role in the determination of cosmological parameters and preferred cosmological models \\cite{spergel:2006}. One of the most important conclusions both from the WMAP data alone, and from joint analyses including other cosmological observables, is that a simple six-parameter flat $\\Lambda$CDM model fits the data very well, and that extended models with additional free parameters do not improve the fit significantly.  Because of the great impact of WMAP data, the 3-year data analysis from the WMAP team has been subject to exhaustive cross-checking. In particular, in refs.\\ \\cite{eriksen:2006} and \\cite{huffenberger:2006} two noticeable issues with the likelihood code as first presented by the WMAP team were pointed out. First, the likelihood approximation used between $13 \\le l \\lesssim 30$ appears to be inadequate, effectively resulting in a $\\sim 5\\%$ power excess in this range compared to an exact treatment. Second, the amplitude for the unresolved point source spectrum used by the WMAP team was found to over-subtract the actual contribution in the data, leading to a power deficit at high $l$'s.  In \\cite{huffenberger:2006}, the effect of these discrepancies were studied for a minimal six-parameter cosmological model. This was done both for WMAP data only, and with additional CMB data from small-scale experiments. Their main finding was an increase in $n_s$, lowering the significance of $n_s \\neq 1$ from $\\sim 2.7\\sigma$ to $\\sim 2.0 \\sigma$.  In this paper we consider the effect on extended cosmological models. We study how the modified WMAP likelihood affects the preferred ranges of the running of the scalar spectral index, the cosmological neutrino mass limits, spatial curvature, and the equation of state for dark energy. We have also taken into account large-scale structure (LSS) and type 1a supernovae (SNIa) data sets, to see whether the shifts in preferred parameter ranges survive a more thorough cosmological analysis. Further, we consider whether the shift of $n_s$ to larger values is robust to changes in cosmological models and data sets.  In the next section, we review both the methods and data we use. In Section \\ref{sec:results}, we report and comment upon our results, before summarizing and concluding in Section \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  In refs.\\ \\cite{eriksen:2006} and \\cite{huffenberger:2006}, two modifications to the three-year WMAP likelihood were presented. Relative to the power spectrum presented by Hinshaw \\cite{hinshaw:2006}, they found a power deficit for low values of $l$ due to inaccurate likelihood approximation, and a small power excess for high values of $l$ due to over-subtracted unresolved point sources.  The impact on the inferred cosmological parameter intervals from these corrections in a minimal six-parameter model was studied in ref.\\ \\cite{huffenberger:2006} using CMB data only. Their single most important result was an increase in the preferred value of $n_s$, lowering the significance of $n_s \\neq -1$ from $\\sim 2.7 \\sigma$ to $\\sim 2.0 \\sigma$. In the present paper, we have extended that analysis to also account for cosmological models including a non-zero running of $n_s$, massive neutrinos, curvature and $w \\neq -1$. We have also added LSS and SNIa data sets to our analysis to see if the parameter shifts induced by the modified WMAP analysis survive when adding more data sets.  We found that the shift of $n_s$ to larger values survives when we add LSS and SNIa data. However, when we apply all data sets in the full 11-parameter model, $n_s$ is not affected much by the modified WMAP analysis anymore. This is mainly due to the extra freedom with $\\alpha_s$.  For the extended models, we found that the preferred values of $\\alpha_s$, $\\Omega_k$ and $w$ are not significantly affected by the modified analysis. When including massive neutrinos we found that the upper limit on $M_\\nu$ when using WMAP data alone was reduced from $M_\\nu < 1.90$eV to $M_\\nu < 1.57$eV. A similar improvement in the $M_\\nu$ limit could not be observed when adding LSS and SN1a data, since the higher preferred value of $n_s$ will allow for larger neutrino masses in the LSS power spectrum.  Since the initial publication of the two re-analysis papers, refs.\\ \\cite{eriksen:2006} and \\cite{huffenberger:2006}, and the present paper, the WMAP team has released a new version of their likelihood code\\footnote{Version v2p2p1} that implements the suggested low-$l$ correction and a revised point-source correction.  Using this updated likelihood code, we find $n_s = 0.959 \\pm 0.016$ for the six-parameter model and the WMAP data only. Including massive neutrinos, this code gives an upper bound of $M_\\nu<1.75\\textrm{eV}$ from WMAP data only. The difference is due to the point source amplitude and corresponding error adopted by the WMAP team, which do not match perfectly that of ref.\\ \\cite{huffenberger:2006}. Unfortunately, full details on the WMAP approach are not currently available, and final assessment of this issue must therefore await the release of the revised WMAP3 papers.  To conclude, the modified analysis does not strengthen the case for non-standard cosmological parameters, and the standard flat $\\Lambda$CDM model still provides an excellent fit to data."
    },
    "0608/astro-ph0608331_arXiv.txt": {
        "abstract": "We present a detailed study of the small frequency separations as diagnostics of the mass of the convective core and evolutionary stage of solar-type stars. We demonstrate how the small separations can be combined to provide sensitive tests for the presence of convective overshoot at the edge of the core. These studies are focused on low degree oscillation modes, the only modes expected to be detected in distant stars. Using simulated data with realistic errors, we find that the mass of the convective core can be estimated to within 5\\% if the total stellar mass is known. Systematic errors arising due to uncertainty in the mass could be up to 20\\%. The evolutionary stage of the star, determined in terms of the central hydrogen abundance using our proposed technique, however, is much less sensitive to the mass estimate. ",
        "introduction": " ",
        "conclusions": "We have presented the results of our study to relate the small frequency separations to the central layers of solar-type stars.  We have shown that the average small separation, \\DZT, can be used as a measure of the mass of the convective core in stars of mass $1.2$--$1.4\\msun$. In addition, we also propose two combinations, \\THETA\\ and \\ETA, of the small separations of different pairs of degrees as new diagnostics of the mass of the convective core, \\mc\\ and the stellar age, in terms of \\xc. We have tested the applicability of our technique to data with errors by estimating the random and systematic errors in \\mc\\ and \\xc\\ through Monte Carlo simulations.  The motivation behind the use of such combinations of \\Dzt\\ and \\Dot\\ stems from the fact that while they both encode slightly different internal phase shifts \\citep{rv00}, their combination might be sensitive to the sharp discontinuity at the boundary of the core. A more thorough theoretical analysis of these diagnostic quantities is required to understand their features better. It turns out that the diagnostic value of \\ETA\\ and \\THETA\\ can be maximised with an appropriate choice of the frequency ranges indicated here. A different choice of the frequency range is not only feasible, but might, in fact, be necessary depending on the availability of data. A detailed error analysis using calibration curves with the available observed frequencies to be used for the averaging would, however, be necessary in order to apply the technique to real data.  It turns out that all the calibration curves depend on convective overshoot in the core, and therefore we would need an independent estimate of the extent of overshoot in order to determine \\mc\\ and \\xc\\ to better than $40\\%$.  We find that the new diagnostic \\ETA\\ itself is, however, a good indicator of the presence of overshoot.  The values of \\ETA\\ are much larger in the presence of overshoot, irrespective of the mass.  The mass of the convective core changes as the star evolves.  For solar-type stars with mass less than $2\\msun$, it increases with age during the first part of evolution on the main sequence before shrinking down slowly in the later part. The values of the small separations, however, decrease monotonically with age, being more sensitive to the general evolution than to the size of the core. We attempt to account for the general change in the mean density with age by scaling the mass of the convective core with the large separation. We find that the average small separation, or the new diagnostic \\ETA, varies almost monotonically with the scaled core mass, \\scmc.  This scaling enables us to calibrate \\DZT\\ or \\ETA\\ against the core mass. It turns out that the random error in determining \\mc\\ through \\DZT\\ is smaller than that through \\ETA. On the other hand, the systematic errors due to inaccuracy in the total mass of the target star in the calibration curves can be as large as up to $20\\%$ for typical mass errors of $10\\%$.  We have also shown that the new diagnostic quantities, \\ETA\\ and \\THETA\\ may be used to estimate the age of a main sequence star. We find that they serve as nearly complementary indicators of the central hydrogen abundance -- while \\THETA\\ is a fairly good indicator of \\xc\\ during the early phase of evolution, \\ETA\\ can be used at more evolved stages. Both of these quantities, as diagnostics of \\xc, are fairly insensitive to mass, which make them better indicators of the evolutionary stage of a star than the small separations themselves. The actual age of the star in terms of years would, of course, still depend heavily on the stellar mass.  The random errors due to errors in frequency actually dominate the systematic errors due to uncertainty in mass. \\xc\\ can be determined to $5\\%$ for typical frequency errors of $\\sim 0.1\\mu$Hz. We note that \\ETA\\ is more sensitive to both \\mc\\ and \\xc\\ when overshoot is included in the models.  While the age of the star determined through \\ETA\\ or \\THETA\\ is quite independent of stellar mass, the determination of the mass of the convective core depends on the total mass of the star. Further, both of these techniques are sensitive to the extent of convective overshoot in the core. We have, in this work, provided a method to test the extent of overshoot through the diagnostic \\ETA\\ itself. The estimate of the stellar mass, however, needs to be obtained independently. For the fortunate cases of binary stars \\citep[e.g., as in $\\alpha$~Cen,][]{pnn99}, the mass is usually known to a high accuracy from dynamic considerations. Even for single stars, a precise location on the \\hrd\\ and a good estimate of the metallicity often helps to constrain the mass. In any case, a detailed correlation analysis between the age, the stellar mass and the extent of overshoot will be needed to correctly estimate the uncertainty in the results obtained from this technique."
    },
    "0608/gr-qc0608140_arXiv.txt": {
        "abstract": "The inspirals of ``small'' ($1 - 100\\,M_\\odot$) compact bodies through highly relativistic orbits of massive (several $\\times 10^5\\,M_\\odot -$ several $\\times 10^6\\,M_\\odot$) black holes are among the most anticipated sources for the LISA gravitational-wave antenna.  The measurement of these waves is expected to map the spacetime of the larger body with high precision, allowing us to test in detail the hypothesis that black hole candidates are described by the Kerr metric of general relativity.  In this article, we will briefly describe how these sources can be used to perform such a test.  These proposed measurements are often described as ``testing relativity''.  This description is at best somewhat glib: Because --- at least to date --- all work related to these measurements assumes general relativity as the theoretical framework in which these tests are performed, the measurements cannot be said to ``test relativity'' in a fundamental way.  More accurately, they test the {\\it nature of massive compact bodies within general relativity}.  A surprising result for such a test could point to deviations from general relativity, and would provide an experimentally motivated direction in which to pursue tests of gravity theories beyond GR. ",
        "introduction": "Compact binary sources in which one member is far more massive than another produce {\\it extreme mass ratio inspirals}, or EMRIs. Binaries of this sort are expected to form by scattering processes in the nuclei of galaxies: multibody interactions in the dense cluster of stars surrounding a galaxy's central black hole will sometimes place a compact object on an orbit that is so strongly bound to the black hole that it is more likely to inspiral due to gravitational-wave (GW) backreaction than it is to be scattered back out by further interactions.  See Clovis Hopman's article in these proceedings {\\cite{clovis}} for a detailed description.  There has been quite an evolution in our thinking about these sources over the past decade.  At the time of the 1st LISA Symposium, it was thought that the inspiralling compact body was most likely to be a white dwarf {\\cite{sigrees,sig97}}; the inferred LISA event rate for white-dwarf-dominated EMRIs was expected to be about an event per year.  It is now believed that measured EMRIs will be dominated by the capture of stellar mass ($\\sim 10\\,M_\\odot$) black holes.  Though the black hole birth rate is much lower than that of white dwarfs, black holes are captured more efficiently due to {\\it mass segregation} (the tendency of heavier stellar objects to sink to the bottom of their gravitational potential well) {\\cite{spitzer}}.  In addition, black hole EMRIs are ``audible'' to much greater distances since the EMRI wavestrain scales with the mass of the smaller body.  These factors, plus physics such as ``resonant relaxation'' (see Hopman's article {\\cite{clovis}} for discussion), have combined to boost the estimated event rate to several hundred events per year!  Capture might in fact be so efficient that galaxies ``eat'' all their EMRI sources early in their lives, leaving few low-redshift EMRIs for LISA unless they are replenished through later stellar evolution (a possibility discussed by Sigurdsson and Rees {\\cite{sigrees}}).  Conversely, the event rate might be so large that {\\it too many} EMRIs will be present in LISA data, and the sources may be confused {\\cite{bc04_conf}}.  GW data will have a lot to say about the astrophysics of these sources.  The relativity community has been attracted to this source in large part because it is an amazingly pristine astrophysical system.  Since the mass ratio of EMRIs is, by definition, very small ($10^{-3} \\lesssim m_1/m_2 \\lesssim 10^{-7}$), it can be treated as a perturbative parameter.  These binaries can thus be analyzed using a variant of perturbation theory, expanding the spacetime (or a suitable surrogate for the spacetime) about the Kerr background of the binary's larger member.  Schematically, the binary's evolution can be regarded as a ``forced geodesic'': Letting $x^\\alpha$ stand for the coordinates of the smaller body, and writing the connection for the background spacetime as $^{\\rm B}{\\Gamma^\\alpha}_{\\beta\\gamma}$, the smaller body follows a trajectory given by \\begin{equation} \\frac{d^2x^\\alpha}{d\\tau^2} + ^{\\rm B}{\\Gamma^\\alpha}_{\\beta\\gamma} \\frac{dx^\\beta}{d\\tau}\\frac{dx^\\gamma}{d\\tau} = f^\\alpha_{\\rm SF}\\;. \\label{eq:schematic} \\end{equation} If the right-hand side were zero, this would simply be the geodesic equation describing motion in the background spacetime.  The {\\it self force} $f^\\alpha_{\\rm SF}$ that appears instead describes how the motion is changed due to the small body's interaction with its own distortion to the background spacetime.  The challenge becomes computing this self force and its impact upon the small body's motion {\\cite{mst97,qw97}}, and then computing the GWs that this motion generates.  See Poisson {\\cite{ericLR}} for a particularly readable review.  For our purposes, it suffices to note that computing the self force without any additional simplifying assumptions (other than extreme mass ratio) has proven to be quite a challenge.  Yasushi Mino {\\cite{mino03}} demonstrated that {\\it if} the evolution can be regarded as ``adiabatic'' (in a sense to be quantified momentarily), then the self force can be simplified to something that is relatively tractable: \\begin{equation} f^\\alpha_{\\rm SF} = \\frac{1}{2}\\left(f^\\alpha_{\\rm ret} - f^\\alpha_{\\rm adv}\\right)\\;. \\end{equation} (This result is identical in form to the flat space electromagnetic self force originally derived by Dirac {\\cite{dirac}}.)  In contrast to the exact expression, which requires integrating a Green's function over the entire past world line of the orbiting body, the ``retarded'' and ``advanced'' contributions are relatively simple to evaluate.  The assumption of adiabaticity reduces that integration to an averaging procedure; what remains can be used to describe the binary's evolution using tools that are not much more complicated than those which have been developed for simpler cases.  See {\\cite{hdff,dfh,sthn}} for further discussion.  The adiabatic assumption that makes it possible to use Mino's ``half advanced minus half retarded'' prescription amounts to a separation of time scales.  We require that the change in any quantity that characterizes the orbit (such as energy or angular momentum) be ``small'' over a single ``orbit''.  Thus, over short timescales, we can treat the system as following an orbit that neglects radiative backreaction; over long timescales, we treat the inspiral as a ``flow'' of the system through a sequence of these orbits.  In practice, we currently treat the short timescale orbits as geodesics of the background spacetime.  This construction throws away ``conservative'' components of the self force --- a time symmetric, non-dissipative piece of $f^\\alpha_{\\rm SF}$ that pushes the orbit away from a geodesic of the background spacetime even in the absence of radiative backreaction.  Neglect of this piece will lead to a systematic phase error in models of the waveforms {\\cite{ppn05}}; this will have to be fixed to maximize the quality of LISA measurements (though it may be adequate for initial EMRI searches {\\cite{hf_inprep,fav_inprep}}).  Waveforms produced by this approximation are being developed now {\\cite{dh06}}, and appear to be sufficiently accurate for developing and testing LISA EMRI data analysis algorithms.  Some examples are shown in Figs.\\ {\\ref{fig1}} and {\\ref{fig2}}.  \\begin{figure} \\includegraphics[height=.4\\textheight]{generic.eps} \\caption{Generic EMRI waveforms.  Top panel shows waveform generated by a small body in a moderately strong field, moderately eccentric, highly inclined orbit about a rapidly spinning black hole; bottom panel shows waveform for a significantly eccentric and strong field orbit at shallower inclination angle about the same hole.  See Ref.\\ {\\cite{dh06}} for detailed discussion.  Superplotted on these waveforms are the ``kludge'' waves described in Ref.\\ {\\cite{kludge2}} (from which this plot is taken).  The kludge matches the strong-field calculation so well that the differences can barely be noted in this plot.} \\label{fig1} \\end{figure}  \\begin{figure} \\includegraphics[height=.4\\textheight]{inspiral.eps} \\caption{Circular EMRI inspiral waveform.  This plot shows waves in the (unrealistic) limit of zero eccentricity, but includes the effect of radiative backreaction; the upper curve shows the orbital radius as a function of time, $r(t)$.  (The waves shown in Fig.\\ {\\ref{fig1}} are of sufficiently short duration that the effect of backreaction is neglected.)  This example is for inspiral into a rapidly rotating hole with spin parameter $a = 0.9$.  Superplotted on this waveform is a ``kludge'' inspiral as described in Ref.\\ {\\cite{kludge2}} (from which this plot is taken).  As in Fig.\\ {\\ref{fig1}}, the kludge matches the strong-field embarrassingly well.} \\label{fig2} \\end{figure}  It's worth noting that recent work {\\cite{kludge1,kludge2}} has shown that almost all of the features of the waveforms developed in the strong field analysis can be reproduced using a very simple ``kludge'' approximation in which a {\\it flat spacetime} wave formula is coupled to the geodesic motion of a small body orbiting a black hole.  Indeed, the agreement between the two approaches is somewhat embarrassing --- we have superplotted waveforms produced by the ``kludge'' on the strong field results in Figs.\\ {\\ref{fig1}} and {\\ref{fig2}}.  The differences are very small; in particular, the phase coherency of the two approaches is remarkable.  Since the ``kludge'' waves can be computed {\\it many} times quicker than the strong-field waveforms (cpu minutes rather than cpu hours), we expect that the ``kludge'' approach will play a key role in scoping out and developing LISA EMRI data analysis. ",
        "conclusions": ""
    },
    "0608/astro-ph0608107_arXiv.txt": {
        "abstract": "The precise knowledge of the atmospheric neutrino fluxes is a key ingredient in the interpretation of the results from any atmospheric neutrino experiment. In the standard atmospheric neutrino data analysis, these fluxes are theoretical inputs obtained from  sophisticated numerical calculations. In this work we present an alternative approach to the determination of the atmospheric neutrino fluxes based on the direct  extraction from the experimental data on neutrino event rates. ",
        "introduction": "\\label{intro}  One of the most important breakthroughs in particle physics, and the only solid evidence for physics beyond the Standard Model, is the discovery -- following a variety of independent experiments \\cite{neutrev2} -- that neutrinos are massive and consequently  can oscillate among their different flavor eigenstates. The flavor oscillation hypothesis has been supported by an impressive wealth of experimental data,  one of the most important pieces of evidence coming from atmospheric neutrinos.  Atmospheric neutrinos are originated by the collisions of cosmic rays with air nuclei in the Earth atmosphere. In the collision mostly pions (and some kaons) are produced and they subsequently decay into electron and muon neutrinos and anti-neutrinos.  These neutrinos are observed in underground experiments using different techniques~\\cite{sk}. In particular in the last ten years, high precision and large statistics data has been available from the SuperKamiokande experiment~\\cite{sk} which has clearly established the existence of a deficit in the $\\mu$-like atmospheric events with the expected distance and energy dependence from $\\nu_\\mu\\rightarrow \\nu_\\tau$ oscillations with oscillation parameters $\\Delta m^2 \\sim 2\\times 10^{-3}$ eV$^2$ and $\\tan^2\\theta=1$.  This evidence has also been confirmed by other atmospheric experiments such as MACRO and Soudan 2.  The expected number of atmospheric neutrino events depends on a variety of components: the atmospheric neutrino fluxes, the neutrino oscillation parameters and the neutrino-nucleus interaction cross section. Since the main focus of atmospheric neutrino data interpretation has been the determination of neutrino oscillation parameters, in the standard analysis the remaining components of the event rate computation are inputs taken from other sources. In particular, the fluxes of atmospheric neutrinos are taken from the results of numerical calculations, like those of Refs. \\cite{honda,bartol}, which make use of a mixture of primary cosmic ray spectrum measurements, models for the hadronic interactions and simulation of particle propagation \\cite{fluxrev}.  The attainable accuracy in the independent determination of the relevant neutrino oscillation parameters from  non-atmospheric neutrino experiments makes it possible to attempt an inversion of the strategy in the atmospheric neutrino analysis: to use the oscillation parameters as inputs in the data analysis in order to extract from data the atmospheric neutrino fluxes.  There are several motivations for such direct determination of the atmospheric neutrino flux from experimental data.  First of all it would provide a cross-check of the standard flux calculations as well as of the size of the associated uncertainties (which being mostly theoretical are difficult to quantify).  Second, a precise knowledge of atmospheric neutrino flux is of paramount importance for high energy neutrino telescopes \\cite{icecubelect}, both because they are the main background and they are used for detector calibration.  Finally, such program will quantitatively expand the physics potential of future atmospheric neutrino experiments.  Technically, however, this program is challenged by the absence of a generic parametrization of the energy and angular functional dependence of the fluxes which is valid in all the range of energies where there is available neutrino data.  In this contribution we present the first results on this alternative approach to the determination of the atmospheric neutrino fluxes: we will determine these fluxes from experimental data on atmospheric neutrino event rates, using all available information on neutrino oscillation parameters, cross-sections, as well as information from the different sources of experimental and theoretical uncertainty. The problem of the unknown functional form for the neutrino flux is bypassed by the use of neural networks as interpolants, since they allow us to parametrize the atmospheric neutrino flux without having to assume any functional behavior. Additional details on our approach can be found in \\cite{paper}.  Indeed the problem of the {\\it deconvolution} of the atmospheric flux from experimental data on event rates is rather close in spirit to the determination of parton distribution functions in deep-inelastic scattering from experimentally measured structure functions \\cite{qcdbook}.  For this reason, in this work we will apply for the determination of the atmospheric neutrino fluxes a general strategy originally designed to extract parton distributions in an unbiased way with faithful estimation of the uncertainties\\footnote{ This strategy has also been successfully applied with different motivations in other contexts like tau lepton decays \\cite{tau} and B meson physics \\cite{bmeson}} \\cite{f2ns,f2nnp,pdf2}. ",
        "conclusions": "Now we discuss the results of our determination of the atmosferic neutrino fluxes. Note that any functional of the neutrino flux (like for example event rates and its associated uncertainty) can be computed taking the appropriate moments over the constructed probability measure \\be \\la \\mathcal{F}\\lp \\Phi\\lp E_{\\nu}\\rp\\rp \\ra_{\\rep}= \\frac{1}{N_{\\rep}}\\sum_{k=1}^{N_{\\rep}} \\mathcal{F}\\lp \\Phi^{(\\net)(k)}\\lp E_{\\nu}\\rp\\rp \\ . \\ee In Fig. \\ref{fluxref} we show our results for the atmosferic neutrino flux as compared with the Honda and Bartol computations as well as with some direct measurements from Amanda \\cite{amandadata}. In Fig. \\ref{fluxosc} we show that the effects of the uncertainties in the neutrino oscillation parameters are rather small with respect the intrinsic flux uncertainties. Finally, we show in Fig. \\ref{dataplot} how our neural network parametrization successfully reproduces the features of experimental data.  \\begin{figure} \\centering \\includegraphics*[scale=0.48]{fluxref.epsi} \\caption{Comparison of the $1-\\sigma$ uncertainty band of our parametrization of the atmospheric neutrino flux with recent computations and AMANDA data.} \\label{fluxref} \\end{figure}  \\begin{figure} \\centering \\includegraphics*[scale=0.48]{oscplot.ps} \\caption{Results for fits varying the default neutrino oscillation parameters from their central values within their 1-$\\sigma$ uncertainties \\cite{globalfits}.} \\label{fluxosc} \\end{figure}  \\begin{figure} \\centering \\includegraphics*[scale=0.48]{data6.ps} \\caption{Comparison of the neural network results with experimental data for the Multi-GeV $\\nu_{\\mu}$ data sample.} \\label{dataplot} \\end{figure}  The work presented in this contribution can be extended in several directions. First of all one could reduce the uncertainty in the large $E_{\\nu}$ region by incorporating in the fit atmosferic neutrino measurements from neutrino detectors like Amanda \\cite{amandadata}. Second, we could use at even larger energies simple parametrizations of the flux to obtain a better estimation of backgrounds at neutrino telescopes. Finally, we could estimate the required statistics that would be required in forthcoming atmosferic neutrino experiments in order to determine from data also the zenith angle and flavor dependence of the neutrino flux."
    },
    "0608/astro-ph0608277_arXiv.txt": {
        "abstract": "We present the constraints on inflationary parameters in a flat $\\Lambda$CDM universe obtained by WMAP three year data release, plus smaller scale CMB and two LSS data sets, 2dF and SDSS (treated separately). We use a Markov Chain Monte Carlo (MCMC) technique combined with an analytic description of the inflationary spectra in terms of the horizon flow functions (HFF). By imposing a consistency condition for the tensor-to-scalar ratio, we study the constraints both on single field standard inflation and on inflation with the violation of the null energy condition, which leads to a blue spectrum for gravitational waves. For standard inflation, the constraint on the tensor-to-scalar ratio we obtain from CMB data and 2dF05 is: $r_{0.01} < 0.26$ at 2 $\\sigma$ cl. Without the consistency condition between the tensor-to-scalar ratio and the tensor slope, the constraints on the tensor amplitude is not significantly changed, but the constraints on the HFFs are significantly relaxed. We then show that when the third HFF $\\epsilon_3$ is allowed to be non-zero and to be of order unity, a large negative (at $2 \\sigma$) value for the running of the scalar spectral index in standard inflation is found in any set of data we consider. ",
        "introduction": "We use the publicly available Markov Chain Monte Carlo (MCMC) package \\cite{LB} which makes use of the Einstein-Botzmann code CAMB \\cite{camb} for the computation of theoretical CMB anisotropies PSs and transfer functions. We restrict ourselves to flat spatial sections and to three species of massless neutrinos. We use as initial conditions for cosmological fluctuations the growing adiabatic mode. We include the CMBsmall data set as by default implemented in COSMOMC. Note that small scale data by the same experiments have been used by the WMAP team in Ref. [9], although there may be additional cross-correlations between the data sets which we consider and WMAP3 data. We restrict our analysis to linear transfer functions and therefore the LSS data are automatically truncated by the MCMC to $k_{\\rm cut}/h \\simeq 0.15$ ($0.2$) for 2dF (SDSS). The convergence diagnostic on runs with multiple chains is the R statistic by Gelman and Rubin implemented in the MCMC.  We use the $\\ln(k/k_*)$ expansion of the logarithm of the parametrizations in Eqs. (\\ref{i_spectra}), in order to avoid that the spectrum becomes negative for some values of the inflationary parameters explored by MCMC \\cite{LLMS}. According to Eqs. (\\ref{i_spectra}), the $\\ln(k/k_*)$ expansion of the power spectra is: \\begin{equation} \\label{plex} \\ln {P (k)\\over P_0(k_*)} = b_0 + b_1 \\ln \\left(k\\over k_*\\right) + \\frac{b_2}{2} \\ln^2\\left(k\\over k_*\\right) + \\dots . \\end{equation} and we use both the GFM \\cite{GS} and MCE \\cite{LLMS} to relate the coefficients $b$s to the HFFs. The coefficients for the scalar spectrum are: \\begin{eqnarray} \\label{eqn:bs0} b_{{\\rm S}0} &=& - 2\\left(C + 1\\right)\\epsilon_1 - C \\epsilon_2 + \\left(- 2C + {\\textstyle\\frac{\\pi^2}{2}} - 7\\right) \\epsilon_1^2 \\nonumber\\\\ & & + \\left({\\textstyle\\frac{\\pi^2}{8}} - 1\\right) \\epsilon_2^2 + \\left(- X^2 - 3X + {\\textstyle\\frac{7\\pi^2}{12}} - 7 + 2 \\Delta_{{\\rm S} 0} \\right) \\epsilon_1\\epsilon_2 \\nonumber \\\\ & & + \\left(-{\\textstyle\\frac 12}X^2 + {\\textstyle\\frac{\\pi^2}{24}} + \\Delta_{{\\rm S} 0} \\right) \\frac{d \\epsilon_2}{d N} \\, \\\\ \\label{eqn:bs1} b_{{\\rm S}1} &\\equiv&  n_S - 1 = - 2 \\epsilon_1 - \\epsilon_2 - 2 \\epsilon_1^2 -\\left(2\\,X+3\\right)\\,\\epsilon_1\\,\\epsilon_2 -X\\, \\frac{d \\epsilon_2}{d N} \\, \\\\ b_{{\\rm S}2} &\\equiv& \\alpha_S = - 2 \\epsilon_1 \\epsilon_2 - \\frac{d \\epsilon_2}{d N} \\label{eqn:bs2}, \\end{eqnarray} and for tensors are: \\begin{eqnarray} b_{{\\rm T}0} &=& - 2\\left(C + 1\\right)\\epsilon_1 + \\left(- 2C + {\\textstyle\\frac{\\pi^2}{2}} - 7\\right) \\epsilon_1^2 \\nonumber\\\\ & & + \\left(-X^2 - 2X + {\\textstyle\\frac{\\pi^2}{12}} - 2 + \\Delta_{{\\rm T} 0} \\right) \\epsilon_1\\epsilon_2 , \\label{eqn:bt0} \\\\ b_{{\\rm T}1} &\\equiv&  n_T = - 2\\epsilon_1 - 2\\epsilon_1^2 -2\\,\\left(X+1\\right)\\,\\epsilon_1\\,\\epsilon_2 , \\\\ \\label{eqn:bt1} b_{{\\rm T}2} &\\equiv&  \\alpha_T = - 2\\epsilon_1\\epsilon_2, \\label{eqn:bt2} \\end{eqnarray} where $C \\equiv \\ln 2+\\gamma_{\\rm E}-2\\approx-0.7296$ ($\\gamma_{\\rm E}$ is the Euler-Mascheroni constant) and $d \\epsilon_2 / d N = \\epsilon_2 \\epsilon_3$. At second order the next-to-leading coefficients depend on the method of approximation used to study cosmological perturbations during inflation: $X=C$ and $\\Delta_{{\\rm S} 0} = \\Delta_{{\\rm T} 0} = 0$ are obtained within the GFM \\cite{GS,LLMS}; $X= D \\equiv \\frac{1}{3}-\\ln 3\\approx-0.7652$, $\\Delta_{{\\rm S} 0} = (D-C) (D + \\ln \\, 2) -1/18$, $\\Delta_{{\\rm T} 0} = 2D(D-C) -1/9$, within the MCE \\cite{CFKLV}. We therefore implement the above expansion in a modified version of the inflationary routines developed by Samuel Leach and public at http://astronomy.sussex.ac.uk/$\\sim$sleach/inflation/camb\\_inflation.html.  As pivot scale we choose $k_*=0.01 \\, {\\rm Mpc}^{-1}$ as in \\cite{leachliddle} (see however \\cite{LPLM}); this choice differs from those by the WMAP team ($k_*=0.002 \\, {\\rm Mpc}^{-1}$ \\cite{peiris,spergel} and $k_*=0.05 \\, {\\rm Mpc}^{-1}$ \\cite{spergel_1}) and is better placed at the center of the range of wavenumbers probed by cosmological perturbations studied in the linear regime. Although the CMB spectrum includes at least a convolution of the power-spectrm with spherical Bessel functions, an effective pivot $\\ell_*$ corresponding to $k_*$ can be defined \\cite{huthesis,LLMS}: \\be k_*= \\frac{3}{2} \\frac{h \\sqrt{\\Omega_{\\rm m}}}{1+0.084 \\ln\\Omega_{\\rm m}}\\ell_* \\times 10^3 \\,{\\rm Mpc}^{-1} \\,. \\ee Since $\\ell_* \\simeq 145$ corresponding to $k_*=0.01 \\, {\\rm Mpc}^{-1}$ does not represent an interesting reference scale to measure the tensor-to-scalar ratio, we shall also quote in our results both $r_{k_*}$ and the ratio of tensor to scalar contribution to $C_\\ell$ and $l=10$, denoted by $R_{10}$ ($\\equiv C_{10}^T/C_{10}^S$).  As parameter describing reionization we shall vary the redshift $z_{\\rm re}$, which is related to the optical depth $\\tau$ by an integration over the redshift $z$ \\cite{LB} \\be \\tau = \\sigma_T \\int_0^{z_{\\rm re}} d z \\frac{n_e (z)}{H(z) (1+z)^2} \\,, \\ee with $\\sigma_T$ as the Thompson cross section and $n_e$ is the electron density number. As for the tensor-to-scalar ratios $r$ and $R_{10}$, we also give the constraints on $\\tau$ as a derived parameter. ",
        "conclusions": "We have analyzed the observational constraints from CMB and LSS on the inflationary expansion. By parametrizing the inflationary spectra as power-law (i.e. to first order in slow-roll parameters), we have updated the WMAP1 constraints of Ref. \\cite{leachliddle}: the better determination of reionization obtained by the 3 year data release disfavour now both the Harrizon-Zeldovich model (with no tensor) as inflationary models predicting $n_s = 1$ with respect to the first year data. Our analysis shows that models with a natural exit from inflation (which stay in the $\\epsilon_2 > 0$ region of the $(\\epsilon_2 \\,, \\epsilon_1)$ plane according to \\cite{dominik}) are now favoured by WMAP3.  When running is included, we have shown how the prior on $\\epsilon_3$ plays a crucial role in re-obtaining the WMAP team results within an analysis employing analytic power-spectra depending on ($H \\,, \\epsilon_i$) with $i=1-3$. Only allowing a broad prior on $\\epsilon_3$ a fairly large negative running in agreement with the WMAP results is obtained. We have found that is more convenient to use $d \\epsilon_2 / d N = \\epsilon_2 \\epsilon_3$ rather than $\\epsilon_3$. The value for the running $\\alpha_S$ found here implies $\\alpha_S \\ln (k/k_*)/2 > (n_S -1)$ for Fourier modes far from the pivot scale. It would be therefore interesting to perform a comparison of data with inflationary predictions at third order in the HFF. Note that the value of the running found here is different from zero at $\\sim 2 \\sigma$ level, and its constraints may also be affected by Ly$\\alpha$ data \\cite{viel,seljak}, which we do not consider here.  We have also studied the impact of theoretical priors on constraining the inflationary expansion: still little is known about high-energy physics and prejudices in interpreting cosmological data may hide interesting physics.  We have first relaxed the hypothesis of a canonical scalar field as the inflaton, which leads to the consistency relation in Eq. (\\ref{ttos}). By considering a generic Lagrangian, the constraints on the $(\\epsilon_1 \\,, \\epsilon_2)$ plane degrade significantly, while the limits on $r$ depend on the slope of the tensor spectrum. This analysis shows that a KG inflaton (with sound speed $c_S=1$) is well inside the confidence contours of the present constraints, although the present data do not allow to constrain $c_S$.  We have then explored how CMB and LSS constrain inflationary models with a blue spectrum, a more radical departure than considering $r$ as a free parameter. We have restricted ourselves to KG field with the \"wrong\" sign for the kinetic term, in order to do not consider the tensor-to-scalar ratio as an independent parameter. We find that both standard inflation and NEC violating inflation are equally good fits to the data (the best fits for these two classes of models have the same $\\chi^2$ comparing with WMAP3+CMBsmall+2dF02). For $n_T > 0$, the constrained region on the $(\\epsilon_1 \\,, \\epsilon_2)$ are different from the standard case, leading to $0.02 \\lesssim \\epsilon_2 \\lesssim 0.08$ with $|\\epsilon_1| \\lesssim 0.3$ at 2$\\sigma$ cl. A simple exponential potential potential is disfavoured at 2 $\\sigma$ cl and larger values for $r$ are allowed with respect to the standard case with $n_T < 0$.   \\vspace{1cm}  {\\bf Acknowledgements:} We wish to thank Samuel Leach for help and numerous suggestions on the use of MCMC. We also thank R. Easther, W. Kinney, A. Liddle, H. Peiris, R. Trotta and A. Vikman for comments and discussions. An anonymous referee is acknowledged for her/his patient and critical reading of the manuscript which led to an improvement of this paper. FF is partially supported by INFN ISs BO11 and PD51. FF thanks the Galileo Galilei Institute for Theoretical Physics for the hospitality and the INFN for partial support during the completion of this work Part of the results presented here have been obtained on CINECA Linux cluster under the agreement INAF/CINECA. We acknowledge the use of the package {\\sc cosmomc} and the use of the Legacy Archive for Microwave Background Data Analysis (LAMBDA). Support for LAMBDA is provided by the NASA Office of Space Science."
    },
    "0608/astro-ph0608041_arXiv.txt": {
        "abstract": "{The linear growth rate and saturation level of magnetic fields for Weibel instabilities driven by  ion temperature anisotropy, defined as $\\alpha=(T_\\perp/T_\\parallel)-1$ where $T_\\perp$ and $T_\\parallel$ are ion temperatures perpendicular and parallel to the wave vector, are derived in the small $\\alpha$-limit. It is shown that the ratio of the saturated magnetic energy to the initial ion energy scales as the fourth power of the electron to ion mass ratio, $m/M$, for an initially unmagnetized plasma with $\\alpha\\leq M/m$.  Particle-in-cell simulations confirm the mass scaling and also show that the electron energy gain is of the same order of magnitude as the magnetic field energy. This implies that the Weibel instabilities cannot provide a faster-than-Coulomb collisionless mechanism to equilibrate ion-electron plasmas with ions initially much hotter than electrons, a key component in low-luminosity astrophysical accretion flows. The results here also show that the large $\\alpha$-limit formulae used in the study of magnetic field generation in collisionless shocks are only valid if $\\alpha\\geq M/m$.} ",
        "introduction": "\\subsection{Weibel instabilities and astrophysical shocks}  { Weibel instabilities, whereby an initially anisotropic plasma distribution relaxes to convert the free energy in the anisotropy into magnetic fields e.g. \\cite{w59,davidson72,ml99} have received much recent attention in the study of astrophysical plasmas, particularly in the context of influence on the observed emission associated with collisionless astrophysical shocks. These include shocks of the crab pulsar wind \\cite{yang94}, gamma-ray bursts (GRB)\\cite{ml99}, and between or within  galaxy clusters \\cite{msk06,markevitch06}.  The former two are relativistic. The roles that the Weibel instabilities may play via their ability to grow magnetic fields  are (i) to magnetically trap particles near the shock, facilitating the shock formation itself; (ii) to account for synchrotron or jitter radiation that would otherwise be absent and hard to explain without a source of field growth; (iii)  to suppress  thermal conductivity, as needed to explain observed temperature inhomogeneities \\cite{msk06} in the ion-electron plasmas of Galaxy clusters; and (iv) possibly producing a faster-than Coulomb coupling or electron acceleration in the post-shock regions near shocks within or between Galaxy clusters. This is important for the estimate of Galaxy cluster masses and therefore fractional dark matter mass fraction estimates \\cite{markevitch06}. The latter are inferred from measured electron temperatures, the assumption that proton temperatures are equal to the X-ray inferred electron temperatures, and virial equilibrium.   The plasma compositions for the astrophysical applications range from  ion-electron to positron- electron pair plasmas, to a mixture between the two. In shocks, the systematic velocity anisotropy is typically assumed to be strong, and is produced by fast particle streaming in the shock propagation direction compared to the directon parallel to the shock.  Current fluctuations parallel to  the direction of  shock propagation lead to current filaments that amplify the magnetic field in the shock plane, which then  further enhances the current filaments. }   \\subsection{Weibel instabilities and two-temperature plasmas}  {We introduce the question of whether the Weibel instability may also facilitate a  faster-than-Coulomb coupling between ions and electrons in two-temperature plasmas even unshocked. The free energy in anisotropic need not be the result of particle streaming perpendicular to  shocks, but would also result from anisotropy in the local distribution of random particle velocities of the  species supplying the free energy in the absence of shocks. The study of the Weibel instability without shocks also has a long history \\cite{davidson72} but has not been studied for plasmas containing ions and electrons at  very different average eneriges,the motivation for which we now describe further.}  Left alone, a two-temperature ion-electron plasma, with the ions at a higher temperature than the electrons,  will eventually approach an equilibrium in which the two species have the same temperature. Coulomb collisions provide a minimum equilibration rate, but whether a faster-than-Coulomb collisionless equilibration mechanism exists\\cite{begelman,blackman99,quataert99} for specific particle distributions and external conditions is both a fundamental question in plasma physics and important for understading aspects of collisionless astrophysical phenomena such as radiatively inefficient accretion flows (RIAF) (e.g. Ref.\\cite{quataert01}).    The survival of a two-temperature plasma for a Coloumb equilibration time scale is important for  geometrically thick RIAFs of which advection dominated accretion flows (ADAFs) have been the archetype \\cite{ichimaru77,rees82,narayan95,narayan98}. In contrast to standard geometrically thin accretion discs (e.g. Ref.\\cite{shakura73}), which can comfortably allow 10\\% of the binding energy of the accreting material to be converted into observable radiation around black holes, two-temperature accretion disks can in principle reduce the photon luminosity by many orders of magnitude for the same accretion rate. Indeed,  the measured luminosities from the central engines of some nearby elliptical galaxies seem to be  3 to 5 orders of magnitude smaller than expected for their estimated accretion rates. Similar quiescient phases of accretion are also observed in some states of X-ray binaries.  There is also X-ray evidence that in some supernova remnants the ion temperature is much higher than the electron temperature\\cite{vink05}.   In the simplest ADAF model, the reduced luminosity results from three key features: (1) energy dissipated by viscosity is assumed to heat the protons, (2) only electrons radiate efficiently, and (3) if the only means of energy transfer between electrons and protons is Coulomb collisions, then a sufficiently collisonless plasma can be accreted onto the central object  before the electrons receive enough energy to radiate the energy dissipated by  the accretion. The gravitational binding energy of the accreting material is then retained as internal energy of the hot protons, and can be quiescently advected through the event horizon of a black hole. Since the associated internal energy per ion in an ADAF is of order $m_pc^2 \\sim$ 1GeV, the electrons are required to stay at a temperature much lower than the GeV ion temperature to ensure a low luminosity for a given accretion rate. ADAFs are thick discs because the thermal energy of the ions puffs up the disk. More general models of two temperature accretion involving  outflows and convective feedback are also popular\\cite{qua_nar99,bla_beg99,naretal02}.   The assumption that heat transfer from the ions to the electrons occurs only by Coloumb collisions is necessary to ensure that a negligible fraction of the dissipated thermal energy is transferred to the electrons during the accretion infall time. This assumption has yet to be validated or invalidated (e.g. Ref.\\cite{pariev05}). Doing so requires understanding the subtle plasma physics of the interactions among the ions, electrons,  and electromagnetic fields. A systematic way to make  progress is to test this assumption under a range of  specific conditions. If the assumption is found to be violated for particular circumstances, then the associated instability and faster-than-Coulomb coupling would  limit the applicability of  two-temperature accretion paradigm. If the assumption is not violated for a particular set of conditions, this would not prove the viablity of two-temperature plasmas in  all circumstances and more cirumstances would need to be tested. While the latter issue has discouraged some people from working on this problem, we think that a series of rigorous tests using different initial and boundary conditions is needed:  Were the assumption to emerge as consistently valid for a range of different initial conditions, this would gradually solidify the robustness of the assumption for an increasingly wider range of viable conditions, and justify two-temperature accretion flows for those regimes tested.  In this spirit,  we study one test of the Coulomb equilibration assumption here, namely, whether ion-electron energy transfer can be sped up via Weibel instabilities driven by an ion distribution anisotropy. For a real disk, the microphysics by which gravitational potential energy is converted to ion thermal energy is not fully understood, but presumably  involves turbulence generated by a collisionless version \\cite{sharma03} of a magneto-rotatioanl instability \\cite{balbus98}. MHD turbulence is believed to anisotropically cascade \\cite{goldreich97} and it is likely in a collisionless system, the proton temperature may also incur anisotropies on small scales, transiently induced by the cascading turbulence. Similar to an anisotropic electron distribution, an anisotropic ion distribution  can also drive a Weibel instability \\cite{w59,davidson72} which amplifies fluctuating magnetic fields. How fast and how much energy the electromagnetic fields and the electrons then gain from the isotropizing protons is the subject of our present work.  \\subsection{Growth and Saturation of the Weibel Instability}  {Since  electrons gain energy from  electromagnetic fields generated  the Weibel instability, the electron energy gain is closely related to the energy level of the magnetic field, the dominant electromagnetic field at saturation. The strength of the saturated magnetic field level from the Weibel instabilities is therefore important for all of the astrophysical applications discussed above:  two temperature accretion plasmas as well as the origin of magnetic fields in Gamma-ray bursts\\cite{ml99,wa04} and galaxy clusters \\cite{msk06}. The saturation mechanism is also a fundamental plasma physics question.  If energy transfered to electrons  were a significant fraction of the initial ion energy during an instability growth time for the two temperature plasmas discussed in the previous section, this transfer would be a much more efficient process than Coulomb collisions for temperature equilibration. The is because   the ratio of maximum growth rate of the ion-driven Weibel instability $\\gamma_{max}$ (see Eq.\\ref{gmax} below) to collisional ion-electron equilibration rate $\\nu_{eI}$ is \\begin{equation} \\label{governu} {\\gamma_{max}\\over\\nu_{eI}}=7.2\\alpha^{3/2}{n\\lambda_{D}^3\\over ln\\Lambda}\\sqrt{{T_e\\over Mc^2}}\\left(1+{T_{I}m\\over T_{e}M}\\right)^{3/2}. \\end{equation} Here $\\alpha \\ge 0$ is a measure of ion temperature anisotropy (defined more precisely in Sec.\\ref{theory}), $n$ the electron density, $\\lambda_{D}$ the Debye length, $ln\\Lambda$ the Coulomb logarithm, $T_e$ and $T_I$ the electron and ion temperatures (multiplied by the Boltzmann constant, as throughout the paper), respectively, $m$ and $M$ the electron and ion masses, respectively, and $c$ the velocity of light. For a proton-electron plasma with $n=10^{12} cm^{-3}, T_I=10^{12}K, \\alpha=1,$ and $T_e=T_Im/M, {\\gamma_{max}/\\nu_{eI}}\\sim 3\\times10^7$. This ratio is large, but  the question is what fraction of the ion energy will be transfered to the electrons when the instability saturates?    There have been  two saturation mechanisms discussed in the literature: (i) the deviation from the equilibrium orbits of the particles with an anisotropic distribution \\cite{davidson72,yang94,wa04} and (ii) the magnetization of the ions \\cite{ml99} or electrons \\cite{msk06}. These two mechanisms gave the same magnetic field level in the electron-driven case (for either ion-electron or pair plasmas) but differ in the ion-driven case \\cite{wa04}. In the present paper we will  show and reaffirm that  the orbit deviation is the correct magnetic saturation mechanism for the Weibel instabilities.  We will derive the growth rate and magnetic saturation level in the small anisotropy limit in Sec.\\ref{theory} which shows that the fraction of the ion energy transfered to the magnetic field via the Weibel instabilities scales as $(m/M)^4$ for a single mode. Connection will also be made to previous results in the large anisotropy limit. Particle-in-Cell (PIC) simulations in Sec.\\ref{sim} with reduced ion mass will confirm this mass scaling and also will show that the electron energy gain is of the same order of magnitude as the magnetic field energy. For an initially unmagnetized plasma, this implies that the Weibel instabilities are  insignificant to the ion-electron temperature equilibration process. However, in Sec.\\ref{dis} we discuss how the transfered energy fraction via the Weibel instabilities may greatly increase for an initially more strongly magnetized plasma. We conclude  in Sec. \\ref{dis}.} ",
        "conclusions": "\\label{dis} In summary, the linear growth rate of the Weibel instability driven by the proton temperature anisotropy has been calculated in the low anisotropy limit [Eqs.(\\ref{rate}) and (\\ref{gmax})]. This growth rate is also used to derive an expression for the ratio of saturated magnetic field energy to the initial proton kinetic energy, $E_B/E_{kp}$, in the low anisotropy limit [Eq.(\\ref{ebk})]. These results show that $E_B/E_{kp}$ scales as $(m/M)^4$ when $\\sqrt{T_e/m}\\sim\\sqrt{T_{Ix}/M}$ and is extremely small for $\\alpha=1$. 1.5D PIC simulations with a reduced ion to electron mass ratio support  the scaling relations and also show that the electron energy gain from the proton anisotropy is comparable to the gain in  magnetic field energy. Although in three dimensions one additional mode polarization is allowed and the number of unstable modes can double, we do not expect orders of magnitude change in the electron energy gain. From these analyses and simulations we conclude that the proton-driven Weibel instability will not  provide an efficient faster-than-Coulomb electron-ion temperature equilibration mechanism in an intially unmagnetized two temperature plasma with low to moderate anisotropy.  {We have also shown that the low anisotropy formulae derived here approach to the previous high anisotropy results when $\\alpha\\rightarrow M/m$ [Eqs.(\\ref{large_rate}) and (\\ref{ebk2})]. This makes explicit the implied assumption of $\\alpha\\geq M/m$ used in the discussion of magnetic field generation in collisionless shocks through the Weibel instability\\cite{wa04,msk06}. It also provides guidance on how to properly scale PIC simulations results with reduced $M/m$. For example, PIC simulations on shocks in galaxy clusters with Mach number $M_{sh}=20$, electron thermal velocity $v_{th,e}=0.05c$, and ion thermal velocity $v_{th,e}=0.005c$ were performed using $M/m=100$ in Ref.\\cite{msk06}. Since $\\alpha=M_{sh}^2(m/M)(v_{th,e}/v_{th,i})^2=400>M/m$, the simulations showed an $E_B/E_{kp}\\approx10^{-3}$, consistent with $m/(16M)$ predicted from Eq.(\\ref{ebk2}). However, if the same $M, v_{th,e},$ and $v_{th,i}$ were used for $M/m=1836$, $\\alpha=22<M/m$ and Eq.(\\ref{ebk}) would predict an $E_B/E_{kp}\\approx 4\\times 10^{-6}$.}  Finally we wish to point out that the situation may be greatly different for  an initially magnetized plasma in which electrons are confined to local magnetic field lines transversely.  For modes with ${\\bf k}$ perpendicular to the equilibrium ${\\bf B}_0$, the electrons would be  practically immobile in the direction perpendicular to both ${\\bf B}_0$ and ${\\bf k}$. As explained in Sec.\\ref{theory}, the Weibel instability is driven by the current induced by the perturbed magnetic field $B_z$, represented by the term proportional to $\\alpha$ in Eq.\\ref{full_disp}. However, additional currents with a different phase are also driven by the accompanied electric field $E_y$, which are represented by the second and third terms on the left of  Eq.\\ref{full_disp}. These currents are present even for an isotropic distribution and oppose mode growth. If the electrons were magnetized and could not respond to $E_y$ to produce a curent, the second term would not be present and the growth rate would be independent of $m/M$. This indicates that the required equilibrium magnetic field should be such that $\\omega_{ce}>>\\omega_{pI}>>\\omega_{cI}$ for this to happen, where $\\omega_{ce}$ and $\\omega_{cI}$ are the electron and ion cyclotron frequencies in the equilibrium magnetic field. In the case of the ion-driven (electricstatic) two-stream instability with magnetized electrons, the results were indeed found to be independent of the mass ratio\\cite{papa71}. If this also were true for the Weibel instability more generally, the saturated magnetic fields could be $E_B/E_{kp}\\approx8\\times10^{-3}\\alpha^5/[(1+\\alpha)^5(3+2\\alpha)]=2.5\\times10^{-4}$ for $\\alpha=5.4$, significantly larger than that predicted by Eq.(\\ref{ebk}).  However,  the electrons still cannot gain significant energy even in this enhanced magnetic field."
    },
    "0608/astro-ph0608088_arXiv.txt": {
        "abstract": "We present new, deep 8.5-GHz VLA observations of the nearby, low-luminosity radio galaxy 3C\\,296 at resolutions from 0.25 to 5.5\\,arcsec.  These show the intensity and polarization structures of the twin radio jets in detail. We derive the spectral-index distribution using lower-frequency VLA observations and show that the flatter-spectrum jets are surrounded by a sheath of steeper-spectrum diffuse emission. We also show images of Faraday rotation measure and depolarization and derive the apparent magnetic-field structure.  We apply our intrinsically symmetrical, decelerating relativistic jet model to the new observations. An optimized model accurately fits the data in both total intensity and linear polarization. We infer that the jets are inclined by $58^\\circ$ to the line of sight.  Their outer isophotes flare to a half-opening angle of 26$^\\circ$ and then recollimate to form a conical flow beyond 16\\,kpc from the nucleus.  On-axis, they decelerate from a (poorly-constrained) initial velocity $\\beta = v/c \\approx 0.8$ to $\\beta \\approx 0.4$ around 5\\,kpc from the nucleus, the velocity thereafter remaining constant. The speed at the edge of the jet is low everywhere. The longitudinal profile of proper emissivity has three principal power-law sections: an inner region (0 -- 1.8\\,kpc), where the jets are faint, a bright region (1.8 -- 8.9\\,kpc) and an outer region with a flatter slope.  The emission is centre-brightened.  Our observations rule out a globally-ordered, helical magnetic-field configuration. Instead, we model the field as random on small scales but anisotropic, with toroidal and longitudinal components only. The ratio of longitudinal to toroidal field falls with distance along the jet, qualitatively but not quantitatively as expected from flux freezing, so that the field is predominantly toroidal far from the nucleus.  The toroidal component is relatively stronger at the edges of the jet.  A simple adiabatic model fits the emissivity evolution only in the outer region after the jets have decelerated and recollimated; closer to the nucleus, it predicts far too steep an emissivity decline with distance. We also interpret the morphological differences between brightness enhancements (``arcs'') in the main and counter-jets as an effect of relativistic aberration. ",
        "introduction": "\\label{intro}  The case that the jets in low-luminosity, FR\\,I radio galaxies \\citep{FR74} are initially relativistic and decelerate to sub-relativistic speeds on kiloparsec scales rests on several independent lines of evidence, as follows. \\begin{enumerate} \\item Proper motions corresponding to speeds comparable with or exceeding $c$ have been observed directly on pc and kpc scales in FR\\,I jets \\citep{Giov01,Biretta95,Hard03}. \\item FR\\,I sources must be the side-on counterparts of at least a subset of the BL Lac population, in which evidence for bulk relativistic flow on pc scales is well established (e.g. \\citealt{UP95}). \\item FR\\,I jets show side-to-side asymmetries which decrease with distance from the nucleus, most naturally explained by Doppler beaming of emission from an intrinsically symmetrical, decelerating, relativistic flow \\citep{LPdRF}, in which case the brighter jet is approaching us.  There is continuity of sidedness from pc to kpc scales and superluminal motion, if observed, is in the brighter jet. \\item The lobe containing the brighter jet is less depolarized than the counter-jet lobe, consistent with Faraday rotation in the surrounding halo of hot plasma if the brighter jet is approaching \\citep{L88,Morganti97}. \\item Theoretical work has demonstrated that initially relativistic jets can decelerate from relativistic to sub-relativistic speeds without disruption provided that they are not too powerful, the surrounding halo of hot plasma providing a pressure gradient large enough to recollimate them \\citep{Phi83,Bic94,Kom94,BLK96}. \\end{enumerate} We have developed a method of modelling jets on the assumption that they are intrinsically symmetrical, axisymmetric, relativistic, decelerating flows. By fitting to deep radio images in total intensity and linear polarization, we can determine the three-dimensional variations of velocity, emissivity and magnetic-field ordering.  We first applied the technique to the jets in the radio galaxy 3C\\,31 \\citep[LB]{LB02a}. With minor revisions, the same method proved capable of modelling the jets in B2\\,0326+39 and B2\\,1553+24 \\citep[CL]{CL} and in the giant radio galaxy NGC\\,315 \\citep[CLBC]{CLBC}. We used the derived velocity field for 3C\\,31, together with {\\em Chandra} observations of the hot plasma surrounding the parent galaxy, to derive the run of pressure, density, Mach number and entrainment rate along its jets via a conservation-law analysis \\citep{LB02b}. We also examined adiabatic models for 3C\\,31 in detail \\citep{LB04}.  We now present a model for the jets in 3C\\,296, derived by fitting to new, sensitive VLA observations at 8.5\\,GHz.  In Section~\\ref{obs}, we introduce 3C\\,296 and summarize our VLA observations and their reduction. Images of the source at a variety of resolutions are presented in Section~\\ref{descrip}, where we also describe the distributions of spectral index, rotation measure and apparent magnetic field derived by combining our new images with lower-frequency data from the VLA archive. We briefly recapitulate our modelling technique in Section~\\ref{model} and compare the observed and model brightness and polarization distributions in Section~\\ref{results}. The derived geometry, velocity, emissivity and field distributions are presented in Section~\\ref{physical}.  We compare the result of model fitting for the five sources we have studied so far in Section~\\ref{sourcecomp} and summarize our results in Section~\\ref{summary}.  We adopt a concordance cosmology with Hubble constant, $H_0$ = 70\\,$\\rm{km\\,s^{-1}\\,Mpc^{-1}}$, $\\Omega_\\Lambda = 0.7$ and $\\Omega_M = 0.3$. Spectral index, $\\alpha$, is defined in the sense $S(\\nu) \\propto \\nu^{-\\alpha}$ and we use the notations $P = (Q^2+U^2)^{1/2}$ for polarized intensity and $p = P/I$ for the degree of linear polarization. We also define $\\beta = v/c$, where $v$ is the flow velocity. $\\Gamma = (1-\\beta^2)^{-1/2}$ is the bulk Lorentz factor and $\\theta$ is the angle between the jet axis and the line of sight. ",
        "conclusions": "\\label{summary}  We have made deep, 8.5-GHz images of the nearby FR\\,I radio galaxy 3C\\,296 with the VLA at resolutions ranging from 0.25 to 5.5\\,arcsec FWHM, revealing new details of its twin jets. In particular we see several thin, discrete brightness enhancements (``arcs'') in both jets. These appear to have systematically different morphologies in the main and counter-jets.  A comparison with lower-frequency images from archive data shows that the flat-spectrum jets are surrounded by a broad sheath of steeper-spectrum diffuse emission, possibly formed by a backflow of radiating plasma which has suffered significant synchrotron losses as in models of FR\\,II sources. The spectral index of the jets initially flattens slightly with distance from the nucleus (from $\\alpha \\approx 0.62$ to $\\alpha \\approx 0.53$) as in other FR\\,I jets. We have also imaged the Faraday rotation and depolarization over the source. The counter-jet is more depolarized than the main jet, but the differences in depolarization and Faraday rotation variance between the two lobes are small. The rms fluctuations in rotation measure are larger close to the nucleus in both lobes.  We have shown that many features of the synchrotron emission from the inner $\\pm$40\\,arcsec of the jets in 3C\\,296 can be fit accurately and self-consistently on the assumption that they are intrinsically symmetrical, axisymmetric, decelerating, relativistic flows. The functional forms we use for the model are similar to those developed in our earlier work, as are most of the derived physical parameters.  In this sense, a strong family resemblance is emerging between the physical properties that we derive for the outflows in these FR\\,I sources.  \\begin{description} \\item [{\\bf Geometry}] The jets can be divided into a flaring region, whose radius $x$ is well fitted by the expression $x = a_2 z^2 + a_3 z^3$, where $z$ is the distance from the nucleus along the axis, and a conical outer region. We infer an angle to the line of sight of $58^{+3}_{-2}$\\,degree.  The boundary between the two regions is then at 15.8\\,kpc from the nucleus. \\item [{\\bf Velocity structure}] Where the jets first brighten, their on-axis velocity is $\\beta \\approx 0.8$, but this value is poorly constrained. Rapid deceleration occurs around 5\\,kpc to an on-axis velocity of $\\beta = 0.40^{+0.07}_{-0.03}$ which thereafter remains constant. 3C\\,296 differs from all of the other objects we have studied in having a very low velocity (essentially consistent with 0) at the edges of its jets. \\item [{\\bf Emissivity}] The longitudinal profile of rest-frame emissivity is modelled as three principal power-law sections $\\epsilon \\propto \\rho^{-E}$ with indices of $E = 2.5$ (0 -- 1.8\\,kpc), $2.8$ (2.5 -- 8.9\\,kpc) and $1.0$ (8.9 -- 24.5\\,kpc, the end of the modelled region). The first two regions are separated by a short transition zone in which the emissivity increases discontinuously by a factor of $\\approx$7 and then has a power-law slope $\\sim$2.1 (1.8 -- 2.5\\,kpc; representing the first knots in the main and counter-jets). The rest-frame emissivity is centre-brightened. \\item [{\\bf Magnetic-field configuration}] We model the magnetic field as disordered, but anisotropic, with toroidal and longitudinal components, but our results are equally consistent with an ordered toroidal component.  We rule out a globally-ordered helical field. The ratio of longitudinal to toroidal field decreases with distance from the nucleus; the toroidal component is also more dominant at the edges of the jets. \\item [{\\bf Flux freezing}] The field-component evolution is qualitatively consistent with flux freezing, but much slower than expected in the flaring region, where the adiabatic approximation also fails to fit the emissivity variation. In the outer region, after the jets decelerate, the adiabatic approximation is a reasonable fit. \\item [{\\bf Arc asymmetry}] If the arcs are narrow shells of enhanced emissivity moving with the underlying flow, then the differences in their appearance in the main and counter-jets can be understood as the effect of relativistic aberration, as required for a fully self-consistent model of the jets. \\end{description}"
    },
    "0608/hep-th0608195_arXiv.txt": {
        "abstract": "We study a brane-world cosmological scenario with local inhomogeneities represented by black holes. The brane is asymmetrically embedded into the bulk. The black strings/cigars penetrating the Friedmann brane generate a Swiss-cheese type structure. This universe forever expands and decelerates, as its general relativistic analogue. The evolution of the cosmological fluid however can proceed along four branches, two allowed to have positive energy density, one of them having the symmetric embedding limit. On this branch a future pressure singularity can arise for either (a) a difference in the cosmological constants of the cosmological and black hole brane regions (b) a difference in the left and right bulk cosmological constants. While the behaviour (a) can be avoided by a redefinition of the fluid variables, (b) establishes a \\textit{critical value of the asymmetry} over which the pressure singularity occurs. We introduce the \\textit{pressure singularity censorship} which bounds the degree of asymmetry in the bulk cosmological constant. We also show as a model independent generic feature that the asymmetry source term due to the bulk cosmological constant increases in the early universe. In order to obey the nucleosynthesis constraints, the brane tension should be constrained therefore both from below and from above. With the maximal degree of asymmetry obeying the pressure singularity censorship, the higher limit is $10$ times the lower limit. The degree of asymmetry allowed by present cosmological observations is however much less, pushing the upper limit to infinity. ",
        "introduction": "In the last decade brane-world models, originally motivated \\cite{ADD}-\\cite% {RS2} by string / M-theory, have developed into classical theories of gravitation, alternative to general relativity and reducing to it in a well-defined limit. In these models our observable universe is a co-dimension one brane with tension $\\lambda $ embedded in a 5-dimensional (5D) bulk and gravitation has more degrees of freedom than in general relativity. The fifth dimension is not compactified, as in the Kaluza-Klein theories, but instead curved, in the simplest case due to a bulk cosmological constant $\\widetilde{\\Lambda }$. Cosmological generalizations of the Randall-Sundrum (RS) second model \\cite{RS2} consist of a moving Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) brane embedded into a static Schwarzschild - anti de Sitter bulk. Cosmological dynamics is given by the motion of the brane in the bulk.  Our homogeneous and isotropic universe certainly contains local inhomogeneities. In principle, these can be dealt with the perturbation theory on the FLRW brane. However, due mainly to technical problems involving the boundary conditions to be imposed \\cite{MaartensPert}, \\cite% {MaartensLR} this is not a completed task for the moment. There is still much to learn until a complete perturbative description will be achieved. In order to get some insight into how the inhomogeneities could affect the cosmological evolution on the brane, exact toy models are worth to study.  This method has a great tradition from the early days of general relativity: in the simplest, Einstein-Straus (Swiss-cheese) model \\cite{ES}, in which Schwarzschild spheres of constant comoving radius are glued into the FLRW space-time, the luminosity-redshift relation is modified as compared to the FLRW universe \\cite{Kantowski}. The dynamics of brane-world models being much more complicated, the question arises, whether similar solutions can be found. A first investigation in \\cite{NoSwissCheese} although presents the correct form of the junction conditions between the Schwarzschild and FLRW brane regions, due to an error reaches the wrong conclusion that no such brane-worlds exist. In reality, the conclusion holds only in a static setup, while dynamical Swiss-cheese brane-worlds do exist. (This is pointed out in the published erratum). The simplest such brane-world model was later presented in \\cite{SwissCheese}. In this model, Schwarzschild voids are immersed in an evolving FLRW brane with flat spatial sections $k=0$ and cosmological constant $\\Lambda $. This Swiss-cheese brane is embedded symmetrically into the bulk.  The bulk for this Swiss-cheese brane has a complicated structure. While the FLRW brane can be embedded into a 5D anti de Sitter (AdS5) bulk, the same does not apply for the Schwarzschild regions. Indeed the perturbative analysis of the gravitational field of a spherically symmetric source in the weak field limit, when the bulk is the AdS5 space-time, has shown corrections scaling as $r^{-3}$ to the Schwarzschild potential \\cite{RS2}, \\cite{GT}-\\cite{GKR}. Therefore the Schwarzschild brane black hole (showing no correction to the Schwarzschild potential by definition) is not embeddible into AdS5. Rather the Schwarzschild black holes extend as black strings into the bulk \\cite{ChRH}. Gravity wave perturbations of such a black-string brane-world were recently discussed in \\cite{SCMlet}. The Gregory-Laflamme instability \\cite{GL} may cause the decay of the black string into a black cigar \\cite{Gregory}, but it has also been shown that under very mild assumptions, classical event horizons cannot pinch off \\cite% {HorowitzMaeda}. The Swiss-cheese type brane structure thus emerges of black strings / cigars penetrating a cosmological brane. The bulk contains transition zones between the black string / cigar regions and the AdS5 bulk regions (Fig \\ref{Fig0}). \\begin{figure}[tbp] \\vskip 1 cm \\center \\includegraphics[bb=50 50 400 400, height=7cm]{Fig0.eps} \\vskip 1 cm \\caption{Schematic representation of the asymetrically embedded Swiss-cheese brane-world. The FLRW regions are bounded by AdS5 bulk regions with different cosmological constant. The brane Schwarzschild black holes extend into the bulk as black strings / cigars. On each side of the brane transition zones of different shape separate the black strings / cigars from the AdS5 bulk regions. } \\label{Fig0} \\end{figure}  The evolution of the scale factor in this model surprisingly proceeds as in general relativity: the Swiss-cheese brane \\ with $k=0$ forever expands and forever decelerates. At a technical level, this happens because the non-vanishing source terms in the effective Einstein equation combine into an effective dust \\cite{SwissCheese}, re-establishing formally the general relativistic Einstein-Straus model with the effective dust source. At a more general level this means that the motion of the brane in the AdS5 bulk is not affected by the inhomogeneities introduced by the puncturing black strings.  However the cosmological fluid on the brane evolves in a different way in the presence of inhomogeneities. Due to the quadratic source term two branches of evolution emerge, one of them physical, with positive energy density $\\rho $. At early times of the cosmological evolution, when the quadratic term dominates, the energy density is less than in the general relativistic case and a negative pressure (tension) $p$ characterizes the cosmological fluid. For $\\Lambda =0$ and after the quadratic source term becomes subdominant, the cosmological fluid evolves asymptotically into the Einstein-Straus dust. The energy density becomes indistinguishable from the general relativistic value and the tension disappears.  Interesting situations occur, whenever $\\Lambda \\neq 0$. With increasing $% \\Lambda $, the evolution of the cosmological fluid changes dramatically.\\ For high enough values of $\\Lambda $ the fluid evolves into a pressure singularity, resembling a sudden future singularity \\cite{Dabrowski}, however here the scale factor $a$ and all of its derivatives stay regular.% \\footnote{% We note that new type of singularities are not unusual in brane-worlds. For example `quiescent'\\ cosmological singularities were described in Ref. \\cite% {ShtanovSahni}, for which the matter density and Hubble parameter stay finite, while all higher derivatives of the scale factor diverge.} This unpleasant behaviour of the fluid can be evaded by redefining the pressure and energy density on the cosmological brane regions in such a way that they absorb $\\Lambda $ \\cite{SwissCheese}. This redefinition generalizes a similar procedure in general relativity.  The simple Swiss-cheese brane-world model presented in \\cite{SwissCheese} can be generalized in many ways. One possibility would be to allow for brane black hole solutions with tidal charge, known both in a static \\cite{tidalRN} and in the rotating stationary axisymmetric case \\cite{Aliev}.\\footnote{% Kerr - anti de Sitter black holes in four and higher space-time dimensions were also recently examined \\cite{Aliev2}. Stars and gravitational collapse on the brane were considered in \\cite{BGM}-\\cite{BraneOppSny2}.}  Another possibility is by allowing for an asymmetric embedding of the brane. While both in the original RS model and in some of its curved generalizations the symmetry of the embedding is assumed, such that the brane is the moving boundary of the bulk, there are many attempts to lift this symmetry, as interesting new features emerge, like a late-time acceleration (see for example \\cite{Induced}). An asymmetric embedding can arise from different black hole masses on the two sides on the brane \\cite% {Kraus}, \\cite{Ida}, \\cite{Davis}, different cosmological constants on the left and right side of the brane \\cite{Deruelle}, \\cite{Perkins} or from both \\cite{Stoica}, \\cite{BCG}, \\cite{Carter}. The asymmetric embedding in brane-worlds was discussed covariantly in \\cite{Decomp} and in the presence of induced gravity and Gauss-Bonnet contributions in \\cite{Induced} and \\cite% {Deruelle}, \\cite{Konya}, respectively.  In line with the above mentioned references, in this paper we propose to discuss \\textit{asymmetrically embedded Swiss-cheese brane-worlds}. In the simplest case the asymmetry is achieved by choosing different values of the bulk cosmological constant in the left and right AdS5 regions in which the FLRW regions of the brane are embedded.  The bulk extension of any Schwarzschild void is left symmetric, therefore the black string / cigar metric is not altered with respect to the symmetric case presented in Ref. \\cite{SwissCheese}. It is reasonable to assume then that the shape of the transition regions connecting the black string / cigar regions to the left or right AdS5 regions is parity-dependent (see Fig \\ref% {Fig0}). The boundaries of these transition regions should be set such that the junction conditions in the bulk are satisfied.  The basic dynamical equation in asymmetrically embedded brane-worlds is an effective Einstein equation (derived in \\cite{Decomp} as a generalization of the result of \\cite{SMS}): \\begin{equation} G_{ab}=-\\Lambda g_{ab}+\\kappa ^{2}T_{ab}+\\widetilde{\\kappa }^{4}S_{ab}-% \\overline{\\mathcal{E}}_{ab}+\\overline{L}_{ab}^{TF}+\\overline{\\mathcal{P}}% _{ab}\\ .  \\label{modEgen} \\end{equation}% The 4-dimensional (4D) and 5D gravitational coupling constants $\\kappa ^{2}$ and $\\widetilde{\\kappa }^{2}$ are related as $6\\kappa ^{2}=\\widetilde{\\kappa }^{4}\\lambda $ and an overbar denotes the average taken over the two sides of the brane. The source term $\\overline{\\mathcal{P}}_{ab}$ (generating in special cases a 'comoving mass' and a bulk pressure \\cite{AT}, \\cite{AT1}) arises from non-standard model bulk fields, like scalar, dilaton, moduli or quantum radiation fields. For simplicity we choose $\\overline{\\mathcal{P}}% _{ab}=0$. We also drop the electric part of the Weyl curvature $\\overline{% \\mathcal{E}}_{ab}$ of the bulk, motivated by the existence of stable black string solutions with vanishing electric part of the Weyl curvature in the two-brane models of Ref. \\cite{RS1}.  The FLRW regions are filled with a perfect fluid $T_{ab}=\\rho \\left( \\tau \\right) u_{a}u_{b}+p\\left( \\tau \\right) a^{2}h_{ab}$ (where $\\tau $ is the comoving time, $u^{a}=\\left( \\partial /\\partial \\tau \\right) ^{a}$ and $% h_{ab}$ the 3-metric with constant curvature of the spatial hypersurfaces $% \\tau $=const), which induces the quadratic source term \\begin{equation} \\widetilde{\\kappa }^{4}S_{ab}=\\kappa ^{2}\\frac{\\rho }{\\lambda }\\left[ \\frac{% \\rho }{2}u_{a}u_{b}+\\left( \\frac{\\rho }{2}+p\\right) a^{2}h_{ab}\\right] \\ . \\label{fS} \\end{equation}% $\\Lambda $ contains a true constant $\\Lambda _{0}$ and a contribution from the asymmetric embedding:% \\begin{eqnarray} \\Lambda &=&\\Lambda _{0}-\\frac{\\overline{L}}{4}~,  \\nonumber \\\\ \\Lambda _{0} &=&\\frac{\\kappa ^{2}\\lambda +\\widetilde{\\kappa }^{2}\\overline{% \\widetilde{\\Lambda }}}{2}~.  \\label{lambda2} \\end{eqnarray}% By applying the definitions given in \\cite{Decomp} the asymmetry source terms arise as% \\begin{eqnarray} \\overline{L} &=&-\\frac{3\\left( 5\\rho +3p+2\\lambda \\right) \\left( \\Delta \\widetilde{\\Lambda }\\right) ^{2}}{8\\left( \\rho +\\lambda \\right) ^{3}}\\ , \\label{Lbar2} \\\\ \\overline{L}_{ab}^{TF} &=&\\frac{-3\\left( \\rho +p\\right) \\left( \\Delta \\widetilde{\\Lambda }\\right) ^{2}}{32\\left( \\rho +\\lambda \\right) ^{3}}\\left( 3u_{a}u_{b}+a^{2}h_{ab}\\right) \\ \\end{eqnarray}% Therefore the Friedmann and Raychaudhuri equations (the non-trivial parts of the effective Einstein equation on Friedmann branes) are changed by asymmetry.  The plan of the paper is the following. In section 2 we discuss the evolution of the scale factor in this asymmetric Swiss-cheese brane-world as derived from the junction conditions on the brane. For $k=0$ we show that the asymmetry does not modify the cosmological evolution. This occurs exactly as in the symmetric case (which in turn is the cosmological evolution of the Einstein-Straus model). Therefore the evolution of $\\rho $ and $p$ has to change.  In section 3 we study the evolution of the fluid. Compared to the symmetric case, the two branches of the symmetric Swiss-cheese model split into four. This happens, because both source terms $\\overline{L}$, $\\overline{L}% _{ab}^{TF}\\propto \\rho ^{-2}$, while $S_{ab}\\propto \\rho ^{2}$, therefore yielding to a quartic polynomial in the effective source. The higher the asymmetry, the more important are the induced modifications as compared to the symmetric case. In order to quantify this, we define a properly chosen dimensionless asymmetry parameter $\\alpha $. Then we show that for any $% \\Lambda \\leq \\kappa ^{2}\\lambda $ $/2$ (including $\\Lambda =0$) there is a \\textit{critical value} $\\alpha _{crit}$, which separates cosmologies with pressure singularities from those with regular evolution during the whole lifetime of the universe. We find that (i) pressure singularities can appear even when $\\Lambda =0$, provided the asymmetry in the embedding is high (ii) pressure singularities are generic for $\\Lambda >\\kappa ^{2}\\lambda /2$. While in the symmetric case there was a single branch allowing for positive energy density, in the asymmetric case there are two. We study the domains of positivity of the energy density on these two branches. Then we illustrate graphically various typical behaviors of $\\rho $ and $p$.  In section 4 we discuss in detail the models with $\\Lambda =0$, which can be obtained after a suitable redefinition of the fluid variables, as proven in Ref. \\cite{SwissCheese}. In subsection 4.a we illustrate the critical behavior of the Swiss-cheese brane-world models with asymmetry by plotting the evolution of the four branches of energy density as function of both time and asymmetry parameter. As $\\Lambda =0$ is chosen, the pressure singularity (coming together with the ill-definedness of the energy density) arises as an exclusive consequence of the asymmetric embedding. In subsection 4.b we introduce the \\textit{pressure singularity censorship} conjecture. This sets limits on the allowable degree of asymmetry in the model. We show that with the highest established lower limit for the brane tension, $\\Delta \\widetilde{\\Lambda }$ can be of the order of the individual values of the bulk cosmological constant $\\widetilde{\\Lambda }_{L,R}$ on the two sides of the brane. Thus in particular the pressure singularity censorship can be satisfied with AdS5 on one side of the brane and Minkowski bulk on the other if the brane tension is high enough. \\  In section 5 we estimate the energy scale for which the asymmetry becomes important during cosmological evolution. We find that in order to satisfy both the nucleosynthesis constraints and the pressure singularity censorship, the minimal value of $\\lambda $\\ derived from the nucleosynthesis constraints has to be supplemented by a $10$ times larger maximal value. It is a generic feature, independent of the Swiss-cheese model considered here that the asymmetry source term of the Friedmann equation generated by different bulk cosmological constants bounds the brane tension from above. Finally we analyze the degree of asymmetry and the likelihood of a cosmological evolution running into a pressure singularity in our universe.  Section 6 contains the Concluding Remarks. ",
        "conclusions": "In this paper we have discussed brane cosmological models with local inhomogeneities (Swiss-cheese brane) in which the brane is embedded asymmetrically into the bulk. The asymmetry is generated by different cosmological constants in the left and right anti de Sitter bulk regions. As the junction conditions on the brane are unchanged by the embedding of the brane, the evolution of the scale factor proceeds exactly as in the symmetric case (given explicitly, both in Ref. \\cite{SwissCheese} and here, for branes with inhomogeneities represented by Schwarzschild black holes and FLRW regions with flat spatial sections). The cosmological evolution in fact proceeds as in the general relativistic Einstein-Straus model, which forever expands and decelerates. The reason is that the totality of source terms (including the asymmetry source term $\\overline{L}_{ab}^{TF}$) in the effective Einstein equation (\\ref{modEgen}) add up to an effective dust source with \\begin{equation} \\rho ^{tot}=\\frac{4}{3\\kappa ^{2}\\tau ^{2}},\\ \\ \\ p^{tot}=0\\ ,  \\label{dust1} \\end{equation}% establishing a formal equivalence with the general relativistic Einstein-Straus model.  Due to the modifications induced by asymmetry in the Friedmann and Raychaudhuri equations, the evolution of the fluid in the asymmetric brane-world scenario is however changed. In the symmetric case the modifications were caused by the source term quadratic in the energy-momentum. As consequence of this quadratic source term, two branches $% \\rho _{\\pm }$ emerged (one of them with positive energy density). With a cosmological constant $\\Lambda $ in the FLRW regions of the brane, the fluid could evolve into a pressure singularity in spite of the regular evolution of the scale factor. This unphysical pressure singularity (and also the evolution into negative energy density) could be conveniently avoided by a proper redefinition of the energy density and pressure of the fluid \\cite% {SwissCheese} such that the new fluid variables absorb the difference in the brane cosmological constants.  The asymmetry between the left and right bulk regions further splits up the possible evolutions of the fluid. With asymmetry, there are four possible branches $\\rho _{\\pm \\pm }$. Only two of them allow for positive energy density and only one of them has well defined limit in the symmetric case. In certain parameter range the evolution of the fluid is regular and approaches at late times the Einstein-Straus universe filled with dust.  However with increasing degree of asymmetry in the left and right bulk cosmological constants, the fluid can evolve into an unphysical negative energy density and / or a pressure singularity, while the scale factor stays regular.\\ This can happen even if $\\Lambda =0$ is chosen, due entirely to the asymmetry in the embedding.  Based on the analysis of this paper and of Ref. \\cite{SwissCheese}, we conclude that pressure singularities could appear as consequence of the difference of the cosmological constants either between the vacuum and FLRW regions (symmetric case) or between the left and right bulk regions with respect to the brane (asymmetric case). When both are present, as a rule, we have seen that asymmetry speeds up the evolution into a pressure singularity.  For generic values of $\\Lambda $ we have introduced a \\textit{critical value} $\\alpha _{crit}$ of a suitably defined asymmetry parameter, given by Eq. (% \\ref{alphacrit}), which separates Swiss-cheese brane-worlds with and without pressure singularities.  We have also analyzed whether such pressure singularities occur in universes similar to ours. By carefully examining the various lower bounds on the brane tension we have proved that the observed value of the cosmological constant is much less than the minimal value of $\\Lambda $ necessary for pressure singularity occurrence in the symmetric case.  In the asymmetric case we have imposed the \\textit{pressure singularity censorship}. This conjecture bounds the possible degree of asymmetry. For the smallest brane tension $\\lambda $ obeying all known constraints, the asymmetry in the bulk cosmological constant can be at most $67\\%$. With higher values of $\\lambda $ the asymmetry can further increase up to $400\\%$ without jeopardizing the regular evolution of the fluid.  We have shown that the nucleosynthesis constraints restrict the brane tension not only from below, as well-known, but also from above, due to the asymmetry source term. The upper bound decreases with $\\lambda $. With the highest asymmetry allowed by the pressure singularity censorship, the upper bound for the brane tension is $10$ times the lower bound.  However for our universe present cosmological observations restrict the value of $\\alpha $ much below the critical value. Therefore the upper bound is $\\infty $ for all practical purposes and the evolution of the fluid on the Swiss-cheese brane representing our universe stays regular. With the upper limit for $\\lambda $ very high, the astrophysical lower limit for $% \\lambda $ can also be obeyed, thus the brane-world universe with the cosmological constant completely replaced by the asymmetry source term can also accommodate compact objects, like neutron stars.  \\ack L\\'{A}G was supported by OTKA grants nos. T046939 and TS044665 and the J% \\'{a}nos Bolyai Fellowship of the Hungarian Academy of Sciences."
    },
    "0608/astro-ph0608511_arXiv.txt": {
        "abstract": "A number of recent studies have estimated the inter-galactic void probability function and investigated its departure from various random models. We study a family of parametric statistical models based on gamma distributions, which do give realistic descriptions for other stochastic porous media. Gamma distributions contain as a special case the exponential distributions, which correspond to the `random' void size probability arising from Poisson processes. The space of parameters is a surface with a natural Riemannian metric structure. This surface contains the Poisson processes as an isometric embedding and a recent theorem~\\cite{gamran} shows that it contains neighbourhoods of all departures from randomness. The method provides thereby a geometric setting for quantifying departures from randomness and on which may be formulated cosmological evolutionary dynamics for galactic clustering and for the concomitant development of the void size distribution.  The 2dFGRS data~\\cite{croton04a,croton04b} offer the possibility of more detailed investigation of this approach than was possible when it was originally suggested~\\cite{vpf,ijtp,gsis} and some parameter estimations are given.  {\\bf Key words:} galaxies: clustering, statistics; cosmology: voids, statistics, parametric models, information geometry. ",
        "introduction": "Several years ago, the author presented an information geometric approach to modelling a space of perturbations of the random state for galactic clustering and cosmological void statistics~\\cite{vpf}.  The present note is intended to update somewhat and draw attention to this approach as a possible contribution to the interpretation of the data from the 2-degree field Galaxy Redshift Survey (2dFGRS), cf Croton et al.~\\cite{croton04a,croton04b}.  The classical random model is that arising from a Poisson process of mean density $\\overline{N}$ galaxies per unit volume in a large box. Then, in a region of volume $V,$ the probability of finding exactly $m$ galaxies is \\begin{equation} P_m =\\frac{(\\overline{N}V)^m}{m!} e^{-\\overline{N}V} \\label{poisson} \\end{equation} So the probability that the given region is devoid of galaxies is $P_0=e^{-\\overline{N}V}.$ It follows that the probability density function for the continuous random variable $V$ in the Poisson case is \\begin{equation} p_{random}(V) = \\overline{N} \\, e^{-\\overline{N}V} \\label{negexp} \\end{equation} In practice of course, measurements will depend on algorithms that specify threshold values for density ranges of galaxies in cells and the lowest range will represent the `underdense' regions which include the voids; Benson et al.~\\cite{benson} discuss this.  A hierachy of $N$-point correlation functions needed to represent clustering of galaxies in a complete sense was devised by White~\\cite{white} and he provided explicit formulae, including their continuous limit. In particular, he made a detailed study of the probability that a sphere of radius $R$ is empty and showed that formally it is symmetrically dependent on the whole hierarchy of correlation functions. However, White concentrated his applications on the case when the underlying galaxy distribution was a Poisson process, the starting point for the present approach which is concerned with geometrizing the parameter space of departures from a Poisson process. Croton et al.~\\cite{croton04b} found that the negative binomial model for galaxy clustering gave a very good approximation to the 2dFGRS, pointing out that this model is a discrete version of the gamma distribution. ",
        "conclusions": ""
    },
    "0608/astro-ph0608382_arXiv.txt": {
        "abstract": "We present a 200 $ksec$ {\\sl Chandra} observation of seven spectroscopically selected, high redshift ($0.75 < z < 1.03$) galaxy groups and clusters discovered by the DEEP2 Galaxy Redshift Survey in the Extended Groth Strip (EGS). X-ray emission at the locations of these systems is consistent with background. The 3$\\sigma$ upper limits on the bolometric X-ray luminosities ($L_X$) of these systems put a strong constraint on the relation between $L_X$ and the velocity dispersion of member galaxies $\\sigma_{gal}$ at $z\\sim 1$; the DEEP2 systems have lower luminosity than would be predicted by the local relation. Our result is consistent with recent findings that at high redshift, optically selected clusters tend to be X-ray underluminous.  A comparison with mock catalogs indicates that it is unlikely that this effect is entirely caused by a measurement bias between $\\sigma_{gal}$ and the dark matter velocity dispersion. Physically, the DEEP2 systems may still be in the process of forming and hence not fully virialized, or they may be deficient in hot gas compared to local systems. We find only one possibly extended source in this  {\\sl Chandra} field, which happens to lie outside the DEEP2 coverage. ",
        "introduction": "Groups and clusters of galaxies are the most massive, dynamically relaxed objects in the universe. Past studies have focused on low to moderate redshifts (for reviews, see, e.g., \\citealp{ros02,voi05}) due to the lack of reliable samples at high redshift. Recently, progress has been made in finding high-$z$ clusters and groups, largely due to improved detection techniques and greater instrument sensitivities over large fields. For instance, the method of observing red sequence galaxies has proven to be efficient in identifying clusters with optical/near-IR imaging data (e.g., \\citealp{gye00}). Also, {\\sl Chandra} and {\\sl XMM}-Newton have started to reveal X-ray clusters at redshift around unity (e.g.,\\citealp{ros02,voi05}). However, until now, no optically selected, spectroscopic sample has existed with sufficient size, sampling density, and redshift accuracy to identify large numbers of groups and clusters at $z > 0.5$.  \\begin{figure*}[thb] \\centerline{\\includegraphics[angle=0,scale=0.8]{f1.eps}} \\caption[h]{True-color image of the {\\sl Chandra} ACIS-I field, with north to the top and east to the left. The image is color-coded so that photons with energies at 0.54 -- 2.05 keV, 2.05 -- 3.50 keV, and 3.50 -- 7.00 keV are in red, green, and blue, respectively. The DEEP2 survey covers the region to the north of the solid line. Circles represent DEEP2-identified systems. The white box labeled ``A'' indicates a possibly extended source.} \\label{f1} \\end{figure*}  The DEEP2 Galaxy Redshift Survey \\citep{dav03} has provided the first large spectroscopic galaxy catalog focused on $z\\sim1$, based on observations of $\\sim 5 \\times 10^4$ galaxies. One of its four survey fields incorporates the Extended Groth Strip (EGS), which is the target of a multiwavelength survey called ``All-wavelength Extended Groth Strip International Survey'' (AEGIS; \\citealp{dav06}). DEEP2 has produced the largest well-defined, spectroscopically selected sample of clusters and groups at high $z$ yet \\citep{ger05}.  In this {\\it  Letter}, we investigate the X-ray properties of the DEEP2 groups and clusters which overlap an archival 200 $ksec$ {\\sl Chandra} observation. The primary goal is to understand the relation between optical and X-ray properties of these systems at $z>0.7$. At low $z$, simple scaling relations between X-ray quantities such as bolometric luminosity ($L_X$) and temperature ($T_X$) and optical quantities such as galaxy velocity dispersion ($\\sigma_{gal}$) have been predicted and well studied (see, e.g., \\citealp{whi97,all98,mar98,fin01,ett02}). However, high-$z$ observations reveal a different trend: optically selected galaxy clusters tend to be less X-ray luminous than those with similar velocity dispersions at low redshifts (\\citealp{cas94,bow97,hol97,lub02,lub04}), while the X-ray selected samples tend to follow the local $L_X$--$\\sigma_{gal}$ relation (see, e.g., \\citealp{ebe01,val04}). The exact reason for the break-down of the $L_X$--$\\sigma_{gal}$ relation is still unclear. While previous studies of optical-X-ray relations have been limited only to a few large clusters, DEEP2 allows us to identify both clusters and groups at $z\\sim 1$. ",
        "conclusions": "In this ${\\it Letter}$,  we find evidence % that spectroscopically selected galaxy groups and clusters at high redshift are X-ray underluminous; i.e., the upper limits of their X-ray luminosities are significantly below predictions from the local $L_X-\\sigma_{gal}$ relation, unlike X-ray selected clusters at these redshifts. This confirms previous results showing that % optically selected clusters at high $z$ are weak X-ray sources.  While bias in the $\\sigma_{gal}$ measurement may contribute partly to the non-detection, it is unlikely to be the sole explanation. Physically, there are two basic reasons why spectroscopically selected systems at $z\\sim1$ might fall below the $L_X-\\sigma_{gal}$ relation followed by X-ray selected clusters at these redshifts.  The first is that the DEEP2 systems may be dynamically younger objects than those clusters which are easily detected in the X-ray.  If the DEEP2 systems are are not yet fully virialized, their gas may tend to be in cooler subclumps rather than at the temperature predicted from spherical-collapse models, and their galaxy velocity distribution may be multimodal rather than Gaussian, increasing the measured $\\sigma_{gal}$ (see, e.g., \\citealp{fre96,bow97,lub04}).  An alternative is that these galaxy groups and clusters may be deficient in the hot gas that is necessary to produce detectable X-ray emission.  DEEP2 systems do appear to be generally above the threshold mass beyond which cooling is predicted to be inefficient and a hot atmosphere forms (\\citealp{bir03,ker05,cro06}). However, they may have passed that threshold only relatively recently in their history and may still be in the process of building a hot halo of substance. In this case $L_X$ would be expected to remain low until a sufficient amount of hot gas had been accreted. We note that the average blue galaxy fraction is higher in DEEP2 groups than in local systems, reflecting the more common presence of the cold gas which feeds star formation \\citep{ger06}. al. 2006, in prep.)  Our {\\sl Chandra} observation is $\\sim 200\\ ksec$. Two deeper ($\\sim 1 Msec$) exposures, the {\\sl Chandra} Deep Field South and North, yield a total of 18 \\citep{gia02} and 6 \\citep{bau02} extended sources, respectively. Most of the sources are early type galaxies, poor groups and clusters. However, an early analysis on the {\\sl Chandra} Deep Field North with similar exposure did not detect any extended emission \\citep{hor01}.   Our study highlights the importance of understanding groups and clusters of galaxies at high redshift. These systems present very different properties when compared with their counterparts at low redshift, indicating that complex processes govern the formation and evolution of large scale structures, which in turn have significant impact on galaxy evolution.  We plan to investigate larger DEEP2 samples in the future to advance this work.  {\\it Acknowledgements:} We thank Martin White, Joanne Cohn and Sandy Faber for useful discussions and comments. TF was supported by the NASA through {\\sl Chandra} Postdoctoral Fellowship Award Number PF3-40030 issued by the {\\sl Chandra} X-ray Observatory Center, which is operated by the Smithsonian Astrophysical Observatory for and on behalf of the NASA under contract NAS~8-39073. Also see Davis et al.~(2006) in this volume for full acknowledgements."
    },
    "0608/astro-ph0608661_arXiv.txt": {
        "abstract": "From observations with the GMOS multi-slit spectrograph on the Gemini North telescope, we have obtained spectra for 39 globular cluster candidates in the Virgo giant elliptical galaxy NGC 4649 (M60), of which 38 are confirmed globular clusters. The clusters extend out to a radius of 260\\arcsec\\ (3.5 effective radii). We find no rotation of the globular cluster system, with an upper limit of v/$\\sigma$ $<$ 0.6 at a confidence level of 95\\%. The globular cluster velocity dispersion is constant with radius, within the uncertainties.  We fit isotropic models to the globular cluster and stellar kinematics; these models yield a M/L$_V$ around 16 at 200\\arcsec\\ radius (16 kpc), an increase of a factor of two from the central M/L. We also use the mass profile as derived from X-rays to determine the orbital structure. Using axisymmetric orbit-based models and the X-ray mass profile, we find the orbital distribution is close to isotropic within 100\\arcsec, and becomes tangentially biased beyond.  Furthermore, when using the X-ray profile, we find a better fit to the kinematics compared to using a constant M/L model. Thus, both isotropic and axisymmetric orbit-based models give support for the presence of a dark matter halo in NGC 4649. ",
        "introduction": "\\label{Introduction}  Globular clusters (GCs hereafter) are an excellent way to study the formation, evolution, and dark matter (DM) content of galaxies of all types (e.g. Ashman \\& Zepf 1998; Harris 2001). They are particularly well-suited for dynamical studies of early-type galaxies, since they are found in substantial numbers out to large radii around such galaxies and are bright enough to be seen at large distances. GCs complement other dynamical probes such as the stellar light, planetary nebulae (PNe), and X-ray data (e.g. Gerhard 2005; Romanowsky 2005). Integrated stellar spectroscopy has the advantage of large sample sizes and high signal-to-noise, which allow the derivation of higher-order moments of the velocity distribution, but the faintness of the stellar halo restricts measurement to within a few effective radii of galaxy centers (e.g. Kronawitter et al. 2000; Gerhard et al. 2001). X-ray data can probe further out, and recent high-quality data from ROSAT, ASCA, and {\\it Chandra} have shown that extended dark halos are a common feature of all luminous ellipticals (Mathews \\& Brighenti 2003; Humphrey et al. 2006). However, accurate X-ray mass determinations require high-quality gas density and temperature profiles, along with the assumption of hydrostatic equilibrium. As well, only massive ellipticals have hot diffuse gas halos for mass estimates.  PNe also have extended spatial distributions, and are easily identified because of their strong emission in the [OIII] 4959/5007 lines. Romanowsky et al. (2003), using the Planetary Nebula Spectrograph (Douglas et al. 2002), found that three intermediate-luminosity ellipticals (NGCs 821, 3379, and 4494) have declining PNe velocity dispersion profiles, consistent with simple models with little or no DM.  This result is in conflict with the conventional model of galaxy formation, whereby galaxies form in dark matter halos. However, numerical simulations by Dekel et al. (2005) show that spiral-spiral mergers may lead to halo stars having highly radial orbits, as they are tidally ejected from the inner regions during the merger. Observed PNe in ellipticals, as part of this stellar halo, would then have line-of-sight velocity dispersions lower than that of the underlying DM. Older tracers such as GCs are expected to have higher velocity dispersions, more representative of the DM. In support of this, Sambhus, Gerhard \\& Mendez (2006) have recently found that the bright and faint PNe in NGC 4697 have different kinematics, and suggest that the bright PNe are a younger population perhaps formed in tidal structures. Our group (Pierce et al. 2006a) has found that the GC velocity dispersion profile in NGC 3379 increases with radius, and non-parametric, isotropic models yield a significant increase in the M/L ratio at large radii. A similar result for the NGC 3379 GCs has also been found by Bergond et al.(2006), from FLAMES/VLT data. Note also that Peng, Ford \\& Freeman (2004a) found a relatively low M/L ratio of 10$-$15 for NGC 5128 from PNe velocities. It is thus very important to compare GC and PNe kinematics in a larger number of galaxies, to gain a more complete understanding of host galaxy dynamics.  GC velocities have been used to study the dynamics of several early-type galaxies, including M87 (Cohen 2000; Cote et al. 2001), NGC 4472 (Sharples et al. 1998; Zepf et al. 2000; Cote et al. 2003), NGC 1399 (Richtler et al. 2004), NGC 5128 (Peng, Ford \\& Freeman 2004b), and NGC 3379 (Pierce et al. 2006a; Bergond et al. 2006); see Zepf (2003) for a review. From these studies it is consistently found that the GC velocity dispersions are either constant or increasing with radius. This is strong evidence for DM halos in these galaxies, supporting X-ray data. Merger simulations (Hernquist \\& Bolte 1993; Vitvitska et al. 2002; Bekki et al. 2005) predict that ellipticals should have significant amounts of angular momentum at large radius. However, rotation is seen in some GC systems (e.g. NGC 5128; M87), but not in others (e.g. NGC 1399; NGC 3379; NGC 4472). In some cases, the lack of detected rotation may be due to small numbers of GCs; but NGC 4472 and NGC 1399 have hundreds of GC velocities, allowing tight constraints on rotation except at the outer reaches of the datasets. The situation is even more complex when the kinematics of metal-poor (MP) and metal-rich (MR) GCs are studied separately. In most GC systems, the MP and MR GCs are different in terms of their rotation and velocity dispersions, and these differences often depend on galactocentric radius. Usually the MP GCs are dynamically hotter than the MR GCs, as might be expected given their generally shallower spatial profiles (e.g. NGC 5128, NGC 1399, NGC 4472). In some cases (e.g. NGC 4472, M87) the blue and red clusters may be counter-rotating at some radii. In only a few galaxies with the largest GC samples can we say anything about GC orbital anisotropy; in NGC 1399, NGC 4472, and M87 the GCs are consistent with isotropic orbits (though there are hints that the orbital behavior may be different for MP and MR sub-samples).  This paper continues our Gemini/GMOS study of the GC system of NGC 4649 (M60), a bright elliptical located in a subcluster in the eastern part of the Virgo cluster (see Forbes et al. 2004 for further details of NGC 4649 and photometric studies of its GC system). NGC 4649 is well-studied in X-rays, with ROSAT (Trinchieri, Fabbiano \\& Kim 1997; Di Matteo \\& Fabian 1997; Bohringer et al. 2000; O'Sullivan, Forbes, \\& Ponman 2001) and {\\it Chandra} (Irwin, Athey \\& Bregman 2003; Sarazin et al. 2003; Swartz et al. 2004; Randall, Sarazin \\& Irwin 2004; Gutierrez \\& Lopez-Corredoira 2005; Humphrey et al. 2006) data. Trinchieri, Fabbiano \\& Kim (1997), Bohringer et al. (2000), and Humphrey et al. (2006) all find evidence for an extended dark matter halo in NGC 4649; Humphrey et al. (2006) find an enclosed mass of ~1 $\\times$ 10$^{12}$ M$_\\odot$ and a M/L $\\sim$ 5 in the K band at 260\\arcsec\\ (20 kpc). {\\it Chandra} data have also revealed over 100 low mass x-ray binaries (LMXRBs) in NGC 4649 (Sarazin et al. 2003; Randall, Sarazin \\& Irwin 2004), and we will use our larger GMOS GC imaging database (Forbes et al. 2004) to match with the {\\it Chandra} LMXRBs in a future paper (Bridges et al. 2006, in preparation).  Previous studies of the stellar kinematics in NGC 4649 have found major axis stellar rotation of 100$-$120 km/sec out to 45$-$100\\arcsec, a high central velocity dispersion of 350$-$400 km/sec, and a declining velocity dispersion profile (Fisher, Illingworth \\& Franx 1995; de Bruyne et al. 2001; Pinkney et al. 2003); de Bruyne et al. also find modest minor axis rotation of $\\sim$ 30 km/sec. de Bruyne et al. use 3-integral axisymmetric models to find that a dark matter halo gives the best fit to the kinematic data, though a constant M/L ratio model is also consistent with the data. We will use the data from Pinkney et al. to combine with our GC velocities and X-ray data (Humphrey et al. 2006) to study the dynamics of NGC 4649 out to large radius (Section \\ref{mlratio}).  Forbes et al. (2004) presented GMOS photometry for $\\sim$ 1000 GC candidates, and confirmed previous findings of bimodality in the GC colour distribution.  Forbes et al. also found that the red GCs are more concentrated to the galaxy centre than the blue GCs, as is found in many other galaxies. The underlying galaxy light has a similar density profile slope as the red GCs. These GMOS images were used to select GC candidates for follow-up spectroscopy using GMOS in multi-slit mode (see Section~\\ref{selection}). This paper presents spectra for 38 confirmed NGC 4649 GCs, and discusses the GC kinematics (rotation and velocity dispersion profiles; Section~\\ref{kinematics}) and dark matter content of NGC 4649 (Section~\\ref{mlratio}). In Section ~\\ref{discussion} we discuss our main results in more detail, while Section~\\ref{conclusions} presents our conclusions. A companion paper (Pierce et al. 2006b) uses spectroscopic line indices to address the ages, abundances, and abundance ratios of the confirmed GCs. We adopt an SBF distance of D = $17.29^{+0.58}_{-0.56}$ Mpc for NGC 4649 (Mei et al. 2006, in preparation), taken from the ACS Virgo Cluster Survey (Cote et al. 2004); this gives an image scale of 11.9\\arcsec/kpc. This is slightly larger than the distance of 16.8 Mpc we used in our photometric study of the NGC 4649 GCs (Forbes et al. 2004). ",
        "conclusions": "\\label{conclusions}  We have used Gemini/GMOS to obtain spectra for 38 confirmed globular clusters in NGC 4649. The recession velocities of these clusters have been used to study the cluster kinematics and dark matter content of NGC 4649. Our main results are:  \\begin{itemize}  \\item The GC velocity dispersion profile is constant with radius, for all GCs, and also for the blue and red GCs separately.  \\item We detect no rotation in the GC system, both for all clusters, and for blue and red samples separately. We are able to place upper limits on the ratio of rotation velocity to velocity dispersion v/$\\sigma$ $<$ 0.6 for all GCs. Further data to improve our spatial and azimuthal coverage, and our sample size, are needed to make further progress.  \\item Both spherical, isotropic and axisymmetric, orbit-based dynamical models strongly support the presence of a dark matter halo in NGC 4649. For the isotropic models there is excellent agreement between the stellar plus GC M/L profile and the X-ray M/L profile from Humphrey et al. (2006).  \\item Within $\\sim$ 100\\arcsec\\ radius, the GC orbits are close to isotropic, while at larger radius the orbits become tangentially biased.  \\end{itemize}"
    },
    "0608/astro-ph0608457_arXiv.txt": {
        "abstract": "We estimate the average group morphological and dynamical characteristics of the Percolation-Inferred Galaxy Group (2PIGG) catalogue within $z\\le 0.08$, for which the group space density is roughly constant. We quantify the different biases that enter in the determination of these characteristics and we devise statistical correction procedures to recover their bias free values. We find that the only acceptable morphological model is that of prolate, or triaxial with pronounced prolatness, group shapes having a roughly Gaussian intrinsic axial ratio distribution with mean $\\sim 0.46$ and dispersion of $\\sim 0.16$. After correcting for various biases, the most important of which is a redshift dependant bias, the median values of the virial mass and virial radius of groups with 4 to 30 galaxy members, is: ${\\overline M}_v \\sim 6 \\times 10^{12} \\; h_{72}^{-1} M_{\\odot}$, ${\\overline R}_v \\sim 0.4 \\;h^{-1}_{72}$ Mpc, which are significantly smaller than recent literature values that do not take into account the previously mentioned biases. The group mean crossing time is $\\sim 1.5$ Gyr's, independent of the group galaxy membership. We also find that there is a correlation of the group size, velocity dispersion and virial mass with the number of group member galaxies, a manifestation of the hierarchy of cosmic structures. ",
        "introduction": "Groups of galaxies are the lowest level cosmic structures, after galaxies themselves, in the hierarchy  that leads to the largest  virialized structures, the clusters of galaxies. There have been a number of recent attempts to construct objectively selected group and cluster catalogues from magnitude limited redshift or photometric galaxy surveys, like the combination of the Updated Zwicky Catalogue and the Southern Sky Redshift Survey (UZC-SSRS2), the Sloan digital sky survey (SDSS), the two-degree field redshift survey (2dFRGS),  the digital Palomar observatory sky survey; eg.,  Ramella et al (2002); Merch\\'an \\& Zandivarez (2002, 2005); Gal et al. (2003); Bachall et al. (2003); Goto et al. (2004); Lee et al. (2004); Lopes et al. (2004); Eke et al. (2004); Tago et al (2006); Berlind et al (2006). Most of the studies that use galaxy redshift information, apply the so-called friends-of-friends algorithm (FoF) to the galaxy redshift data by using a variable linking length that takes into account the drop of the redshift selection function with increasing redshift.  It appears that most galaxies are found in groups and they are therefore extremely important in our attempts to understand the cosmic structure formation processes. Since virialization will tend to sphericalize initial anisotropic distributions of matter, the shape of different cosmic structures is an indication of their evolutionary stage. Furthermore, the group shape, size and velocity dispersion are important factors in determining galaxy member orbits and interaction rates, which are instrumental in understanding galaxy evolution processes. Alot of theoretical and observational studies have dealt with the intrinsic shape of cosmic structures and their dependence on different cosmological backgrounds, environments, evolutionary stage, etc. (e.g., Carter \\& Metcalfe 1980; Plionis, Barrow \\& Frenk 1991; Cooray 2000; Basilakos, Plionis \\& Maddox 2000; Zeldovich, Einasto \\& Shandarin 1982; de Lapparent, Geller \\& Huchra 1991; Plionis, Valdarnini \\& Jing 1992; Oleak et al. 1995; de Theije, Katgert \\& van Kampen 1995, Jaaniste et al. 1998; Sathyaprakash et al. 1998; Valdarnini, Ghizzardi \\& Bonometto 1999; Basilakos, Plionis \\& Rowan-Robinson 2001; Jing \\& Suto 2002; Kasun \\& Evrard 2005; Allgood et al. 2006; Paz et al. 2006; Sereno et al. 2006). In the case of groups of galaxies a study of the UZC-SSRS2 group catalogue (Plionis, Basilakos \\& Tovmassian 2004) have shown that they are quite elongated prolate-like systems.  In this paper we use the recently constructed 2PIGG group catalogue (Eke et al. 2004), based on the 2-degree field redshift survey, to estimate the group projected and intrinsic shape, size, velocity dispersion, virial mass and crossing-time distributions after correcting for a number of systematic biases. ",
        "conclusions": "Using Monte-Carlo and N-body simulations we have investigated the different biases that enter in the determination of the 2PIGG morphological and dynamical characteristics and we have devised statistical correction procedures to recover their bias free values.  Within a redshift that the 2PIGG groups have a roughly constant space number density ($z\\le 0.08$), we derived the average morphological and dynamical characteristics (size, velocity dispersion, virial mass and crossing times), which we find to be significantly smaller than those of other recent studies, exactly because we have taken into account the redshift-dependant bias. Within $z\\le 0.08$, the median value of the  group virial radius, virial mass and crossing time is $\\sim 6 \\times 10^{12} \\; h^{-1}_{72} M_{\\odot}$, $\\sim 0.4 \\; h^{-1}_{72}$ Mpc and $\\sim 1.5$ Gyr's, respectively.  Assuming that groups constitute a homogeneous spheroidal population, we numerically invert the projected, discreteness free, axial ratio distribution to obtain the corresponding intrinsic one. The only acceptable model is that of a prolate, or a triaxial with pronounced prolatness, distribution with a mean axial ratio of $\\beta\\simeq 0.46$ and a dispersion of $\\sim 0.16$.  We have also found that there is a correlation of the size and the velocity dispersion with the number of group members, which is probably a natural outcome of the hierarchy of cosmic structures."
    },
    "0608/astro-ph0608112.txt": {
        "abstract": "{On the order of 20\\% of the very metal-poor stars in the Galaxy exhibit large carbon enhancements. It is important to establish which astrophysical sites and processes are responsible for the elemental abundance patterns of this early Galactic population. } {We seek to understand the nature of the progenitors of three main-sequence turnoff Carbon-Enhanced Metal-Poor (CEMP) stars, \\csthree, \\cstwo, and \\cstwonine, based on a detailed abundance analysis. } {From high-resolution VLT/UVES spectra ($R\\sim 43,000$), we determine abundances or upper limits for Li, C, N, O, and other important elements, as well as $^{12}$C/$^{13}$C isotopic ratios. } {All three stars have $\\rm -3.30 \\le [Fe/H] \\le -2.85$ and moderate to high CNO abundances. \\cstwo\\ is one of the most carbon-rich CEMP stars known ([C/Fe] = +3.2), while \\cstwonine\\ (one of the few N-enhanced metal-poor stars known) is one of the most nitrogen rich ([N/Fe] = +3.0). Oxygen is very high in \\csthree\\ ([O/Fe] = +2.35) and in \\cstwo\\ ([O/Fe] = +1.35). All three stars exhibit [Sr/Fe] $< 0$; Ba is not detected in \\cstwo\\ ([Ba/Fe] $< -0.53$), but it is moderately enhanced ([Ba/Fe] $\\sim 1$) in the other two stars. \\cstwo\\ displays one of the largest sodium overabundances yet found in CEMP stars ([Na/Fe] = +2.8). \\cstwo\\ has $^{12}$C/$^{13}$C = 9, similar to most other CEMP stars without enhanced neutron-capture elements, while $^{12}$C/$^{13}$C $\\ge 40$ in \\csthree. \\csthree\\ and \\cstwonine\\ have A(Li) $\\sim 1.7$, below the Spite Plateau, while Li is not detected in \\cstwo. } {\\cstwo\\ is a CEMP-no star, but the other two stars are in no known class of CEMP star and thus either constitute a new class or are a link between the CEMP-no and CEMP-s classes, adding complexity to the abundance patterns for CEMP stars. We interpret the abundance patterns in our stars to imply that current models for the presumed AGB binary progenitors lack an extra-mixing process, similar to those apparently operating in RGB stars. } % ",
        "introduction": "\\begin{figure}[th] \\rotatebox{0}{\\resizebox{8.5cm}{!}{\\includegraphics{5440fig1.eps}}} \\caption{Observed (crosses) and synthetic spectra (lines) of \\csthree, \\cstwo, and \\cstwonine\\ in the region of the Li~I 670.7 nm doublet. Li is not detected in \\cstwo.} %\\label{} \\end{figure}  A considerable fraction (20-25\\%) of the very metal-poor stars in the Galaxy ([Fe/H] $\\le -2.0$) are strongly enhanced in carbon and often in nitrogen as well ([C,N/Fe] $>$ 1.0; Beers \\& Christlieb 2005; Lucatello et al. 2006). It appears that multiple nucleosynthetic pathways for the production of C in low-metallicity stars must be invoked, almost certainly involving different astrophysical sites. Most stars in the range $-2.6 \\le \\rm{[Fe/H]} \\le -2.0$ appear to be associated with C production in Asymptotic Giant Branch (AGB) companions that have transferred material to the star that is now observed (Aoki et al. 2006a). Even lower-metallicity stars, especially those with [Fe/H] $< -3.0$, may contain C produced by massive, rapidly-rotating stars with [Fe/H] $< -6.0$; see, e.g., Hirschi et al. (2006); Karlsson et al. (2006); Piau et al. (2006). In order to better understand the role of these various possible mechanisms, detailed investigations are needed of the elemental abundances in Carbon-Enhanced Metal-Poor (CEMP) stars ([C/Fe] $> +1.0$; see Beers \\& Christlieb 2005).  Recent high-resolution spectroscopic studies of CEMP stars selected from the two large objective-prism surveys for low-metallicity stars, the HK \\citep{BPS1985,BPS1992,beers99} and Hamburg-ESO surveys \\citep{HES}, have revealed a wide variety of elemental abundance patterns (e.g, \\citealt{norris1997b,norris1997c, bonifacio1998,aoki2002a, aoki2002b, aoki2002c}; Barklem et al. 2005; Beers \\& Christlieb 2005). Many of these stars (roughly 70\\% - 80\\%; Aoki et al. 2003, 2006a) exhibit strong s-process enhancements (CEMP-s), and are thought to be the result of AGB mass transfer to the presently observed companion (e.g., \\citealt{aoki2002c, lucatello2003, sivarani2004, barbuy2005}).  The observed radial-velocity variations among the CEMP-s stars indicate that most, if not all, could be members of binary systems (\\citealt{tsang2004, lucatello2005}). Even within this group of CEMP-s stars, the observed s-process patterns are not identical, however. For example, the scatter in their observed Sr, Ba, Pb, and C abundances at a given [Fe/H] is large (e.g., \\citealt{aoki2002c, sivarani2004, barbuy2005}). This diversity presumably arises because of the range in masses of the star that underwent AGB evolution (see Herwig 2005 for a recent review), as well as due to details of internal mixing and possible variations in the level of dilution of the material passed to the star we now observe.   \\begin{figure}[th] \\rotatebox{0}{\\resizebox{8.5cm}{!}{\\includegraphics{5440fig2.eps}}} \\caption{Observed and synthetic spectra of our stars around the 388.3 nm CN band (symbols as in Fig. 1). The strong CN bands in \\cstwo\\ indicate a larger N excess than in \\csthree. In \\cstwonine\\ the CN bands are weak, but the derived [N/H] is high because [C/H] is low and \\teff\\ is high. The 336 nm NH bands are also detected in \\cstwonine\\ and \\cstwo\\ (not shown). } %\\label{} \\end{figure}  \\begin{table*}[ht] \\caption{Log of observations} \\begin{tabular}{lclllll}  \\hline Star         & Coordinates     &       Date     &  MJD-24000.5   & Exposure  & Radial Velocity  & S/N\\\\ &(2000)&y/m/d&&(sec)&\\kms&at 430 nm\\\\ \\hline  CS~31080-095   &  04:44:21.8 $-$45:13:57 &   2000-10-19   &  51836.2572450 & 2400& $-$43.64 &93 \\\\ &                         &   2000-10-19   &  51836.2864673 & 2400& $-$43.30 &72 \\\\ &                         &   2001-10-19   &  52201.3331806 & 3600& $-$43.76 &82 \\\\ &                         &   2001-10-20   &  52202.3388291 & 3000& $-$43.76 &52 \\\\ &                         &   2001-11-06   &  52219.2454194 & 3600& $-$43.66 &92 \\\\ &                         &   2001-11-09   &  52222.2693614 & 1000& $-$43.34 &48 \\\\ &&&&&\\\\ CS~22958-042   &  02:01:07.5 $-$57:17:07 &   2001-11-05   &  52218.1900217 & 3600& $+$165.21& 48 \\\\ &                         &   2001-11-05   &  52218.2360571 & 3600& $+$162.51& 35 \\\\ &&&&&\\\\ CS~29528-041   &  02:29:25.1 $-$18:13:30 &   2001-09-06   &  52158.3486878 & 5000& $-$285.1 & 50\\\\ &                         &   2001-10-21   &  52203.2220389 & 3260& $-$285.4 & 33 \\\\ &                         &   2001-11-07   &  52220.2338976 & 3600& $-$285.0 & 37 \\\\ \\hline \\end{tabular} \\end{table*}   \\begin{figure}[th] \\rotatebox{0}{\\resizebox{8.5cm}{!}{\\includegraphics{5440fig3.eps}}} \\caption{Observed spectra (crosses) for the two stars with meaningful measurements or limits on the $^{12}$C/$^{13}$C ratio. Solid, dotted, dashed, and dash-dotted lines show our synthetic spectra for $^{12}$C/$^{13}$C = 99, 40, 9, and 3, respectively. For \\cstwonine\\, all $^{12}$C/$^{13}$C ratios between 3 and 99 are within a 2 $\\sigma$ detection limit; hence, no $^{12}$C/$^{13}$C is given for this star.} \\end{figure}   %\\scriptsize \\begin{table}[h] \\caption{Photometry and adopted atmospheric parameters} \\begin{tabular}{lrrr} \\hline Parameter          & CS~31080-095   & CS~22958-042  & CS~29528-041\\\\ \\hline $V$            &  12.989       & 14.516       & 14.60\\\\ $B-V$          &  0.521        & 0.479        & 0.42\\\\ $U-B$          & $-$0.291      &  \\nodata     & $-$0.19\\\\ $V-R$          &  0.317        & 0.294        & \\nodata \\\\ $V-I$          &  0.617        & 0.614        & \\nodata \\\\ %J             &  11.942       & 13.515       & \\\\ %H             &  11.670       & 13.288       & \\\\ %K             &  11.592       & 13.213       & \\\\ $V-K$          &  1.397        & 1.303        & 1.292\\\\ $J-H$          &  0.272        & 0.227        & 0.314 \\\\ $J-K$          &  0.350        & 0.302        & 0.339 \\\\ &               &              & \\\\ E(B$-$V)       &  0.009        & 0.025        & 0.030 \\\\ &               &              & \\\\ \\teff($B-V$)$_{0}$   &  5671         & 5897         & 6378 \\\\ \\teff($V-R$)$_{0}$   &  5964         & 6221         &\\nodata \\\\ \\teff($V-I$)$_{0}$   &  6020         & 6146         &\\nodata \\\\ \\teff($R-I$)$_{0}$   &  6224         & 6135         &\\nodata \\\\ \\teff($V-K$)$_{0}$   &  6077         & 6344         & 6199 \\\\ \\teff($J-H$)$_{0}$   &  5940         & 6277         & 5947 \\\\ \\teff($J-K$)$_{0}$   &  5822         & 6156         & 6266 \\\\ &               &              & \\\\ Adopted: & & & \\\\ \\teff(Fe lines)      & 6050          & 6250         & 6150 \\\\ \\logg                & 4.50          & 3.50         & 4.00 \\\\ \\feh                 & $-$2.85       &$-$2.85       & $-$3.30 \\\\ \\mic(\\kms)           & 1.0    & 1.5          & 1.3 \\\\ \\hline \\end{tabular} \\end{table} %\\normalsize   Other CEMP stars exhibit r-process enhancements (e.g., CS~22892-052; Sneden et al. 2003); Lucatello (private communication) reports several additional likely CEMP-r stars from an analysis of the HERES sample of Barklem et al. (2005). Another class of carbon-enhanced metal-poor stars are the CEMP-r/s stars, which exhibit the presence of neutron-capture elements associated with {\\it both} the r- and s-process (e.g., CS~22948-027, CS~29497-034, Hill et al. 2000 and Barbuy et al. 2005; CS 29526-110, CS~22898-027, CS~31062-012, Aoki et al. 2002c; HE~2148-1247, Cohen et al. 2003; CS~29497-030, Ivans et al. 2005).  Most CEMP-r/s stars also exhibit very high Pb abundances (the so-called ``lead stars,\" see Aoki et al. 2002c and van Eck et al. 2003). The discovery of CS~22957-027 \\citep{norris1997c,bonifacio1998} and a number of similar stars which exhibit {\\it no} enhancements of neutron-capture elements \\citep{aoki2002a,aoki2002b,depagne,2005astro.ph..1535P,2005Natur.434..871F}, led Beers \\& Christlieb (2005) to introduce the CEMP-no class of stars.  As shown below, at least one of the three stars studied here is classified as a CEMP-no star. The other two stars in our study exhibit intermediate Ba excesses ([Ba/Fe] $>$ 0), but no other strong s-process-element enrichment. Thus, they fail the Beers \\& Christlieb (2005) criterion for the Ba ratios in CEMP-no stars (which requires [Ba/Fe] $<$ 0.0), but they may well be related to this class.  The origin of CEMP-no stars remains unclear at present. They could, for example, be the result of pollution by AGB stars which, for some reason, have been unable to activate the s-process, but have still dredged up significant amounts of carbon to their outer atmospheres. Alternatively, the high CNO abundances in these stars could be the products of SNe; such models have been proposed at least in the case of HE~0107-5240 \\citep{2002Natur.419..904C,2003Natur.422..834B} and HE~1327-2326 \\citep{2005Natur.434..871F,2005Sci...309..451I}.  Measurements of C, N, and O abundances provide crucial information on the nucleosynthetic and mixing history of the progenitors of CEMP stars. The $^{12}$C/$^{13}$C ratio is a sensitive indicator of the mixing processes experienced by carbon-enhanced stars; Li provides another valuable diagnostic tool. Lithium is not detected in most CEMP stars, but is sometimes found at a level below the Spite Plateau (\\citealt{norris1997a, aoki2002b}); both observations suggest that some mechanism induced strong mixing and Li burning in the past. However, at least two known CEMP stars, LP~706-7 \\citep{norris1997b} and CS~22898-027 \\citep{TB1992} exhibit Li abundances corresponding to the Spite Plateau. Perhaps surprisingly, both of these stars are neutron-capture-element-rich stars (\\citealt{norris1997b, aoki2002c}), although significant radial-velocity variation has yet to be found in either star.  In this paper we present a detailed abundance analysis for three main-sequence turnoff CEMP stars, \\csthree, \\cstwo, and \\cstwonine, based on high-resolution, high S/N observations obtained with the ESO VLT/UVES. Main-sequence turnoff CEMP stars are of particular interest because their present evolutionary stage precludes the significant internal mixing that is expected for cooler CEMP stars, hence their atmospheres better preserve the nucleosynthetic signatures of their progenitors.  We describe our observations in Sect. 2. Details of our abundance analysis procedures are reported in Sect. 3; in Sect. 4 we present the results, and in Sect. 5 we examine possible scenarios for the nucleosynthesis histories of these objects and compare them with the properties of previously observed CEMP stars. Our conclusions are presented in Sect. 6.  Throughout this paper, the abbreviations VMP, EMP, UMP, HMP, and MMP (Very, Extremely, Ultra, Hyper, and Mega Metal Poor) denote stars with [Fe/H] below $-$2.0, $-$3.0, $-$4.0, $-$5.0, and $\u0096$6.0, respectively (Beers \\& Christlieb 2005). ",
        "conclusions": "The origin of the observed abundance patterns of CEMP stars remains enigmatic. Fig. 6 summarises how the abundances in our stars compare to the predicted abundances in AGB ejecta of very low metallicity, presumably transferred to the observed stars from former AGB binary companions. It should be recognized that the stars have lower [Fe/H] than the models. Because all the overabundances shown are dominated by the primary production in AGB stars, this could be corrected to first order by just adding the Fe abundance difference (in logarithmic units). However, a stringent comparison of models and actual stellar observations must include the problems of mass transfer and dilution. For that reason, one should not pay too much attention to absolute overabundances of 0.5 -- 1 dex, but rather focus on the abundance patterns.  In our small sample, \\cstwonine\\ is probably the most straightforward case; all available evidence points toward pollution from a HBB AGB star. \\cstwonine\\ is one of the very few known NEMP stars, the existence of which is predicted by standard stellar evolution in conjunction with the binary mass-transfer scenario for CEMP-s stars. The moderate Ba overabundance in this star may be due to a polluting AGB star of large mass. The CNO abundances would also be compatible with pollution from a massive, rotating star. The s-process in these (or indeed any VMP AGB) stars is not yet understood. However, the observed Li, Na, and Ba abundances clearly favour the AGB alternative.  Assuming that the measured rapid velocity variation for \\cstwo\\ is due to a binary orbital motion of  short period, the AGB connection strongly suggests that this star is a member of a post-common-envelope binary. This {\\it must} be the case if the overabundances are attributed to mass-transfer from an AGB star.  We can only speculate about the effects of common-envelope phase on the nuclear signatures in a CEMP star that formed from this mechanism. One scenario could involve several thermal pulses with efficient dredge-up causing the observed overabundances of C, N, O and Na. However, before the s-process production can get fully underway, AGB evolution is terminated by the onset of common-envelope evolution. This could explain the absence of Ba in this star.  Another CEMP star with a short orbital period is HE~0024-2523 (Lucatello et al. 2003), which must be a post-common envelope object as well. However, in that star Ba (and other heavy elements) are overabundant without doubt. These two short-period binaries have contrasting Na signatures: Na is underabundant in HE~0024-2523, but \\cstwo\\ is one of the most Na-rich stars yet found.  Both HE~0024-2523 and \\cstwo\\ have very large C/O ratios, although O is significantly overabundant: C/O$ = 48$ in \\cstwo, and 100 in HE~0024-2523. A very large C/O ratio may present a problem for the AGB mass-transfer scenario, because a non-HBB AGB star can only deliver material with a lower C/O ratio than found anywhere inside of it. The material dredged up to the surface, usually from the intershell, contains from one to a few percent primary O by mass, almost independent of metallicity, and also maybe $20 - 25\\%$ carbon by mass\\footnote{These C and O abundances are for models without extra convective mixing. If such mixing is included both C and O increase, but the C/O ratio does not change much.}. Thus, C/O $<32$ is a reasonable limit for the intershell layers of AGB stars (including VMP AGB stars). Even if the AGB star dumps this maximally polluted material onto the CEMP star we now observe, the resulting C/O ratio cannot become larger than in the intershell layer. The C/O ratio could be a problem for the two known short-period binaries, but perhaps especially for HE~0024-2523.  Clearly, common-envelope evolution introduces a whole new family of possible evolutionary outcomes. Nevertheless, it is noteworthy that we may have found the second short-period CEMP binary, which seems to be significantly different from the previously known case, HE~0024-2523. One of the two might in fact not be a post-common envelope system, but have an inert low-mass MS companion instead. However, VMP AGB stars remain the best present candidate nuclear production sites for explaining both the s-process enhancement in HE~0024-2523 and the high Na in \\cstwo.  Finally, \\csthree\\ could be an AGB mass-transfer object if VMP non-HBB AGB stars experience some extra mixing. This could account for the observed Li, C, and N abundances. The O abundance in this star is certainly high compared to the predictions of this scenario, but cannot rule it out. The problem is the low Na abundance, which excludes AGB progenitors above 2 \\msun. Perhaps a less massive AGB star is involved, in which the $^{22}$Ne source is not activated at all; the Ba could then still come from the $^{13}$C neutron source. The larger $^{12}$C/$^{13}$C ratio ($> 40$) in this star is consistent with this picture, and inconsistent with the obvious alternatives; the $^{12}$C/$^{13}$C ratios in HBB AGB stars ($^{12}$C/$^{13}$C $\\sim 10$) as well as the wind ejecta of rotating massive stars ( $^{12}$C/$^{13}$C $\\sim 4$) are much lower than the observed limit.  This discussion has clearly shown that our current understanding of the origin of the abundance patterns in CEMP stars is still in its infancy. The large overabundances of some elements in CEMP stars probe one or more important primary nuclear production sites, including those that may have set the stage for the following galactic chemical evolution. Yet, we remain far from having identified even the main players in this game. Tasks to be addressed include a better understanding of mixing in stars of all masses, and of the impact of binary evolution on nucleosynthesis.  A significant increase in the numbers of CEMP stars with accurate, detailed elemental abundances is also an obvious next step."
    },
    "0608/astro-ph0608169_arXiv.txt": {
        "abstract": "We report the detection of ongoing star formation in the prominent tidal arms near NGC\\,3077 (member of the M\\,81 triplet).  In total, 36 faint compact \\ion{H}{2} regions were identified, covering an area of $\\sim4\\times6$ kpc$^2$. Most of the \\ion{H}{2} regions are found at \\ion{H}{1} column densities above $1\\times10^{21}$\\,cm$^{-2}$ (on scales of 200\\,pc), well within the range of threshold columns measured in normal galaxies. The \\ion{H}{2} luminosity function resembles the ones derived for other low--mass dwarf galaxies in the same group; we derive a total star formation rate of 2.6$\\times10^{-3}$ \\Msun\\,yr$^{-1}$ in the tidal feature. We also present new high--resolution imaging of the molecular gas distribution in the tidal arm using CO observations obtained with the OVRO interferometer. We recover about one sixth of the CO flux (or M$_{H2}\\sim2\\times10^6$\\,\\Msun, assuming a Galactic conversion factor) originally detected in the IRAM 30\\,m single dish observations, indicating the presence of a diffuse molecular gas component in the tidal arm. The brightest CO peak in the interferometer map (comprising half of the detected CO flux) is coincident with one of the brightest \\ion{H}{2} regions in the feature. Assuming a constant star formation rate since the creation of the tidal feature (presumably $\\sim3\\times10^8$\\,years ago), a total mass of $\\sim7\\times10^5$ \\Msun\\ has been transformed from gas into stars. Over this period, the star formation in the tidal arm has resulted in an additional enrichment of $\\Delta$Z$>$0.002. The reservoir of atomic and molecular gas in the tidal arm is $\\sim3\\times10^8$ \\Msun, allowing star formation to continue at its present rate for a Hubble time.  Such wide--spread, low--level star formation would be difficult to image around more distant galaxies but may be detectable through intervening absorption in quasar spectra. ",
        "introduction": "Studies of the processes that lead to star formation are critical for our understanding of how galaxies form and evolve with time. In massive galaxies, star formation is triggered through a variety of processes: e.g., compression of molecular cloud cores in spiral density waves, tidal interactions, and bar formation. On scales of a few hundred parsecs, star formation is observed where the gas density exceeds a threshold, whose value varies with the critical density for local gravitational stability (e.g., Martin \\& Kennicutt 2001). Models that rely on thermal instabilities to determine the star formation threshold predict a constant threshold column density (Schaye 2004).  Studies of star formation (SF) in tidal arms of interacting galaxies offer the possibility to study SF processes at conditions that are distinctly different from those in normal galactic disks (e.g., given the absence of spiral density waves and bars).  Nearby, interacting systems give us the unique opportunity to study the properties of tidal features (e.g.  atomic and molecular gas content and \\ion{H}{2} region properties) in detail.  Tidal dwarf candidates have been detected in tidal tails of interacting galaxies in a number of cases (e.g., see list in Braine et al.\\ 2001); but selection effects limit these to bright, and therefore fairly evolved, tidal dwarfs. The detection of SF regions around interacting systems (`intergalactic' \\ion{H}{2} regions) has been reported in Ryan--Weber et al.\\ (2004). Given the much higher incidence of interacting galaxies in the past, a significant fraction of today's dwarf galaxies may have formed out of tidal debris (Hunsberger \\etal 1996, Hunter \\etal 2000, Okazaki \\& Taniguchi 2000, Duc et al.\\ 2004).  Also of interest in this context is the self-enrichment of these tidal debris, and their possible incidence as intervening metal line systems.  The nearby M\\,81 triplet (consisting of M\\,82, M\\,81 and NGC\\,3077) is an ideal testbed to study tidal arms around interacting galaxies in detail.  The \\ion{H}{1} map of Yun et al. (1994) shows a geometrical cross section of roughly 50 kpc $\\times$ 100~kpc at column densities of N(HI)$>$1.8$\\times$10$^{20}$ cm$^{-2}$.  The impressive tidal arm near NGC\\,3077 is the subject of the present study (total \\ion{H}{1} mass: $\\sim5\\times10^8$\\,\\Msun, Walter et al.\\ 2002).  According to various numerical simulations (Thomasson \\& Donner 1993, Brouillet \\etal 1991, Yun \\etal 1993) this tail has been created by a recent ($\\sim3\\times10^8$\\,yr) interaction with M\\,81.  Barbieri et al. (1974) detected `a fragmentary complex of almost stellar objects' in this region (the `Garland', see also Karachentsev \\etal 1985, Sharina 1991, Sakai \\& Madore 2001). The presence of extended molecular gas in this tidal feature has been reported by Walter \\& Heithausen (1999) and Heithausen \\& Walter (2000). In this paper we report on the discovery of low--level star formation over roughly 24~kpc$^2$ and present new high--resolution observations of the molecular gas distribution in this feature.  In Section 2 we summarize our observations obtained at the 2\\,m telescope at Kitt Peak, and at the OVRO millimeter interferometer. In Sec.~3 we present the distribution of the \\ion{H}{2} regions, the \\ion{H}{2} luminosity function and the high--resolution CO maps; Sec.~4 summarizes our results. ",
        "conclusions": "We present the discovery of wide-spread, low-level star formation in the tidal feature near NGC\\,3077 (total of 36 \\ion{H}{2} regions). The \\ion{H}{1} column densities exceed $1\\times10^{21}$\\,cm$^{-2}$ over the $\\sim$24\\,kpc$^2$ region where the \\ion{H}{2}\\ regions are found.  The total star formation rate in the tidal arm is SFR=$2.3\\times10^{-3}$\\,M$_\\odot$ yr$^{-1}$.  These properties place the complexes amongst the faintest star forming regions detected in tidal arms around galaxies so far (cf. Ryan-Weber et al.\\ 2004).  The presence of a huge reservoir of molecular gas ($\\sim 10^7$\\,\\Msun, Heithausen \\& Walter 2000, using the Galactic conversion factor) as well as atomic gas (M$_{\\rm HI}\\sim3\\times10^8$\\,\\Msun\\ in the region shown in Fig.~2) suggests that star formation may continue at the current rate for a Hubble time. If some 10 percent of the gas were transformed into stars, the total luminosity of the tidal dwarf will eventually be of order $5\\times10^{7}$\\,M$_{\\rm \\odot}$ (here we simplisticly adopt a stellar mass--to--light ration of 1\\,M$_{\\rm \\odot}$/L$_{\\rm \\odot}$), similar to other faint dwarf galaxies in the M\\,81 group.  It is interesting to note that no clear correlation exists between the sites of on--going star formation and gas column densities (both molecular and atomic): \\ion{H}{2} regions are not always found where the \\ion{H}{1} or molecular column density is highest (e.g., Fig.~3). However, star formation is only seen where the \\ion{H}{1} column density reaches values larger than $\\sim 1.0\\times10^{21}$\\,cm$^{-2}$ (Tab.~1) -- this value is consistent with what is typically found in (dwarf) galaxies (e.g., Skillman 1987, Walter \\& Brinks 1999, Martin \\& Kennicutt 2001).  Also, the derived \\ion{H}{2} luminosity function is, within the errors, similar to other low-mass dwarf galaxies in the same group (such as Holmberg~I). No \\ion{H}{2} regions are found with log(L[erg\\,s$^{-1}$])$>$37.4 -- this is more than an order of magnitude fainter than what is found in more massive dwarf galaxies in the same group (e.g. IC\\,2574) and bigger spiral galaxies (e.g., in M51 luminosities reach values of log(L)[erg\\,s$^{-1}$]=39, e.g., Rand 1993). None of the \\ion{H}{2} regions are coincident with the X--ray sources found in the same region of the sky (Ott, Martin \\& Walter 2003).  It is difficult to estimate when star formation started in the tidal arm. From the fact that we see H$\\alpha$ emission today it is clear that star formation is occuring well after the creation of the structure some $\\sim3\\times 10^8$\\,yr ago (i.e., `in situ' formation). If we assume a constant star formation rate over the last $3\\times10^8$\\,yr we estimate a total mass of newly formed stars of $\\sim7\\times10^{5}$\\,M$_{\\odot}$.  This star formation is enriching the ISM of the tidal tail with heavy elements: Taking our derived assembled stellar mass since the creation of the feature, and a yield of 2\\%, we derive a metal production of 1.4$\\times$10$^4$\\,\\Msun\\ over the last $3\\times10^8$\\,yr. We divide this number by the total amount of gas within the $\\sim$4$\\times$6\\,kpc region enclosed by the 10$^{21}$\\,cm$^{-2}$ contour (3$\\times10^8$\\,\\Msun) to get a metallicity of 4.7$\\times10^{-5}$, or roughly Z$\\sim$0.002 Solar, well below the values measured for the gas phase in the most metal--poor dwarfs.  We note that the total metallicity in the tidal feature is likely higher as models for the tidal arm show that it was stripped from the outskirts of NGC\\,3077 (Thomasson \\& Donner 1993, Brouillet \\etal 1991, Yun \\etal 1993).  Measurements of the gas phase metallicity in the tidal feature would thus clearly help to shed light on its origin and history.  We conclude that wide--spread low--level star formation may be a common phenomenon in tidal \\ion{H}{1} tails with column densities that exceed $1\\times10^{21}$\\,cm$^{-2}$, leading to chemical enrichment in these tails.  This type of low-level star formation in tidal debris should be detectable, if it is common, as intervening absorption in quasar spectra.  The tidal system discussed here would be classified as a damped Lyman--alpha absorber (DLA) with a cross section of roughly 30 kpc$^2$. A tidal origin for (at least some) DLAs would naturally explain their lack of enhanced [$\\alpha/$Fe] abundance ratios, which suggest that star formation proceeded relatively slowly (Pettini et al.\\ 1999).  Direct observations of DLAs at low redshift are consistent in that they indicate associations with a wide variety of galaxy types and relatively small impact parameters (of order 10\\,kpc, e.g., Kanekar \\& Chengalur 2003).  For comparison, the tidal system discussed here has a mean projected distance of $\\sim$5\\,kpc from NGC 3077 ($\\sim$50\\,kpc distance from M\\,81). This interesting object should be further explored via metallicity measurements and simulations that would establish the fate of the metals formed in this system."
    },
    "0608/astro-ph0608443_arXiv.txt": {
        "abstract": "{We propose a useful method for mapping large-scale velocity fields in the solar photosphere. It is based on the local correlation tracking algorithm when tracing supergranules in full-disc dopplergrams. The method was developed using synthetic data. The data processing the data are transformed during the data processing into a suitable coordinate system, the noise is removed, and finally the velocity field is calculated. Resulting velocities are compared with the model velocities and the calibration is done. From our results it becomes clear that this method could be applied to full-disc dopplergrams acquired by the Michelson Doppler Imager (MDI) onboard the Solar and Heliospheric Observatory (SoHO). ",
        "introduction": "The solar photosphere is a very dynamic layer of the solar atmosphere. It is strongly influenced by the underlying convection zone. Despite years of intensive studies, the velocity fields in the solar photosphere remain not very well known. The evidence of the vigorous and sometimes chaotic character of the motions of observed structures in the photosphere (sunspots, granules, and other features) came already from the first systematic studies made in 19th century, of which let us at least mention the discovery of the solar differential rotation (Carrington \\cite{carrington}). Motions in the photosphere are strongly coupled with magnetic fields. The large-scale velocity fields are very important for studies of the global solar dynamo.  An attempt to describe the differential rotation by a parabolic dependence did not make for clear results. Coefficients of the parabola differed according to traced objects and also changed with time, when tracing one type of object (reviewed e.~g. by Schr\\\"oter \\cite{schroeter}). The character of the differential rotation has never been in doubt.  It follows from these arguments that a temporally variable streaming of the plasma exists on the surface of the Sun, which can be roughly described by the differential rotation. This streaming has a large-scale character, large-scale plasma motions were studied for example on the basis of tracking the magnetic structures (e. g. Ambro\\v{z} \\cite{ambroz_a}, \\cite{ambroz_b}). The long-term Doppler measurements done by the MDI onboard SoHO make it possible to extend the studies of large-scale velocities in the solar photosphere. The knowledge of the behaviour of velocities in various periods of the solar activity cycle could contribute to understanding the coupling between the velocity and magnetic fields and of the solar dynamo function.  There are at least three methods calculating photospheric velocities:  \\begin{figure*} \\resizebox{0.5\\textwidth}{!}{\\includegraphics{4124fig1.eps}} \\resizebox{0.5\\textwidth}{!}{\\includegraphics{4124fig2.eps}} \\caption{Comparison of the real dopplergram (left, observed by MDI/SoHO) and the simulated one (right). Both images are visually very similar. The black colour means line-of-sight velocity $-700$~m\\,s$^{-1}$ (towards an observer) while white represents $+700$~m\\,s$^{-1}$.} \\label{svanda_fig:dopplergrams} \\end{figure*}  \\begin{enumerate} \\item \\emph{Direct Doppler measurement} -- provides only one component (line-of-sight) of the velocity vector. These velocities are generated by local photospheric structures, amplitudes of which are significantly greater than amplitudes of the large-scale velocities. The complex topology of such structures complicates an utilisation for our purpose. Analysing this component in different parts of the solar disc led to very important discoveries (e. g. supergranulation -- Hart \\cite{hart} and Leighton et al. \\cite{leighton62}).  \\item \\emph{Tracer-type measurement} -- provides two components of the velocity vector. When tracing some photospheric tracers, we can compute the local horizontal velocity vectors in the solar photosphere. Tracking motions of sunspots across the solar disc led to the discovery of the differential rotation (Carrington \\cite{carrington}).  \\item \\emph{Local helioseismology} -- provides a full velocity vector. Although local helioseismology (e. g. Zhao \\cite{zhao}) is a very promising method, it is still in progress and until now does not provide enough reliable results. \\end{enumerate}  Since the photosphere is a very thin layer (0.04~\\% of the solar radius), the large-scale photospheric velocity fields have to be almost horizontal. Then, the tracer-type measurement should be sufficient for mapping the behaviour of such velocities. In this field the local correlation tracking (LCT) method is very useful.  This method was originally designed for the removal of the seeing-induced distortions in image sequences (November \\cite{november86}) and later used for mapping the motions of granules in the series of white-light images (November \\& Simon \\cite{november88}). The method works on the principle of the best match of two frames that record the tracked structures at two different instants. For each pixel in the first frame, a small correlation window is chosen and is compared with a somewhat displaced window of the same size in the second frame. A vector of displacement is then defined as a difference in the coordinates of the centres of both windows when the best match is found. The velocity vector is calculated from this displacement and the time lag between two frames.  LCT was recently used for tracking many features in various types of observations, especially for tracking the granules in high-resolution white-light images (e. g. Sobotka et al. \\cite{sobotka99}, \\cite{sobotka00}).  The method needs a tracer -- a significant structure recorded in both frames, the lifetime of which is much longer than the time lag between the correlated frames. We decided to use the supergranulation pattern in the full-disc dopplergrams, acquired by the MDI onboard SoHO. We assume that supergranules are carried as objects by the large-scale velocity field. This velocity field is probably located beneath the photosphere, so that the resulting velocities will describe the dynamics in both the photospheric and subphotospheric layer. The existence of the supergranulation on almost the whole solar disc (in contrast to magnetic structures) and its large temporal stability make the supergranulation an excellent tracer.  The resulting velocities cover the whole solar disc. Hence the velocity field also describes a motion of the supergranulation in the areas occupied by active regions or by magnetic field concentrations. This fact allows the results to be used to study the mutual motions of substructures like sunspots, magnetic field in active regions, background fields, and the quiet photosphere.  Supergranules are structures with a strong convection coupling. The mean size of the supergranular cell is approx. 30~Mm (e. g. Wang \\& Zirin \\cite{wang}), with the size and shape dependent on the phase of the solar cycle (e. g. S\\' ykora \\cite{sykora}). Supergranules are quite stable with a mean lifetime of approx. 20~hours (e. g. Leighton \\cite{leighton64}). Information about the distribution function of the lifetime was published by  DeRosa et al. (\\cite{derosa}). The internal velocity field in the supergranular cell is predominantly horizontal (e. g. Hathaway et al. \\cite{hathaway}) with the amplitude approx. 300~m\\,s$^{-1}$. Due to the horizontality of the internal velocity field, the supergranules can be observed in dopplergrams on the whole solar disc except for its centre.  We do not propose that this method capable of measuring velocities of order 1~m\\,s$^{-1}$, but we do expect that the large-scale velocities will have magnitudes at least one order greater. We also have to take the largest-scale velocities into account like the differential rotation or meridional circulation, which have velocities of at least 10~m\\,s$^{-1}$. If for example we take differential rotation described by the formula $\\omega {\\rm[deg/day]}= 13.064-1.69 \\sin^2 b-2.35 \\sin^4 b$ (Snodgrass \\cite{snodgrass}), we have to expect a velocity approx. 190~m\\,s$^{-1}$ in $b=60\\,^\\circ$ in the Carrington's coordinate system. The main goal of our study (and the proposed method of mapping the velocities should be proxy for this purpose) is to separate the superposed velocity field into the components and to investigate their physics.  This paper is the first one of a series about the large-scale photospheric velocity fields. In the following papers we shall apply the method here suggested to observed data and describe the properties of the large-scale velocity fields in different periods of the solar activity cycle. ",
        "conclusions": "We have developed a method for mapping the velocity fields in the solar photosphere based on the local correlation tracking (LCT) algorithm. The method consists in two main steps. In the first main step the mean zonal velocities are calculated and, on the basis of expansion to the Fay's formula (\\ref{svanda_eq:fay}), the differential rotation is removed. In the second main step, the LCT algorithm with an enhanced sensitivity is applied. Finally, the differential rotation (obtained in the first step) is added to the vector velocity field obtained in the second main step. Both main steps can be divided into a few substeps, which are mostly common. \\begin{enumerate} \\item ``Derotation'' of the data series using the Carrington's rotation rate in the first step and using the calculated differential rotation in the second step.  \\item Transformation of the data series into the Sanson-Flamsteed coordinate system to remove the geometrical distortion caused by the projection to the disc.  \\item The $k$-$\\omega$ filtering (with the cut-off velocity 1\\,500~m\\,s$^{-1}$) for suppression of the noise coming from the evolutionary changes of supergranules, of the numerical noise, and for the partial removal of the ``blind spot'' (an effect at the centre of the disc).  \\item Application of the LCT: the lag between correlated frames is 4~hours, the correlation window with FWHM 60\\arcsec, the measure of correlation is the sum of absolute differences and the nine-point method for calculation of the subpixel value of displacement is used. The calculated velocity field is averaged over the period of one day. \\end{enumerate}  \\noindent Finally, the magnitude of the calculated vector field is corrected using formula (\\ref{svanda_eq:calibration}) obtained from the tests on the synthetic data.  According to our test, the method provides a very reliable tool for mapping the velocity fields with the spatial resolution of 60\\arcsec{} (43,5~Mm) on the solar disc and with the accuracy of 15~m\\,s$^{-1}$ for velocity magnitudes under 100~m\\,s$^{-1}$ and 25~m\\,s$^{-1}$ for velocity magnitudes greater than 100~m\\,s$^{-1}$. The method is ready to use on the real full-disc dopplergrams acquired by the MDI onboard SoHO. Our method shows that the usability of the LCT method is wider than originally expected."
    },
    "0608/astro-ph0608675_arXiv.txt": {
        "abstract": "% Gravitational lensing in a modified gravity (MOG) is derived and shown to describe lensing without postulating dark matter. The recent data for merging clusters identified with the interacting cluster 1E0657-56 is shown to be consistent with a weak lensing construction based on MOG without exotic dark matter. ",
        "introduction": "A relativistic modified gravity (MOG) theory~\\cite{Moffat,Moffat2} has been proposed to explain the rotational velocity curves of galaxies and the X-ray data for clusters of galaxies with a modified Newtonian acceleration law, without non-baryonic dark matter. A fitting routine for galaxy rotation curves has been used to fit a large number of galaxy rotational velocity curve data, including low surface brightness (LSB), high surface brightness (HSB), dwarf galaxies and elliptical galaxies with both photometric data and a two-parameter core model without non-baryonic dark matter~\\cite{Brownstein}. The fits to the data are remarkably good and for the photometric data only the one parameter, the mass-to-light ratio $\\langle M/L\\rangle$, is used for the fitting, once two parameters are universally fixed for galaxies and dwarf galaxies. A large sample of mass profile X-ray cluster data has also been fitted~\\cite{Brownstein2} without dark matter. It has been shown that MOG can fit the Cosmic Microwave Background acoustical oscillation peaks data in the power spectrum without dark matter and provide an explanation for the accelerated expansion of the universe~\\cite{Moffat2}.  The MOG requires that Newton's constant G, the coupling constant $\\omega$ that measures the strength of the coupling of a skew field to matter, and the mass $\\mu$ of the skew field vary with distance and time, so that agreement with the solar system and the binary pulsar PSR 1913+16 data can be achieved, as well as fits to galaxy rotation curve data and galaxy cluster data. In MOG ~\\cite{Moffat,Moffat2}, the action contains the Einstein-Hilbert action based on a symmetric pseudo-Riemannian metric, an action formed from a vector field $\\phi_\\mu$ called the phion field which produces a ``fifth'' force skew field, and an action for scalar fields that leads to {\\it effective} field equations describing the variations of $G$, $\\omega$ and $\\mu$.  In the following, we shall investigate the gravitational lensing in MOG. The variation of $G$ leads to a consistent description of relativistic lensing effects for galaxies without non-baryonic dark matter. We study the weak gravitational lensing of the merging, interactive cluster 1E0657-56 at a redshift $z=0.296$ that has recently been claimed to enable a direct detection of dark matter, without alternative gravitational theories~\\cite{Clowe,Clowe2,Clowe3}. We will show that the lensing of distant background galaxies predicted by MOG is consistent with the data from the interacting cluster 1E0657-56. ",
        "conclusions": "The gravitational potential determined from the STVG MOG~\\cite{Moffat} has a unique signature for a merging of galaxy clusters.  Other alternative models based on Milgrom's MOND~\\cite{Milgrom} and Bekenstein's and Sander's relativistic generalizations of MOND~\\cite{Bekenstein,Sanders,Sanders2} have been studied by several authors~\\cite{Mortlock,Zhao,Angus}. In a MOND-type model and in the Bekenstein and Sanders models, the MOND critical acceleration $a_0$ is expected to satisfy $a\\gg a_0\\sim 1.2\\times 10^{-8}\\,{\\rm cm}\\,{\\rm s}^{-2}$ inside the interacting cluster, while the MOND modified acceleration law comes into play outside the interacting cluster for $a\\ll a_0$. It therefore seems difficult to understand how the lack of peaking of the off-set central X-ray gas cloud compared to the more pronounced peaking of the outer galaxy mass distribution can be explained by a MOND-like model. In view of this, MOND-like models would be expected to predict an Einstein (Newtonian) gravitational field for the weak lensing of the interacting cluster, requiring dark matter to fit the data. It is already known that MOND does not fit the mass profiles of X-ray clusters without dark matter~\\cite{Sanders,Sanders2}, whereas MOG has been shown to fit a large number of mass profiles of X-ray clusters without dark matter~\\cite{Brownstein2}. With the accumulation of more data for 1E0657-56, this interacting cluster could distinguish MOG from other alternative gravity theories which purport to fit galaxy and clusters of galaxies data without dark matter. A more detailed study of the MOG prediction for the interacting cluster based on a fitting to published data will be considered in a future publication.  We learn from the results presented here that {\\it one should not draw premature conclusions about the existence of dark matter} without a careful analysis of alternative gravity theories and their predictions for galaxy lensing and cluster lensing, in particular, for the interacting cluster 1E0657-56.  \\vskip 0.4 true in {\\bf Acknowledgments} \\vskip 0.2 true in  This work was supported by the Natural Sciences and Engineering Research Council of Canada."
    },
    "0608/astro-ph0608505_arXiv.txt": {
        "abstract": "In the era of high precision CMB measurements, systematic effects are beginning to limit the ability to extract subtler cosmological information.  The non-circularity of the experimental beam has become progressively important as CMB experiments strive to attain higher angular resolution and sensitivity. The effect of non-circular beam on the power spectrum is important at multipoles larger than the beam-width.  For recent experiments with high angular resolution, optimal methods of power spectrum estimation are computationally prohibitive and sub-optimal approaches, such as the Pseudo-$C_l$ method, are used. We provide an analytic framework for correcting the power spectrum for the effect of beam non-circularity and non-uniform sky coverage (including incomplete/masked sky maps). The approach is perturbative in the distortion of the beam from non-circularity allowing for rapid computations when the beam is mildly non-circular. When non-circular beam effect is important, we advocate that it is computationally advantageous to employ `soft' azimuthally apodized masks whose spherical harmonic transform die down fast with $m$. ",
        "introduction": "The fluctuations in the Cosmic Microwave background (CMB) radiation are theoretically very well understood, allowing precise and unambiguous predictions for a given cosmological model~\\cite{bonLH,hu_dod02}. Consequently, measurement of CMB anisotropy has spearheaded the remarkable transition of cosmology into a precision science. The transition has also seen the emergence of data analysis of large complex data sets as an important and challenging component of research in cosmology. Increasingly sensitive, high resolution, measurements over large regions of the sky pose a stiff challenge for current analysis techniques to realize the full potential of precise determination of cosmological parameters. The analysis techniques must not only be computationally fast to contend with the huge size of the data, but, the higher sensitivity also limits the simplifying assumptions that can be then invoked to achieve the desired speed without compromising the final precision goals. There is a worldwide effort to push the boundary of this inherent compromise faced by the current CMB experiments that measure the anisotropy in the CMB temperature and its polarization.  Accurate estimation of the angular power spectrum, $C_{l}$, is arguably the foremost concern of most CMB experiments. The extensive literature on this topic has been summarized in literature~\\cite{hu_dod02,efs04}. For Gaussian, statistically isotropic CMB sky, the $C_l$ that corresponds to the covariance that maximizes the multivariate Gaussian PDF of the temperature map, $\\Delta T(\\q)$ is the Maximum Likelihood (ML) solution. Different ML estimators have been proposed and implemented on CMB data of small and modest sizes~\\cite{gor94,gor_hin94,gor96,gor97,max97,bjk98}. While it is desirable to use optimal estimators of $C_l$ that obtain (or iterate toward) the ML solution for the given data, these methods are usually limited by the computational expense of matrix inversion that scales as $N_d^3$ with data size $N_d$~\\cite{bor99,bon99}.  Various strategies for speeding up ML estimation have been proposed, such as, exploiting the symmetries of the scan strategy~\\cite{oh_sper99}, using hierarchical decomposition~\\cite{dor_knox01}, iterative multi-grid method~\\cite{pen03}, etc. Variants employing linear combinations of $\\Delta T(\\q)$ such as $a_{lm}$ on set of rings in the sky can alleviate the computational demands in special cases~\\cite{harm02,wan_han03}.  Other promising `exact' power estimation methods have been recently proposed~\\cite{knox01,wan04,jew04}.  However there also exist computationally rapid, sub-optimal estimators of $C_l$. Exploiting the fast spherical harmonic transform ($\\sim N_d^{3/2}$), it is possible to estimate the angular power spectrum $C_l= \\sum_m |a_{lm}|^2/(2l+1)$ rapidly~\\cite{yu_peeb69,peeb73}. This is commonly referred to as the Pseudo-$C_l$ method~\\cite{wan_hiv03}. \\footnote{Analogous approach employing fast estimation of the correlation function $C(\\q\\cdot\\qp)$ have also been explored~\\cite{szap01,szap_prun01}.}  It has been recently argued that the need for optimal estimators may have been over-emphasized since they are computationally prohibitive at large $l$ . Sub-optimal estimators are computationally tractable and tend to be nearly optimal in the relevant high $l$ regime. Moreover, already the data size of the current sensitive, high resolution, `full sky' CMB experiments such as WMAP have been compelled to use sub-optimal Pseudo-$C_l$ related methods~\\cite{Bennett:2003, hin_wmap03}.  On the other hand, optimal ML estimators can readily incorporate and account for various systematic effects, such as noise correlations, non-uniform sky coverage and beam asymmetries. The systematic correction to the Pseudo-$C_l$ power spectrum estimate arising from non-uniform sky coverage has been studied and implemented for CMB temperature~\\cite{master} and polarization~\\cite{Brown:2004}.  The systematic correction for non circular beam has been studied by us~\\cite{mit04}. Here we extend the results to include non-uniform sky coverage.    It has been usual in CMB data analysis to assume the experimental beam response to be circularly symmetric around the pointing direction. However, any real beam response function has deviations from circular symmetry. Even the main lobes of the beam response of experiments are generically non-circular (non-axisymmetric) since detectors have to be placed off-axis on the focal plane. (Side lobes and stray light contamination add to the breakdown of this assumption). For highly sensitive experiments, the systematic errors arising from the beam non-circularity become progressively more important.  Dropping the circular beam assumption leads to  major complications at every stage of analysis pipeline. The extent to which the non-circularity affects the step of going from the time-stream data to sky map is very sensitive to the scan-strategy. The beam now has an orientation with respect to the scan path that can potentially vary along the path. This implies that the beam function is inherently time dependent and difficult to deconvolve.  Even after a sky map is made, the non-circularity of the effective beam affects the estimation of the angular power spectrum, $C_l$, by coupling the power at different multipoles, typically, on scales beyond the inverse angular beam-width.  Mild deviations from circularity can be addressed by a perturbation approach~\\cite{TR:2001,fos02} and the effect of non-circularity on the estimation of CMB power spectrum can be studied (semi) analytically~\\cite{mit04}.  Fig.~\\ref{clerr} shows the predicted level of non-circular beam correction in our formalism for elliptical beams with {\\it fwhm} beam-width of $0.22^\\circ$ compared to the non-circular beam corrections computed in the recent data release by WMAP~\\cite{hin_wmap06}. \\begin{figure} \\centering \\includegraphics[angle = 0,scale = 0.45 ]{clerr_comp.eps} \\caption{The predicted non-circular beam correction for a CMB experiment with elliptical Gaussian beam with {\\it fwhm} beam-width of $0.22^\\circ$ and two values of eccentricity, $e=0.6$ and $e=0.4$ are shown as dashed-dot lines (labeled Th. Pr.). The solid curves are the non circular beam corrections estimated by the WMAP team for the Q,V and W channels. The {\\it fwhm} beam-width of $0.22^\\circ$ corresponds closely to the cleanest and highest resolution W-band beam-width. For larger beam-widths, the theoretical curves would roughly shift left in multipole by the ratio of beam-widths.} \\label{clerr} \\end{figure}   To avoid contamination of the primordial CMB signal by Galactic emission, the region adjoining the Galactic plane is masked from maps. If the Galactic cut is small enough, then the coupling matrix will be invertible, and the two-point correlation function can be determined on all angular scales from the data within the uncut sky~\\cite{mch02}.  Hivon et al.~\\cite{master} present a technique (MASTER) for fast computation of the power spectrum taking accounting for the galactic cut, but for circular beams. {\\em In our present work, we present analytical expressions for the bias matrix of the Pseudo-${C_l}$ estimator for the incomplete sky coverage, using a non-circular beam.} ",
        "conclusions": "\\label{disc}  The assumptions of non-circular beam leads to major complications at every stage of the data analysis pipeline.  The extent to which the non-circularity affects the step of going from the time-stream data to sky map is very sensitive to the scan-strategy. The beam now has an orientation with respect to the scan path that can potentially vary along the path. This implies that the beam function is inherently time dependent and difficult to deconvolve.  We extend our analytic approach for addressing the effect of non-circular experimental beam function in the estimation of the angular power spectrum $\\mathcal{C}_l$ of CMB anisotropy, which also includes the effect of the galactic cut in the entire sky map. Non-circular beam effects can be modeled into the covariance functions in approaches related to maximum likelihood estimation~\\cite{max97,bjk98} and can also be included in the Harmonic ring~\\cite{harm02} and ring-torus estimators~\\cite{wan_han03}. However, all these methods are computationally prohibitive for high resolution maps and, at present, the computationally economical approach of using a Pseudo-$C_l$ estimator appears to be a viable option for extracting the power spectrum at high multipoles~\\cite{efs04}. The Pseudo-$C_l$ estimates have to be corrected for the systematic biases.  While considerable attention has been devoted to the effects of incomplete/non-uniform sky coverage, no comprehensive or systematic approach is available for non-circular beam.  The high sensitivity, `full' (large) sky observation from space (long duration balloon) missions have alleviated the effect of incomplete sky coverage and other systematic effects such as the one we consider here have gained more significance. Non-uniform coverage, in particular, the galactic masks affect only CMB power estimation at the low multipoles. The analysis accompanying the recent second data from WMAP uses the hybrid strategy~\\cite{efs04} where the power spectrum at low multipoles is estimated using optimal Maximum Likelihood methods and Pseudo-$C_l$ are used for large multipoles~\\cite{hin_wmap06,sperg_wmap06}.   The non-circular beam is an effect that dominates at large $l$ comparable to the inverse beam width~\\cite{mit04}. For high resolution experiments, the optimal maximum likelihood methods which can account for non-circular beam functions are computationally prohibitive.  In implementing the Pseudo-$C_l$ estimation, we have included both the non-circular beam effect and the effect of non-uniform sky coverage. Our work provides a convenient approach for estimating the magnitude of these effects in terms of the leading order deviations from a circular beam and azimuthally symmetric mask. The perturbation approach is very efficient. For most CMB experiments the leading few orders capture most of the effect of beam non-circularity~\\cite{TR:2001}. Our results highlight the advantage of azimuthally smoothed masks (mild deviations from azimuthal symmetry) in reducing computational costs.  The numerical implementation of our method can readily accommodate the case when pixels are revisited by the beam with different orientations. Evaluating the realistic bias and error-covariance for a specific CMB experiment with non-circular beams would require numerical evaluation of the general expressions for $A_{l\\lp}$ using real scan strategy and account for inhomogeneous noise and sky coverage, the latter part of which has been addressed in this present work.  It is worthwhile to note in passing that that the angular power $C_l$ contains all the information of Gaussian CMB anisotropy only under the assumption of statistical isotropy.  Gaussian CMB anisotropy map measured with a non-circular beam corresponds to an underlying correlation function that violates statistical isotropy. In this case, the extra information present may be measurable using, for example, the bipolar power spectrum~\\cite{haj_sour03,haj_sour05,haj_sour06,bas06}. Even when the beam is circular the scanning pattern itself is expected to cause a breakdown of statistical isotropy of the measured CMB anisotropy ~\\cite{master}. For a non-circular beam, this effect could be much more pronounced and, perhaps, presents an interesting avenue of future study.  In addition to temperature fluctuations, the CMB photons coming from different directions have a random, linear polarization. The polarization of CMB can be decomposed into $E$ part with even parity and $B$ part with odd parity.  Besides the angular spectrum $C_l^{TT}$, the CMB polarization provides three additional spectra, $C_l^{TE}$, $C_l^{EE}$ and $C_l^{BB}$ which are invariant under parity transformations. The level of polarization of the CMB being about a tenth of the temperature fluctuation, it is only very recently that the angular power spectrum of CMB polarization field has been detected. The Degree Angular Scale Interferometer (DASI) has measured the CMB polarization spectrum over limited band of angular scales in late 2002~\\cite{kov_dasi02}. The DASI experiment recently published 3-year results of much refined measurements~\\cite{dasi_3y}. More recently, the BOOMERanG collaboration reported new measurements of CMB anisotropy and polarization spectra~\\cite{boom_polar}. The WMAP mission has also measured CMB polarization spectra ~\\cite{kog_wmap03,pag_wmap06}. Correcting for the systematic effects of a non-circular beam for the polarization spectra is expected to become important. Extending this work to the case CMB polarization is another line of activity we plan to undertake in the near future.  In summary, we have presented a perturbation framework to compute the effect of non-circular beam function on the estimation of power spectrum of CMB anisotropy taking into account the effect of a non-uniform sky coverage (eg., galactic mask). We not only present the most general expression including non-uniform sky coverage as well as a non-circular beam that can be numerically evaluated but also provide elegant analytic results in interesting limits.  As CMB experiments strive to measure the angular power spectrum with increasing accuracy and resolution, the work provides a stepping stone to address a rather complicated systematic effect of non-circular beam functions."
    },
    "0608/astro-ph0608219_arXiv.txt": {
        "abstract": "The dip is a feature in the diffuse spectrum of ultra-high energy (UHE) protons in the energy range  $1\\times 10^{18} - 4\\times 10^{19}$ eV, which is caused by electron-positron pair production on the cosmic microwave background (CMB) radiation. For a power-law generation spectrum $E^{-2.7}$, the calculated position and shape of the dip is confirmed with high accuracy by the spectra observed by the Akeno-AGASA, HiRes, Yakutsk and Fly's Eye detectors. When the particle energies, measured in these detectors, are calibrated by the dip, their fluxes agree with a remarkable accuracy. The predicted shape of the dip is quite robust: it is modified very weakly when the discreteness and inhomogeneities in the source distribution are taken into account, and for different regimes of propagation (from rectilinear to diffusive). The cosmological evolution of the sources, with parameters inspired by observations of Active Galactic Nuclei (AGN), also results in the same shape of the dip. The dip is modified strongly when the fraction of nuclei heavier than protons is high at injection, which imposes some restrictions on the mechanisms of acceleration operating in UHECR sources. The existence of the dip, confirmed by observations, implies that the transition from galactic to extragalactic cosmic rays occurs at $E \\lsim 1\\times 10^{18}$~eV. We show that at energies lower than a characteristic value $E_{\\rm cr} \\approx 1\\times 10^{18}$ eV, determined by the equality between the rate of energy losses due to pair production and adiabatic losses, the spectrum of extragalactic cosmic rays flattens in all cases of interest, and it provides a natural transition to a steeper galactic cosmic ray spectrum. This  transition occurs at some energy below $E_{\\rm cr}$, corresponding to the position of the so-called second knee. We discuss extensively the constraints on this model imposed by current knowledge of acceleration processes and sources of UHECR and compare it with the traditional model of transition at the ankle. ",
        "introduction": "\\label{introduction} A very important step toward unveiling the origin of the sources of UHECR is to identify the range of energies where cosmic rays become mainly of extragalactic origin. For such extragalactic cosmic rays, in the hypothesis of a proton dominated spectrum, the propagation in the intergalactic medium induces three features in the spectrum observed at the Earth: 1) the GZK feature \\cite{GZK}, a suppression of the flux at energies in excess of $\\sim 10^{20}$ eV, due to the photopion production interactions of cosmic rays off the CMB photons; 2) a bump (\\cite{HS85} - \\cite{Stanev00}), due to the accumulation of particles below the kinematic threshold for photopion production. As was shown in \\cite{BG88}, while the bump is present in the calculated spectra of single sources, it almost disappears in the diffuse spectrum, because the position of the bump depends upon the source redshift; 3) A dip (\\cite{HS85} - \\cite{BGG3}), generated due to pair production, $p+\\gamma_{\\rm CMB} \\to p+e^++e^-$, where the target is provided by the CMB photons.  The detection of these features would be a definitive test of the extragalactic origin of UHECR and of the fact that they are mainly protons. Since the detection of the GZK feature requires very large statistics of events and, as stressed above, the bump is almost absent in the diffuse spectrum, at present the spectral feature that can be detected more easily is the dip. As we show here (see also \\cite{BGG3} - \\cite{BGG}), the dip (see Fig.~\\ref{fig:mfactor}) is a quite robust prediction of the calculation and we claim that in fact it might have already been observed by the AGASA, Fly's Eye, HiRes and Yakutsk experiments (see \\cite{agasa} - \\cite{NW} for the data). However, the detectability of the dip as a feature of the propagation of cosmic rays on cosmological scales would also imply that the transition from galactic to extragalactic cosmic rays should not take place at the {\\it ankle}, as has been postulated since the end of '70s, when this feature was discovered in the Haverah Park data (see \\cite{ankle} for the recent works).  The traditional explanation of the transition from galactic to extragalactic cosmic rays invokes the intersection between a steep ($E^{-3.1}$) galactic spectrum and a flat ($E^{-\\alpha}$ with $\\alpha=2-2.3$) extragalactic spectrum, at the {\\it ankle}, located at an energy $E_a \\approx 10^{19}$ eV and identified as a flattening of the spectrum in the data of AGASA, HiRes and Yakutsk detectors (see Fig. \\ref{fig:dips} and \\cite{DeMSt} for a general discussion of the transition).  It is important to stress that in the dip scenario the predicted spectrum flattens below and above the dip location (see Fig~\\ref{fig:mfactor}). The high energy flattening, at $E_a \\approx 1\\times 10^{19}$ eV, reproduces perfectly well the observed {\\it ankle}. The low energy flattening, at $E_{\\rm cr} \\approx 1\\times 10^{18}$ eV, obtained for both cases of rectilinear \\cite{BGG} and diffusive propagation \\cite{AB,Lem,AB1}, provides the transition to a steeper galactic component. Note that this property, the intersection of steep and flat spectral components, is the same for both models of transition.  We demonstrate here (see also \\cite{AB1}) that $E_{\\rm cr}$ is connected with the energy scale $E_{\\rm eq} = 2.3 \\times 10^{18}$ eV, where the rates of pair production and adiabatic energy losses are equal.  The visible transition from galactic to extragalactic cosmic rays occurs at $E_{\\rm tr} < E_{\\rm cr}$, and this energy coincides with the position of the {\\em second knee}  (Akeno - $6\\times 10^{17}$~eV, Fly's Eye - $4\\times 10^{17}$~eV, HiRes - $7\\times 10^{17}$~eV and Yakutsk - $8\\times 10^{17}$~eV). The transition at the second knee was also proposed as a consequence of the study of the propagation of galactic cosmic rays (\\cite{Biermann} - \\cite{Dermer}).  The energy region around $E_{\\rm cr} \\approx 10^{18}$ is also expected to correspond to a change in the chemical composition, from a heavy galactic component to a proton-dominated composition of UHECR. While HiRes \\cite{mass-Hires}, HiRes-MIA \\cite{Hi-Mia} and Yakutsk \\cite{Glushkov00} data support this prediction and Haverah-Park \\cite{HP} data do not contradict it at $E \\gsim (1 - 2)\\times 10^{18}$~eV, the Akeno \\cite{mass-Akeno} and Fly's Eye \\cite{FE} data favor a mixed composition, dominated by heavy nuclei (for a review see \\cite{NW} and \\cite{Watson04}).  In this paper we shall use the number density of UHECR sources estimated from the small-scale clustering in the angular distribution of the arriving particles \\cite{ssc}. In \\cite{ssa1,ssa2} it was found that the observed small-scale clustering implies, in the case of rectilinear propagation of particles, a spatial density of UHECR $n_s \\approx (1 - 3)\\times 10^{-5}~{\\rm Mpc}^{-3}$. Approximately the same space density was found from a study of the small-scale clustering for propagation in magnetic fields \\cite{Sato}. This density is of the same order of magnitude as the density of powerful AGN. It is however worth stressing that the simulated AGASA spectra with this source density are incompatible with the observed AGASA spectrum at the $5\\sigma$ level \\cite{danny1}. The potential of future detectors such as the Pierre Auger Observatory to measure the source density from small scale anisotropies has been discussed in \\cite{danny2}. From the phenomenological point of view, there are some observational indications of AGN as the sources of UHECR \\cite{TT}, although the subject is still matter of much debate.  The paper is organized as follows: in Section \\ref{dip} we discuss the physics behind the formation of the dip, and compare our predictions with the data of several experiments. We discuss also the robustness of the prediction of the dip and some physical phenomena which modify the shape of the dip. In Section \\ref{transition} we address the more specific issue of the transition from galactic to extragalactic cosmic rays, stressing the differences between the {\\it dip scenario} and the {\\it ankle scenario}. In the Appendix we discuss the problems of acceleration relevant for the dip scenario.  We conclude in Section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  The dip is a feature in the extragalactic cosmic ray spectrum, that originates from the Bethe-Heitler pair production of protons on the cosmic microwave background  (\\cite{BGG3},\\cite{BGG}). The dip appears at energy $1\\times 10^{18} - 4\\times 10^{19}$~ eV, with shape practically independent of the discreteness and inhomogeneity in the source distribution, independent of local overdensity and deficit of the sources and of the maximum acceleration energy, independent of the presence or absence of magnetic fields in the intergalactic space, and independent of fluctuations in $p\\gamma$ interactions. The cosmological evolution of UHECR sources, with the parameters taken from observations of AGN, do not affect the dip. The only phenomenon which modifies noticeably the shape of the dip is the presence of a large fraction of nuclei, heavier than hydrogen, in the generation spectrum. We consider the small fraction of nuclei, allowed by the dip shape, as an indication of possible mechanisms of acceleration or injection, operating in UHECR sources (Section \\ref{sec:injection}).  The predicted shape of the dip is in excellent agreement with the data of Akeno-AGASA, Fly's Eye, HiRes and Yakutsk detectors (the data of Auger are inconclusive because of the absence of data at $E< 3\\times 10^{18}$~eV, essential for the dip). The energy calibration of these detectors by the position of the dip results in excellent agreement between the measured fluxes.  To our knowledge, the precision of the agreement between the predictions for the dip, which need only two free parameters, and observations (see Fig.~\\ref{fig:dips}) is the best that ever existed in cosmic ray physics.  This, and the results of energy calibration of the detectors by the position of the dip (see Fig.~\\ref{fig:AgHiYa}), makes improbable that we are observing an accidental agreement between the predictions and the observations.  All the discussion presented so far in the paper was based on purely phenomenological and largely model independent grounds. Clearly, in order to establish a connection with existing theories of acceleration or propagation of cosmic rays, the dip needs to be related to models. In passing we notice however that the shape of the dip agrees well with data when AGN with their observed cosmological evolution are considered as sources of UHECR (Section \\ref{evolution}). A possible, though model dependent, correlation between the maximum acceleration energy and the luminosity of the sources might also solve the problem of the excessive energetics required by the dip scenario (Section \\ref{Emax}). A small fraction of heavy nuclei can also be accommodated in the model (Section \\ref{chemical}).  In the dip scenario the transition from galactic to extragalactic cosmic rays takes place at $E_{\\rm cr} \\approx 1\\times 10^{18}$~eV. The natural character of this transition is guaranteed by the fact that a flat extragalactic spectrum at $E < E_{\\rm cr}$ intersects a steep galactic spectrum at $E > E_{\\rm Fe} \\sim 1\\times 10^{17}$~eV. Observationally the transition occurs at the second knee.  An alternative possibility for the transition is given by the {\\em ankle} model (see Section \\ref{transition}).  The ankle model has many attractive features. It gives a simple and natural picture of the transition as the intersection of a steep galactic spectrum with a very flat extragalactic spectrum, resulting from a generation spectrum as predicted by the 'standard' acceleration models ($E^{-2}$ for non-relativistic shock acceleration and $E^{-2.2}$ for relativistic shocks). These generation spectra have no problem with energetics. The model can also easily accommodate a large fraction of heavy nuclei, if observations show that this is needed.  On the other hand the ankle model has weaknesses: the model requires that the galactic component extends to energies in excess of $10^{19}$ eV, in apparent contradiction with the KASCADE data. We add also that it seems to contradict the most current models of cosmic ray acceleration in galactic sources \\cite{Emax}. The model also has difficulties in explaining the $\\sim 10\\%$ of protons observed in Akeno at $E \\sim 10^{17}$~eV (in the dip model these protons are extragalactic).  In conclusion, we want to emphasize the importance of having detectors working in the energy range $3\\times 10^{17} - 1\\times 10^{19}$~eV, such as possible low-energy extensions of Auger, Telescope Array (TA) and low-energy extension of TA (TALE). These detectors can reliably solve the problem of measuring the mass composition in the transition region. It is important to observe that the two basic models of the transition, at the second knee and at the dip, give drastically different predictions for the chemical composition in the energy interval  $1\\times 10^{18} - 1\\times 10^{19}$~eV: proton-dominated in the dip model and iron-dominated in the ankle model. The updated fluorescent method with measurement of atmosphere transparency for each event can reliably distinguish these two models, in spite of existing experimental and theoretical (models of interaction) uncertainties.  One can add another signature common to both the dip and the ankle scenarios (see \\cite{parizot}) namely the appearance of a jump in $x_{max}$ at energy $E_{tr}$, which equals $\\sim 5\\times 10^{17}$ eV for the dip model and $\\sim 10^{19}$ eV for the ankle model. This feature appears because of a sharp transition from a steep galactic (iron dominated) spectrum to a flat extragalactic (proton dominated) component (see Fig. 10). Such a jump is practically absent or much smoother in the mixed composition model. In contrast with the dip and ankle models, the mixed composition model needs a higher accuracy of the fluorescent method to distinguish the mixed extragalactic composition from the proton-dominated or iron-dominated compositions (there are no reliable predictions for the mass composition of extragalactic UHECR).  One should keep in mind the great predictive power of the {\\em spectrum measurements} for the determination of the mass composition of extragalactic UHECR. The modification factors for protons and different nuclei (see Fig.~\\ref{fig:dip-nucl} for comparison) differ very strongly in the energy region $1\\times 10^{18} - 1\\times 10^{20}$~eV and even small admixtures of nuclei can be discovered as a distortion of the proton modification factor. Note that this possibility exists for all models: dip, ankle and intermediate transition.  In principle, the {\\em anisotropy} discriminates the different models of transition. While the dip-transition model predicts an extragalactic anisotropy at $1\\times 10^{18} - 1\\times 10^{19}$~eV, the ankle models predicts the galactic anisotropy at this energy range. However, in the case of an expected iron-dominated composition of galactic cosmic rays both models predict too small anisotropy.  We think that measurements of the spectrum and chemical composition in the energy region $3\\times 10^{17} - 1\\times 10^{19}$~eV can resolve the existing problem of establishing where the transition from galactic to extragalactic cosmic rays takes place and much experimental effort should be put in aiming at this goal."
    },
    "0608/gr-qc0608012_arXiv.txt": {
        "abstract": "We investigate general relativistic effects associated with the gravitomagnetic monopole moment of gravitational source through the analysis of the motion of test particles and electromagnetic fields distribution in the spacetime around nonrotating cylindrical NUT source. We consider the circular motion of test particles in NUT spacetime, their characteristics and the dependence of effective potential on the radial coordinate for the different values of NUT parameter and orbital momentum of test particles. It is shown that the bounds of stability for circular orbits are displaced toward the event horizon with the growth of monopole moment of the NUT object. In addition, we obtain exact analytical solutions of Maxwell equations for magnetized and charged cylindrical NUT stars. ",
        "introduction": "At present there is no any observational evidence for the existence of gravitomagnetic monopole or so-called NUT charge which may be source of gravitomagnetic field in static case. However there are several attempts and proposals to detect it through astronomical observations in galactic, extragalactic range and in Solar system. For example, the anomalous acceleration of Pioneer satellites~\\cite{pioneer} is attempted to explain through the effect of gravitational field of {\\it magnetic mass} on motion of test particles~\\cite{pioneer_nut1,pioneer_nut2}. In recent papers ~\\cite{rahvar,rahnour}, the theoretical effect of magnetic mass in NUT space on the microlensing light curve and the possibility of magnetic mass detection using the gravitational microlensing technique have been studied.  It is not less interesting to study vacuum axial-symmetric solutions of Einstein equations describing the gravitational field of nonrotating NUT source. Nowadays cylindrically symmetric stationary metrics are analyzed due to the fact that they can be used as simple first order approximation model for describing some astrophysical objects.  Here we will concentrate in studying properties of NUT sources related to the cylindrical symmetry. The NUT metric components generated by the solution to the Newtonian Laplacian equation corresponding to a Newtonian potential of a rod of length $2A$ and mass per unit length being equal to ${1}/{2}$ was presented in~\\cite{gauthoff}. In the paper~\\cite{zonoz97} the metric of cylindrical NUT source in the approximation of weak field has been obtained and studied. In the paper~\\cite{avas-bobo} we have shown the metric of Nouri-Zonoz describing spacetime outside a line gravitomagnetic field could be obtained from Papapetrou solution of vacuum equations of gravitational field.  Physical interpretation of NUT solution was also investigated in the literature. In particular, it was shown in recent paper~\\cite{manko} that the sources of the NUT spacetime are two semi-infinite counter-rotating rods of negative masses and a finite static rod of positive mass. It was also noted that this model is the only possibility for the NUT solution to be stationary and possess zero total angular momentum. Although the metric is not, from the mathematical point of view, cylindrically symmetric, it is physically cylindrically symmetric, which is clear from the gravitoelectric and gravitomagnetic fields which are both explicitly cylindrically symmetric functions.  In our preceding paper~\\cite{bobo} the electromagnetic fields of conducting shell embedded  in the weak gravitational field of cylindrical gravitomagnetic monopole were investigated and analytical solutions to the Maxwell equations in this spacetime were found. Here we extend them to the space time of cylindrical NUT source obtained in~\\cite{manko} as exact solution of Einstein equations without making any assumption on the strength of gravitational field. We are interested to study the motion of test particles and electromagnetic fields in NUT space with the aim to get tools for studying new important general relativistic effects which are associated with nondiagonal components of the metric tensor and have no Newtonian analogues due to their origin.  The paper is organized as follows. In section~\\ref{motion} we study the motion of test particles in spacetime of cylindrical NUT source. We investigate the effective potential and consider the influence of NUT parameter on characteristics of particles motion. In section~\\ref{meq} we solve Maxwell equations in the exterior metric of cylindrical NUT source around (i) magnetized star with the dipolar magnetic field and (ii) electric charge put on cylindrical gravitomagnetic monopole. We summarize our results in section~\\ref{conclusion}.  Throughout, we use the units $G=c=1$ and the metric signature $(-,+,+,+)$. The round and square brackets on the indices denote symmetrization and antisymmetrization respectively, e.g., $A_{(\\mu\\nu)}=\\left(A_{\\mu\\nu}+A_{\\nu\\mu}\\right)/2$ and $A_{[\\mu\\nu]}=\\left(A_{\\mu\\nu}-A_{\\nu\\mu}\\right)/2$. ",
        "conclusions": "\\label{conclusion}  We have investigated the general relativistic effects associated with the possible existence of the gravitomagnetic monopole moment of cylindrical NUT source and in particular studied the influence of the gravitomagnetic monopole moment on the motion of the test particles along circular orbits. These effects could be divided into two parts. First effect is related to the fact that due to the influence of the NUT parameter the particle's circular orbits will be not kept in the equatorial plane. Our aim was mainly related to the investigation of the second effect connected with the influence of the NUT parameter on the circular orbits of the particles. It is shown that under the influence of the NUT parameter the particle's orbits are becoming less stable which can be seen from the table presented in the paper.  Analytical general relativistic expressions for the dipolar magnetic field around NUT star are presented. It is shown that in the linear approximation in NUT parameter there is only effect of mass of the star on the value of stationary magnetic field. In the case of charged star there is strong influence of NUT parameter on the value of monopolar electric field. Our results in principle could be combined with analysis of astrophysical data on electromagnetic fields of compact objects in order to get upper limits on the value of NUT parameter."
    },
    "0608/astro-ph0608439_arXiv.txt": {
        "abstract": "A new cosmological scenario driven by a slow rolling homogeneous scalar field whose exponential potential $V(\\Phi)$ has a quadratic dependence on the field $\\Phi$ in addition to the standard linear term is discussed. The derived equation of state for the field predicts a transient accelerating phase, in which the Universe was decelerated in the past, began to accelerate at redshift $z \\sim 1$, is currently accelerated, but, finally, will return to a decelerating phase in the future. This overall dynamic behavior is profoundly different from the standard $\\Lambda$CDM evolution, and may alliviate some conflicts in reconciling the idea of a dark energy-dominated universe with observables in String/M-theory. The theoretical predictions for the present transient scalar field plus dark matter dominated stage are confronted with cosmological observations in order to test the viability of the scenario. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608654_arXiv.txt": {
        "abstract": "{ {This paper develops a method for obtaining the star formation histories of a mixed, resolved population through the use of color-magnitude diagrams (CMDs).  The method provides insight into the local star formation rate, analyzing the observations of the Hipparcos satellite through a comparison with synthetic CMDs computed for different histories with an updated stellar evolution library.} {Parallax and photometric uncertainties are included explicitly and corrected using the Bayesian Richardson-Lucy algorithm.  We first describe our verification studies using artificial data sets.  From this sensitivity study, the critical factors determining the success of a recovery for a known star formation rate are a partial knowledge of the IMF and the age-metallicity relation, and sample contamination by clusters and moving groups (special populations whose histories are different than that of the whole sample). Unresolved binaries are less important impediments.  We highlight how these limit the method.} {For the real field sample, complete to $M_V < 3.5$, we find that the solar neighborhood star formation rate has a characteristic timescale for variation of about 6 Gyr, with a maximum activity close to 3 Gyr ago. The similarity of this finding with column integrated star formation rates may indicate a global origin, possibly a collision with a satellite galaxy.  We also discuss applications of this technique to general photometric surveys of other complex systems (e.g. Local Group dwarf galaxies) where the distances are well known.} ",
        "introduction": "Since the pioneering works of Bahcall \\& Soneira (\\cite{bah}) and Gilmore \\& Reid (\\cite{gil83}), the comparison with deep star counts has been a primary method to disentangle the stellar components of the Milky Way. However, the number of parameters and the observational uncertainties involved in this kind of analysis are often prohibitive.  In contrast, nearby stars are a well defined sample of disk stars and the distances are now well known so one can extract detailed information about the local star formation rate and how this is connected with the whole Galactic disk.  In this framework, the Hipparcos catalog provides a unique opportunity. Before Hipparcos, the local stellar population of bright stars, in main sequence or in red giant phase, was poorly represented; moreover, the distance uncertainties washed out most of the fine structure of the color-magnitude diagrams (CMDs).  After Hipparcos, it was possible to study the color-magnitude diagram for local stars in a statistical sense, beyond the simple comparison between evolutionary tracks and single stars, a method for which is the subject of this paper.  Published studies based on the Hipparcos local sample have concentrated on two approaches. One is the direct comparison between data and artificial CMDs using a likelihood function, e.g.  Bertelli \\& Nasi (\\cite{bert}) and Schr\\\"oder \\& Pagel (\\cite{scho}). The second is the Bayesian approach, e.g. by Hernandez et al. (\\cite{her2}) and Vergely et al. (\\cite{verg}).  Here we suggest an hybrid technique (for an over-complete description see Cignoni 2006). First, we adopt a {Bayesian treatment of the observational uncertainties}: the CMD is converted into an image (binning process) and a Richardson-Lucy algorithm is used to clean the data (see Cignoni \\& Shore 2006 for details). This ``cleaned'' Hipparcos CMD can then be used to recover the local star formation rate (SFR). This is done in different steps, as described in the following sections: (1) an ensemble of synthetic CMDs is generated using Monte Carlo simulations; (2) a likelihood function is minimized for the comparison between theory and observation; and (3) a confidence limit of the result is evaluated with a bootstrap technique.  In section 2 we describe the properties of the selected volume complete sample and how we removed the known stellar clusters, moving groups, and associations. Sections 3 and 4 describe the method: physical inputs adopted for the stellar evolutionary calculations, Monte Carlo generations of artificial CMDs, likelihood function and bootstrap. In section 5 we apply the algorithm to artificial CMDs, showing which parameters are critical for recovering the star formation rate. Section 6 shows the results for the Hipparcos data. Section 7 discusses the dependence of the recovered SFR from a kinematic selection of the data. Finally in section 9 the results are discussed and compared with previous works available in the literature. ",
        "conclusions": "We have used a selection of the Hipparcos stars to recover the local SFR. The analysis is restricted to the stars within 80 pc and brighter than $M_V=3.5$. Numerical experiments with artificial CMDs show that, at these luminosities, neither the IMF nor the binary fraction are critical inputs, while the possibility to recover the SFR is strongly influenced by the adopted age-metallicity relation.  In particular, this result was checked assuming the observational AMR for the solar neighborhood by Nordstr\\\"om et al. (\\cite{nord}): the simulation with artificial CMDs indicate that most of the information about the underlying SFR is still recoverable.  Finally, we applied the algorithm to real Hipparcos data. In contrast with artificial CMDs, the first problem was the presence of observational uncertainties (due to photometric and parallax errors).  To take these uncertainties into account, we applied to the data the Richardson-Lucy technique as introduced in Cignoni \\& Shore (\\cite{cign}), cleaning the Hipparcos CMD from the observational errors.  Then, assuming the observational AMR by Nordstr\\\"om et al. (\\cite{nord}), we have found the most probable SFR from our sample. The result indicates that \\emph{the recent local SF history of the Galactic disk is increasing from the past to the present with some irregularities}. The mean value increases very steeply from 6-7 Gyr ago up to 2 Gyr, in a way qualitatively similar to the findings of Vergely et al. (\\cite{verg}) and Bertelli \\& Nasi (\\cite{bert}). In particular, this result is is quite independent against kinematic selections, suggesting that: \\begin{enumerate} \\item The local contamination of halo and thick disk stars is negligible in the last 6 Gyr and/or these populations are older than 6 Gyr; \\item In the last 5-6 Gyr, all the stellar generations are well sampled; in other words, the recovered local SFR seems not to be biased by dynamical diffusion and the local volume is not ``depleted'' by old disk stars. \\end{enumerate} The timescale of the recovered SFR seems too long (larger than the dynamical timescale) to be attributed to local events: an accretion of a satellite galaxy is suspected."
    },
    "0608/astro-ph0608462_arXiv.txt": {
        "abstract": "Recently, Bell \\cite{bell04} has reanalysed the problem of wave excitation by cosmic rays propagating in the pre-cursor region of a supernova remnant shock front. He pointed out a strong, non-resonant, current-driven instability that had been overlooked in the kinetic treatments by Achterberg \\cite{achterberg83} and McKenzie \\& V\\\"olk \\cite{mckenzievoelk82}, and suggested that it is responsible for substantial amplification of the ambient magnetic field. Magnetic field amplification is also an important issue in the problem of the formation and structure of relativistic shock fronts, particularly in relation to models of gamma-ray bursts. We have therefore generalised the linear analysis to apply to this case, assuming a relativistic background plasma and a monoenergetic, unidirectional incoming proton beam. We find essentially the same non-resonant instability noticed by Bell, and show that also under GRB conditions, it grows much faster than the resonant waves.  We quantify the extent to which thermal effects in the background plasma limit the maximum growth rate. ",
        "introduction": "The acceleration of cosmic rays at the shock front that bounds a supernova remnant is thought to proceed via the diffusive first-order Fermi mechanism (for reviews, see \\cite{blandfordeichler87,drury83,jonesellison91}). In the shock precursor, where the cosmic rays stream at roughly the shock speed, Alfv\\'en waves can grow as a result of interacting at the cosmic ray cyclotron resonance. In turn, these waves provide the pitch-angle scattering essential for the acceleration process \\cite{mckenzievoelk82}. Using a kinetic approach, Achterberg \\cite{achterberg83} found that the presence of cosmic rays has no significant impact on the nature of the plasma modes. However, Bell \\cite{bell04}, using a hybrid MHD/kinetic theory approach, has recently challenged this conclusion. He identifies a low frequency, non-resonant, electromagnetic mode that can be strongly driven by the current induced in the plasma by the streaming cosmic rays. Simple analytic estimates and MHD simulations both suggest that the nonlinear evolution of this instability leads to very strong amplification of the ambient magnetic field. In turn, this may result in a cosmic-ray acceleration rate that is much more rapid than previously thought \\cite{bell05,belllucek}.  Gamma-ray bursts drive highly relativistic outflows with Lorentz factors of 100 or more \\cite{grbreview}. On interacting with the surrounding medium, a shock front forms, but the mechanism by which this happens is controversial \\cite{spitkovsky05}. Observations of the afterglows of gamma-ray bursts suggest that the ambient magnetic field must be amplified substantially at the shock front. The relativistic Weibel or two-stream instability, that generates a small-scale magnetic field in a previously unmagnetized plasma, has been investigated in this connection \\cite{medvedevloeb99}, but appears to saturate at a relatively low amplitude \\cite{lyubarskyeichler05,wiersmaachterberg04}. In recent years, particle in cell (P.I.C.) simulations have been used in studies of collisionless shocks \\cite{spitkovsky05}, in particular with reference to the Weibel instability ~\\cite{hededalnishikawa,nishikawaetal,silvaetal}. The instability that is presented in this paper has not been identified to date. However, relativistic P.I.C. electron-proton simulations in three dimensions still present a major challenge, even to the best computational resources currently available.  When two electron-proton plasmas moving at high speed relative to each other collide and interpenetrate, the first process that takes place is the equilibration of the electrons \\cite{lerchepohlschlickeiser}. However, this releases only a small fraction of the available free energy. The shock front responsible for the thermalisation of the bulk of the energy is mediated by the interaction of the two counterstreaming proton plasmas.  Thus, Pohl \\& Schlickeiser \\cite{pohlschlickeiser} attacked the shock formation problem starting with a monoenergetic relativistic proton beam penetrating a relatively dense, cool background in a direction aligned with the ambient magnetic field. They then computed the isotropisation rate that results from the resonant excitation of Alfv\\'en waves in the background plasma. Here, we use the same initial conditions, but, because this physical situation resembles that in the precursor region of a SNR shock front, we analyse instead the growth rate of the relativistic analogue of the non-resonant mode discovered by Bell \\cite{bell04}. We find it grows much faster than the resonant mode, and, as in the case of SNR, can be expected to generate a substantial magnetic field transverse to the beam direction. However, we find the instability is quite sensitive to damping by thermal effects, once the background plasma is heated. ",
        "conclusions": "From Fig.~\\ref{modeplot} it can be seen that the $\\varepsilon=1$ mode  (grey curves) resonates with the beam particles when $k$ is close to their inverse gyro-radius, i.e., $\\hat{k}\\approx1$. This is the cyclotron resonant mode considered by, amongst others, Pohl \\& Schlickeiser \\cite{pohlschlickeiser}. However, for low to moderate values of the background temperature, its growth rate is very much slower than that of the non-resonant, current-driven mode (black curves). Near the peak of the unstable band of the non-resonant mode, the waves are strongly modified by the beam and do not resemble Alfv\\'en waves: the phase velocity of the non-resonant mode is lower than that of the resonant mode, and both speeds generally exceed $v_{\\rm A}$ substantially. At very large wavelengths, the modes converge and become Alfv\\'en waves. This also happens for large $k$ values outside the unstable band, but only if these do not resonate with the background protons.  Eqs~(\\ref{bell2}) and (\\ref{bell3}), show that in a cold background with $n_{\\rm b}\\approx n_{\\rm p}$, the non-resonant mode has a maximum growth rate comparable to that of electrostatic oscillations, although at much shorter (parallel) wavelengths. Only a relatively modest temperature $\\Theta_{\\rm p}\\approx\\omega_{\\rm c}/\\omega_{\\rm pp}$ suffices to damp such waves, as can be seen from Eq.~(\\ref{thermaleffects}). Physically, this can be understood by noting that, as the background temperature rises, an increasing number of particles react to the wave perturbation as if they were unmagnetized. This reduces the effective current driving the instability. Ultimately, if all particles react in the same way to the perturbation (i.e., all are magnetized or all are unmagnetized) the plasma behaves as a fully compensated system with no driving current. Eq.~(\\ref{thermaleffects}) also shows that, for fixed temperature, the growth rate scales with the cyclotron frequency of the background protons. For $\\Theta_{\\rm p}\\approx1$ and $n_{\\rm b}\\approx n_{\\rm p}$, the maximum growth rate is approximately equal to $\\omega_{\\rm c}$.  The crucial question of the nonlinear evolution of the system has been discussed in the nonrelativistic case by Bell \\cite{bell04} and Milosavljevi\\'c \\& Nakar \\cite{nakar05}. Saturation can be expected when the currents associated with the growing waves become comparable to the current induced by the beam. Applying this to the relativistic case, one finds for the amplitude $B_{\\rm w}$ of the magnetic field of the perturbations \\eqb \\left|\\vec{k}\\wedge\\vec{B}_{\\rm w}\\right| \\approx q_{\\rm b}n_{\\rm b}c\\beta \\eqe If we assume that the largest wavelength in our system is of the order of the gyroradius of the beam protons, we can estimate the maximum magnetic field energy density to be \\eqb {B_{\\rm w}^2\\over 8\\pi} \\approx \\frac{1}{2}n_{\\rm b} \\Gamma m_{\\rm p}c^2 \\eqe  This suggests that the entire energy content of the beam is, at least initially, transferred into transverse field fluctuations. The efficiency of magnetic field generation by this mechanism (i.e., the ratio of the saturated field energy density to that in the incoming stream) is somewhat higher than that attributed to the Weibel instability in the case of a relativistic shock in a pair plasma \\cite{spitkovsky05}. The Weibel instability is not thought to be effective in mediating shocks in electron/proton plasmas \\cite{lyubarskyeichler05}. However, the efficacy of the instability described here depends strongly on the ability of the plasma to convert the energy input at small scales into a larger scale magnetic field.    \\subsection*{Acknowledgements} This research was jointly supported by Cosmogrid and the Max-Planck-Institut f\\\"ur Kernphysik."
    },
    "0608/astro-ph0608181_arXiv.txt": {
        "abstract": "We present new computations of the equilibrium and non-equilibrium cooling efficiencies and ionization states for low-density radiatively cooling gas containing cosmic abundances of the elements H, He, C, N, O, Ne, Mg, Si, S, and Fe. We present results for gas temperatures between $10^4$ and $10^8$~K, assuming dust-free and optically thin conditions, and no external radiation. For non-equilibrium cooling we solve the coupled time-dependent ionization and energy loss equations for a radiating gas cooling from an initially hot, $\\gtrsim 5\\times 10^6$~K, equilibrium state, down to $10^4$ K. We present results for heavy element compositions ranging from $10^{-3}$ to $2$ times the elemental abundances in the Sun. We consider gas cooling at constant density (isochoric) and at constant pressure (isobaric). We calculate the critical column densities and temperatures at which radiatively cooling clouds make the dynamical transition from isobaric to isochoric evolution. We construct ion ratio diagnostics for the temperature and metallicity in radiatively cooling gas. We provide numerical estimates for the maximal cloud column densities for which the gas remains optically thin to the cooling radiation. We present our computational results in convenient on-line figures and tables (see http://wise-obs.tau.ac.il/$\\sim$orlyg/cooling/). ",
        "introduction": "\\label{introduction} The collisionally controlled ionization states and radiative cooling rates of hot ($10^4-10^8$~K) low-density ``coronal'' gas clouds are crucial quantities in the study of the diffuse interstellar and intergalactic medium, including hot gas in supernova remnants, the Galactic halo, active galactic nuclei, galaxy groups and clusters, and warm/hot ``baryonic reservoirs'' in intergalactic clouds and filaments. In this paper we present new computations of the equilibrium and non-equilibrium cooling efficiencies and ionization states for radiatively cooling gas clouds with temperatures between $10^4$ and $10^8$~K. Our main emphasis is on the non-equilibrium cooling and ionization states that occur when an initially hot gas cools radiatively below $\\sim 5\\times 10^6$ K. Below this temperature, cooling can become rapid compared to ion-electron recombination, and the gas at any temperature tends to remain ``over-ionized'' compared to gas in collisional ionization equilibrium (CIE). Because the energy losses below $\\sim 10^7$~K are dominated by atomic and ionic line emissions from many species, the non-equilibrium cooling rates also differ (are suppressed for over-ionized gas) compared to equilibrium cooling. The ``recombination lag'' and departures from equilibrium depend on the abundances of the heavy elements, and we present results for compositions ranging from $10^{-3}$ to 2 times the elemental abundances in the Sun. We include H, He, C, N, O, Ne, Mg, Si, S, and Fe, and we calculate the ion fractions as functions of temperature, for both equilibrium and non-equilibrium cooling, for all ionization stages of these elements.  The non-equilibrium radiative cooling of hot ($>10^4$ K), highly-ionized gas, with emphasis on the associated time-dependent ionization of the metal ions, is a classical problem first investigated by Kafatos (1973), and subsequently by several authors including Shapiro \\& Moore (1976), Edgar \\& Chevalier (1986), Schmutzler \\& Tscharnuter (1993), Sutherland \\& Dopita (1993), and Smith et al.~(1996). Time dependent cooling in colder ($T\\lesssim 10^4$~K) neutral hydrogen gas was studied earlier by Bottcher et al.~(1970), Jura \\& Dalgarno (1972), and Schwarz et al.~(1972) (see also Dalgarno \\& McCray [1972]). These investigations demonstrated that below a few $10^6$~K recombination lags can indeed lead to significant departures from CIE states and the associated equilibrium cooling rates. Computations of hot gas cooling efficiencies and ionization states assuming CIE have been presented in many papers dating back to House (1964), Tucker \\& Gould (1966), Allen \\& Dupree (1969), Cox \\& Tucker (1969), Jordan (1969), followed by Raymond et al.~(1976), Shull \\& van Steenberg (1982), Gaetz \\& Salpeter (1983), Arnaud \\& Rothenflug (1985), Boehringer \\& Hensler (1989), and more recently Sutherland \\& Dopita (1993), Landi \\& Landini (1999), and Benjamin et al.~(2001).  In this paper we reexamine the fundamental problem of non-equilibrium ionization in a time-dependent radiatively cooling gas. We focus on the behavior for pure radiative cooling with no external sources of heat or photoionization. Recombination lags and non-equilibrium ionization also occur in an initially hot gas undergoing rapid expansion and adiabatic cooling (Breitschwerdt \\& Schmutzler 1994; 1999). We do not consider such dynamical effects here.  Our work is motivated by ultraviolet (UV) and X-ray metal absorption-line spectroscopy of hot gas in the Galactic halo (e.g., Sembach \\& Savage 1992; Spitzer 1996), including, more recently, in discrete ionized ``high-velocity clouds'' (Savage et al.~2002; Fox et al.~2004, 2005, Fox, Savage, \\& Wakker~2006, see also Sternberg, McKee, \\& Wolfire~2002; Gnat \\& Sternberg 2004; Maller \\& Bullock 2004). We are also motivated by recent {\\it Hubble Space Telescope (HST)}, {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)}, {\\it Chandra}, and {\\it XMM-Newton}, detections and observations of hot gas around the Galaxy (Collins et al.~2005; Fang et al.~2006), and/or of gas that may be part of a 10$^5$-10$^7$ K ``warm/hot intergalactic medium'' (WHIM) (Tripp, Savage, \\& Jenkins~2000; Shull, Tumlinson, \\& Giroux~2003; Richter et al.~2004; Sembach et al.~2004; Soltan, Freyberg, \\& Hasinger 2005; Nicastro et al.~2005; Savage et al.~2005, Williams et al.~2006; see however Rasmussen et al.~2006). This includes the temperature range where departures from equilibrium cooling are expected to be important. An intergalactic shock heated WHIM is expected to be a major reservoir of baryons in the low redshift universe (Cen \\& Ostriker 1999; Dav{\\'e} et al.~2001; Furlanetto et al.~2005). Important ions for the detection of such gas include \\ion{C}{4}, \\ion{N}{5}, \\ion{O}{6}, \\ion{O}{7}, \\ion{O}{8}, \\ion{Si}{3}, \\ion{Si}{4}, \\ion{Ne}{8}, and \\ion{Ne}{9}.  The early discussions of Kafatos (1973) and Shapiro \\& Moore (1976) focused on time-dependent isochoric (constant density) cooling at a single metallicity. Sutherland \\& Dopita (1993) calculated isobaric (constant pressure) cooling efficiencies for a range of metallicities, but presented limited results for the time-dependent ion-fractions, and did not consider isochoric cooling. Schmutzler \\& Tscharnuter (1993) considered isochoric cooling for a range of metallicities, and presented ion-fractions only for solar metallicity gas. They presented limited results for isobaric cooling. In our calculations, we use up-to-date atomic data, and we present complete results for the time-dependent ion fractions and cooling efficiencies, for both isochoric and isobaric evolutions, for a wide range of metallicities. Because of the ``$PdV$'' work, isobaric cooling is less rapid than isochoric cooling, so departures from equilibrium are somewhat smaller in the isobaric case.  We consider metallicities ranging from $10^{-3}$ to twice the solar abundances of the heavy elements. In our calculations we assume that any initial dust is rapidly destroyed (e.g.~by thermal sputtering) on a time-scale that is short compared to the initial cooling rate (Smith et al.~1996), and we consider dust-free cooling at fixed gas-phase elemental abundances. The sensitivity of the cooling efficiencies to the metallicity was studied by Boehringer \\& Hensler (1989) for equilibrium cooling, and by Schmutzler \\& Tscharnuter (1993) and Sutherland and Dopita (1993) for non-equilibrium conditions as well. Time-dependent effects are enhanced (diminished) as the cooling efficiencies are increased (reduced) for larger (smaller) gas-phase abundances of the heavy elements. Thus, recombination lags are more significant as the metallicity is increased. For a low-density gas, the non-equilibrium ion fractions and cooling efficiencies are independent of the gas density or pressure, because the ratio of the cooling and recombination times is independent of density. The density scales out of the problem. If the gas is assumed to cool, at either constant density or constant pressure, from an initial equilibrium state at high temperature, the only remaining free parameter in the problem is the metallicity $Z$.  We present results for the optically thin limit in which reabsorption of diffuse ``self-radiation'' may be neglected. We verify previous findings that trapping of hydrogen and helium recombination radiation, in optically thick clouds has only a small effect on the cooling and resulting ionization states. We also provide estimates of the maximal cloud column densities for which reabsorption of radiation from the metals (lines and continua) may be neglected in computations of the gas cooling efficiencies. Elsewhere, we will consider the effects of external fields on the non-equilibrium evolution of radiatively cooling gas. The inclusion of such radiation will introduce a pressure/density dependence via the ratio of the photoionization and recombination times.  In a separate paper we will present computations of the integrated metal ion ``cooling columns'' produced in steady flows of cooling gas (e.g.~Edgar \\& Chevalier 1986; Dopita \\& Sutherland 1996; Heckman et al.~2002) including the effects of ``upstream'' cooling radiation on the ``downstream'' ionization states. Such flows occur, for example, in the radiative (isobaric) cooling layers behind steady shock waves (Shull \\& McKee 1979; Draine \\& Mckee 1993).  The outline of our paper is as follows. In \\S 2 we write down the basic equations that we solve in our computations, and we describe our numerical method. In \\S 3 we discuss isochoric versus isobaric cooling, and we derive critical cloud column densities and temperatures at which the dynamical transition from isobaric to isochoric cooling occurs. In \\S 4 we present our computations of the ionization states as functions of gas temperature, for both CIE and non-equilibrium cooling. For time-dependent ionization we present results for both isochoric and isobaric evolutions. In \\S 5 we present our computations of the radiative cooling efficiencies. We compare the CIE and non-equilibrium cooling efficiencies, and determine the temperatures below which non-equilibrium effects become important for each assumed metallicity. In this paper we have generated a large number of figures and tables containing results for metallicities ranging from $10^{-3}$ to 2 times solar.  Some of the figures and tables appear in the main text, and some are available as online data files. Our data sets can be used to construct ion density-ratio diagnostics for radiatively cooling gas for a wide range of conditions. As an example, in \\S 6 we discuss the evolution of the density ratios \\ion{C}{4}/\\ion{O}{6} versus \\ion{N}{5}/\\ion{O}{6} in radiatively cooling gas. We summarize in \\S 7. ",
        "conclusions": "\\label{sum} In this paper we present new computations of the equilibrium and non-equilibrium cooling efficiencies and ionization states for low-density radiatively cooling gas, containing cosmic abundances of the elements H, He, C, N, O, Ne, Mg, Si, S, and Fe. In these calculations we assume pure radiative cooling, with no gas heating or photoionization by external sources. We present results for gas temperatures, $T$, between 10$^4$ and 10$^8$~K. We assume that the gas is dust free, and we consider metallicities $Z$ ranging from 10$^{-3}$ to 2 times the elemental abundances in the Sun. We carry out our computations using up-to-date rate coefficients for all of the atomic recombination and ionization processes, and the energy loss mechanisms.  For temperatures below $\\sim 5\\times 10^6$~K, where ion-electron recombination lags significantly behind the cooling, we explicitly solve the coupled time-dependent ionization and energy loss equations for the cooling gas. For such gas we assume that the cooling is from an initially hot equilibrium state. The basic equations and our numerical method are presented in \\S 2. We calculate the non-equilibrium cooling efficiencies for constant pressure (isobaric) and constant density (isochoric) evolutions. Departures from collisional ionization equilibrium (CIE) are slightly smaller for isobaric cooling.  In \\S 3 we consider the conditions for isochoric versus isobaric cooling. We compute the critical column densities and temperatures at which the cooling time becomes short compared to the dynamical time, and the transition from isobaric to isochoric evolution occurs. These results are displayed in Figure \\ref{choose-case}, and are based on the cooling efficiencies we present in \\S 5.  Because we exclude photoionization by external radiation, both the equilibrium and time-dependent behavior is independent of the assumed density or pressure. The primary parameter for radiative cooling is the metallicity. Departures from equilibrium are largest for high $Z$ where the ratios of the cooling and recombinations times are smallest. At very low metallicity ($Z\\lesssim 10^{-2}$) the cooling rates, and hence also the recombination lags, are independent of the metallicity. In this limit any departures from CIE approach a universal ``primordial'' form. The results of our equilibrium ionization calculations are displayed in Figure 2 in \\S 4. This figure also shows the ionization states for non-equilibrium isobaric and isochoric cooling, for $Z=1$. Results for other metallicities are presented in on-line tables and figures, as summarized in Table 16. Our computational data set can also be found at our website http://wise-obs.tau.ac.il/$\\sim$orlyg/cooling/.  \\setcounter{table}{15} \\begin{deluxetable}{llcc} \\tablewidth{0pt} \\tablecaption{Ionization and Cooling Data} \\tablehead{ \\colhead{Data} & \\colhead{} & \\colhead{Table} & \\colhead{Figure}} \\startdata \\multicolumn{2}{l}{Ion Fractions:} &&\\\\ ~~~CIE                   & & $2$  & $2$\\\\ ~~~$Z=1$, &Isochoric       & $3$  & $2$  \\\\ ~~~$Z=1$, &Isobaric        & $4$  & $2$  \\\\ ~~~$Z=10^{-3}$, &Isochoric & $5$  & $3$  \\\\ ~~~$Z=10^{-3}$, &Isobaric  & $6$  & $3$  \\\\ ~~~$Z=10^{-2}$, &Isochoric & $7$  & $4$  \\\\ ~~~$Z=10^{-2}$, &Isobaric  & $8$  & $4$  \\\\ ~~~$Z=10^{-1}$, &Isochoric & $9$  & $5$  \\\\ ~~~$Z=10^{-1}$, &Isobaric  & $10$ & $5$  \\\\ ~~~$Z=2$, &Isochoric       & $11$ & $6$  \\\\ ~~~$Z=2$, &Isobaric        & $12$ & $6$  \\\\ \\multicolumn{2}{l}{Cooling Efficiencies:} &&\\\\ ~~~CIE                    && $13$ & $8$  \\\\ ~~~Isochoric       && $14$ & $8$  \\\\ ~~~Isobaric        && $15$ & $8$  \\\\ \\enddata \\label{tables} \\end{deluxetable}  Figure \\ref{cool-res-fig} in \\S 5 displays our results for the equilibrium and non-equilibrium cooling efficiencies for $Z$ between 10$^{-3}$ and $2$. We provide simple, or two-piece, power-law fits (equations [14], [15], and [16]) for the equilibrium and non-equilibrium cooling efficiencies at high and low $Z$. Non-equilibrium cooling is suppressed, by factors of 2 to 4, relative to CIE, because of the generally higher ionization state of the rapidly cooling gas. The familiar peaks in the CIE cooling curve are smeared out for non-equilibrium cooling, because individual ions persist for a broader range of temperatures. Overall, our results are in good agreement with previous such calculations. Detailed difference are mainly due to differences in the input atomic data, and assumed abundances of the heavy elements. We make some explicit comparisons with previous computations in \\S 4 and \\S 5.  Ion ratios are useful as diagnostic probes. In \\S 6 we discuss one example, N~{\\small V}/O~{\\small VI} versus C~{\\small IV}/O{\\small VI}, and show how this ratio evolves in radiatively cooling gas, and how it can be used as a probe of metallicity for realistic non-equilibrium conditions.  In our computations we assume that the cooling gas is optically thin. In the Appendix, we provide numerical estimates for the maximal cloud column densities for which this assumption remains valid. We also investigate how reabsorption of hydrogen and helium recombination radiation in optically thick clouds alters the cooling rates and associated ionization states in non-equilibrium cooling gas. For high metallicity the effects of trapping are small, but become more significant ($\\sim$~factors of 2) for low $Z$ gas.  In a companion paper we will present computations of the metal ion ``cooling columns'' produced in steady flows of radiative cooling gas, such as occur in post shock cooling layers, including the effects of ``upstream'' cooling radiation on the ``downstream'' ionization states. We will also consider how photoionization by external background radiation fields alters the ionization and thermal evolution of radiatively cooling gas such as we have considered here."
    },
    "0608/astro-ph0608148_arXiv.txt": {
        "abstract": "We present Spitzer Space Telescope observations of the ``evolved starless core'' L1521F which reveal the presence of a very low luminosity object ($L <$0.07 \\Lsun).  The object, L1521F-IRS, is directly detected at mid-infrared wavelengths ($>$5 \\micron) but only in scattered light at shorter infrared wavelengths, showing a bipolar nebula oriented east-west which is probably tracing an outflow cavity.  The nebula strongly suggests that L1521F-IRS is embedded in the L1521F core.  Thus L1521F-IRS is similar to the recently discovered L1014-IRS and the previously known IRAM 04191 in its substellar luminosity and dense core environment.  However these objects differ significantly in their core density, core chemistry, and outflow properties, and some may be destined to be brown dwarfs rather than stars. ",
        "introduction": "\\label{sec-intro}  Studies of starless dense cores just before they form a protostar are important as they provide the initial conditions for star formation. However, identifying this stage of evolution is not easy.  The majority of cores are starless in the sense that they do not harbor an IRAS point source to a sensitivity of $L \\sim$ 0.1\\Lsun($d$/140 pc)$^2$ (Myers et~al.\\ 1987; Lee \\& Myers 1999), but which starless cores are close to forming a star?  Recently attempts have been made to identify a group of ``evolved'' starless cores with physical, chemical and kinematic properties (e.g., high central densities, molecular depletion, infall asymmetry) that suggest they are close to star formation (Crapsi et~al.\\ 2005a; Kirk et~al.\\ 2005).  It is possible that some of these cores have already begun star formation, containing central sources whose luminosity is too low to have been detected by IRAS.  They are prime targets for sensitive mid-infrared searches for very low luminosity objects with the Spitzer Space Telescope (hereafter Spitzer).  Surprisingly, the first ``starless'' core observed by the Spitzer Legacy program ``From Molecular Cores to Planet Forming Disks'', or c2d (Evans et al.\\ 2003), L1014, was found to contain an embedded source, L1014-IRS with a very low luminosity  $L <$ 0.1 \\Lsun\\ (Young et~al.\\ 2004; Bourke et~al.\\ 2005; Huard et~al.\\ 2006).  Interestingly, L1014 shows very little evidence of being an ``evolved'' core (Crapsi et al.\\ 2005b).  If L1014-IRS and similar objects have accretion rates which are significantly lower than that expected of low mass protostars, then some are likely to be proto-brown-dwarfs.  Thus, L1014-IRS and similar discoveries may pose a significant challenge to our understanding of star formation.  In this letter we report on our Spitzer observations of the object L1521F-IRS which confirm that it is embedded within the dense core L1521F in Taurus (140 pc) and has a very low luminosity.  L1521F-IRS is of particular interest, as L1521F (a.k.a.\\ MC27, Mizuno et~al.\\ 1994; Codella et~al.\\ 1997; Lee, Myers \\& Tafalla 2001) is one of the two best examples of an evolved ``starless'' core, along with L1544 (Crapsi et~al.\\ 2004, 2005a).  L1521F shows a high central density ($\\sim10^6$ \\cc), infall asymmetry, molecular depletion, and enhanced deuterium fractionation (Onishi, Mizuno \\& Fukui 1999; Shinnaga et~al.\\ 2004; Crapsi et~al.\\ 2004). Observations of CO 2-1 show no clear evidence for structured (i.e., bipolar) outflow emission, but line wings are seen in HCO$^+$ 3-2 (Onishi et~al.\\ 1999), which are spatially compact ($<30''$) and may be due to outflow emission similar to that seen from L1014-IRS. ",
        "conclusions": "With the Spitzer Space Telescope we have discovered a low luminosity object L1521F-IRS ($L_{int} = 0.05 \\Lsun$, with an upper limit of $\\sim 0.07$ \\Lsun) within the ``evolved starless core'' L1521F in Taurus.  At short wavelengths ($<$ 5 \\micron) a well defined bipolar scattered light nebula is seen, suggesting the presence of a molecular outflow like those observed from low mass protostars.  L1521F-IRS is located at the density peak of the core, and its luminosity is similar to other low luminosity objects IRAM 04191+1522 and L1014-IRS.  It is unclear at present whether L1521F-IRS will evolve to form a low mass star or a brown dwarf.  These results do not support the idea that core properties correlate with features of early stellar evolution.  Some cores, including L1544, have hallmarks of high density, including freeze-out chemistry, but harbor no embedded sources, while similar cores do have internal sources, such as L1521F. Other cores, including L1014, have modest densities yet harbor VeLLOs.  Still other cores, such as L1521F and IRAM 04191, have similar gas properties and objects of similar luminosity, but their embedded sources have very different mass outflow rates.  It is not clear where VeLLOs fit within the current paradigm of low mass star formation, and it is difficult to imagine a single evolutionary scenario which can explain them all.  The picture may become clearer once a large enough sample of these objects has been studied."
    },
    "0608/astro-ph0608238_arXiv.txt": {
        "abstract": "We investigate structural properties of old, metal-poor globular clusters (GCs) formed at high redshifts ($z>6$) and located inside and outside virialized galaxy-scale halos in clusters of galaxies with the total masses of $M_{\\rm CL}$ based on high-resolution cosmological simulations with models of GC formation. We mainly derive the parameter dependences of physical properties of intracluster GCs (ICGCs) based on the results of 14 cluster models. Our principle results are summarized as follows.   (1) ICGCs are formed as a result of tidal stripping of GCs initially within galaxy-scale halos during hierarchical growth of clusters via halo merging. These ICGCs comprise 20-40 \\% of all GCs in clusters with $1.0 \\times 10^{14} {\\rm M}_{\\odot} \\le M_{\\rm CL} \\le 6.5 \\times 10^{14} {\\rm M}_{\\odot}$ and the number fraction of ICGCs does not depend on $M_{\\rm CL}$ for the above cluster mass range.   (2) The projected radial density profiles (${\\Sigma}_{\\rm GC}$) of ICGCs in clusters with different $M_{\\rm CL}$ can be diverse, though ICGCs have inhomogeneous, asymmetric, and somewhat elongated distributions in most models. If ${\\Sigma}_{\\rm GC} (R) \\propto R^{\\alpha}$, $\\alpha$ ranges from $\\approx -1.5$  to $\\approx -2.5$ for GCs in clusters with the above mass range.    (3) Although total number of GCs within the central 0.05 Mpc ($N_{\\rm GC,0.05}$) and 0.2 Mpc ($N_{\\rm GC,0.2}$) are diverse in different clusters, they can depend weakly on $M_{\\rm CL}$ in such a way that both $N_{\\rm GC,0.05}$ and $N_{\\rm GC,0.2}$ are likely to be larger for clusters with larger $M_{\\rm CL}$.  (4) Total number of GCs per cluster masses (specific frequency of GCs for clusters of galaxies) are more likely to be larger in more massive clusters,  mainly because a larger number of earlier virialized objects can be located in more massive clusters.   (5) Radial density distributions of all GCs including ICGCs and galactic GCs have steeper profiles than those of the background dark matter halos in the central regions of clusters ($R<200$kpc) with different  $M_{\\rm CL}$.  (6) Spatial distributions of old GCs in clusters can depend on the truncation epoch of GC formation ($z_{\\rm trun}$) such that they can be  steeper and more compact in the models with higher $z_{\\rm trun}$.  (7) The mean metallicity of ICGCs in a cluster can be smaller than that of GCs within the cluster member galaxy-scale halos by  $\\sim 0.3$ in [Fe/H]. Metallicity distribution functions (MDFs)  of ICGCs show peak values around [Fe/H] $\\sim -1.6$ and  do not have remarkable  bimodality. ",
        "introduction": "Structural, kinematical, and chemical properties of globular cluster systems (GCSs) have long provided valuable clues to the better understanding of how galaxies form and evolve (e.g.,  Harris 1991; West et al. 2004). In particular, the specific frequencies ($S_{\\rm N}$), colour bimodality, and structural properties of GCSs have been discussed extensively in many different contexts of galaxy formation (e.g., Searle \\& Zinn 1978; Forbes et al. 1997; Ashman \\& Zepf 1998; Brodie et al. 1998;  C\\^ote et al. 2001; Beasley et al. 2002; Bekki et al. 2002). The physical properties of globular clusters (GCs) however have not been so far considered as fossil records that contain valuable information on {\\it how groups and clusters of galaxies were formed and evolved}. Furthermore, GC properties characteristics of giants Es located in the central regions of  clusters (e.g., very high $S_{\\rm N}$) have not been extensively discussed in the context of cluster formation processes via hierarchical merging of smaller groups and clusters of galaxies.  \\begin{table} \\centering \\begin{minipage}{75mm} \\caption{Meaning of acronym} \\begin{tabular}{|l||l|} % INGC &  INtergalactic GCs \\\\ ICGC &  IntraCluster GCs \\\\ IGGC &  IntraGroup GCs \\\\ GGC &  Galactic GCs \\\\ \\end{tabular} \\end{minipage} \\end{table}  \\begin{table*} \\centering \\begin{minipage}{185mm} \\caption{Brief summary of the results for each model} \\begin{tabular}{ccccccccc} (1)&(2)&(3)&(4)&(5)&(6)&(7)&(8)&(9)\\\\ {Model no. \\footnote{The results of the cluster-scale halos with the total masses ($M_{\\rm CL}$) larger than $10^{14} {\\rm M}_{\\odot}$ are described.}} & $M_{\\rm CL}$ ($\\times 10^{14} {\\rm M}_{\\odot}$)  & {$N_{\\rm GC,CL}$ \\footnote{Total number of GCs in a cluster-scale halo.}} & {$S_{\\rm N,CL}$% \\footnote{Specific frequency of cluster GCs. This is defined as total GC number per mass for a cluster of galaxies and given in  units of $ {\\rm number} / 10^{10} {\\rm M}_{\\odot}$.}} & $R_{10}$   (Mpc) & $R_{50}$   (Mpc) & $f_{\\rm ICGC}$  & {$N_{\\rm GC,0.05}$ % \\footnote{Total number of GCs within 0.05 Mpc.}} & {$N_{\\rm GC,0.2}$ % \\footnote{Total number of GCs within 0.2 Mpc.}} \\\\ CL1 &  6.5 & 6868 &  0.11 &  0.026 & 0.76 &  0.29 & 1098 & 2084 \\\\ CL2  & 3.1 & 3449 &  0.11 &  0.014 &  0.21 & 0.31 &  962 &  1699 \\\\ CL3  & 2.9 & 3248  & 0.11 &  0.026 &  0.51 & 0.40 &  479 & 947 \\\\ CL4  & 2.3 & 3788 & 0.17 & 0.012 & 0.12 & 0.18 & 1348 & 2262 \\\\ CL5  & 2.0 & 2251  & 0.11 &  0.013 &  0.50  & 0.27 & 597 & 944 \\\\ CL6  & 1.8 & 2214 &  0.13 & 0.011  & 0.33 & 0.24 & 766 & 1052 \\\\ CL7  & 1.7 & 2197 &  0.13 &  0.068 &  0.32 & 0.37 & 145 &  762\\\\ CL8  & 1.5 & 1905 &  0.12 & 0.030 &  0.23 & 0.27 & 292 & 883 \\\\ CL9  & 1.4 & 1685 &  0.12 &  0.013 &  0.24  & 0.25 & 470 & 801 \\\\ CL10 &  1.3 & 1649 &  0.13 &  0.028  &  0.31 & 0.29 & 311 & 717 \\\\ CL11 &  1.2 & 1347 &  0.11 &  0.028 &  0.23 & 0.29 & 224 & 620 \\\\ CL12 &  1.1 & 795 &  0.07 &  0.035 &  0.37 & 0.35 & 104 & 235 \\\\ CL13  &  1.0 & 1580  & 0.15 &  0.029 &  0.33 & 0.31 & 248 & 640 \\\\ CL14 &  1.0 &  1042 &  0.10 &  0.028 &  0.34 & 0.28 & 186 & 444 \\\\ \\end{tabular} \\end{minipage} \\end{table*}      Observational studies of GCs in clusters of galaxies have suggested that there can be a population of GCs that are bounded by cluster gravitational potentials rather than those of cluster member galaxies (e.g., White 1987; Bassino et al. 2002, 2003; Jord\\'an et al. 2003). Mar\\'in-Franch \\& Aparicio (2003) investigated whether these intracluster GCs (``ICGCs'')  exist in the Coma cluster using surface brightness fluctuations, and concluded that ICGCs are highly unlikely to exist in the Coma. West et al. (1995) suggested that the observed large variations in GC number (or $S_{\\rm N}$) between central giant Es in different clusters are due to the existence of a population of GCs that are not bounded to individual galaxies but move freely throughout the core of clusters galaxies. If very high $S_{\\rm N}$ in central Es in some clusters (e.g., Abell 2052) are due to  larger number of ICGCs in the clusters, it is an essentially important question why only some clusters of galaxies contain larger number of ICGCs.    A new type of sub-luminous and extremely compact ``dwarf galaxy'' has been recently discovered in an `all-object'  spectroscopic survey centred on the Fornax Cluster (Drinkwater et al. 2000; 2003). These ``dwarf galaxies'', which are members of the Fornax Cluster, have intrinsic sizes of $\\sim$ 100 pc and absolute $B$ band magnitude  ranging from $-13$ to $-11$ mag and are thus called ``ultra-compact dwarf'' (UCD) galaxies. Structural and kinematical studies of this new population of very bright ``star clusters'' have also suggested  that these clusters can be also freely floating intracluster objects (Mieske et al. 2004; Drinkwater et al. 2005; Jones et al. 2005). The observed compact spatial distribution and smaller velocity dispersion of this possible intracluster population in the Fornax cluster has not been clearly explained in a self-consistent manner by previous theoretical studies (e.g., Jones et al. 2005).      Physical properties of intracluster stellar objects such as ICGCs and PNe are considered to be sensitive to dark matter properties and cluster-related physical processes (e.g., tidal stripping of GCs and hierarchical growth of clusters) and thus provide some fossil information on formation of galaxies and clusters of galaxies (e.g., Arnaboldi 2004 for a recent review). However there has been little theoretical and numerical works  carried out as to how ICGCs are formed in clusters environments (e.g., Forte et al. 1982; Muzzio 1987; Bekki et al. 2003). These previous models showed that tidal stripping of GCs from cluster member galaxies though galaxy-galaxy and galaxy-cluster interaction is a mechanism for ICGC formation. These previous works however used {\\it fixed} gravitational potentials of already virialized clusters and accordingly could not discuss {\\it how ICGCs in  a cluster  are formed as the cluster grows through hierarchical merging of smaller groups and clusters.}    These previous observational and theoretical studies have raised many questions, the most significant being: (i) How are ICGC  formed under the currently favored cold dark matter (CDM) theory of galaxy formation ? (ii) What physical properties can  ICGCs  have if their formation is closely associated with hierarchical formation of clusters ? (iii) Are there any correlations between physical properties of ICGCs and global properties (e.g., masses and X-ray temperature) of their host clusters of galaxies ? (iv) Is the origin of very high $S_{\\rm N}$ of central giants Es in some clusters closely associated with the formation processes of clusters of galaxies (e.g., galaxy accretion from fields) ? and (v) Are there any differences in radial distributions and kinematics between UCDs and ICGCs ?   In order to answer these questions in a fully self-consistent manner, we need to model both (1) dynamical evolution of GCs  (e.g., tidal stripping and accretion of GCs) during hierarchical growth of clusters of galaxies and (2) formation of GCs at low and high redshifts in clusters environments. Although recent high-resolution cosmological N-body simulations with total particle number of $ \\sim 10^8$ (e.g., Yahagi \\& Bekki 2005; YB) can allow us to address the above point (1), our poor understanding of the above point (2) on physical conditions required for GC formation in galaxies (e.g., Ashman \\& Zepf 1998) would make it difficult for us to have robust conclusions for the above five problems. However,  considering recent significant development in observational studies on GCs in nearby clusters (e.g., Dirsch et al. 2003; Jord\\'an et al. 2003; Richtler et al. 2004), it is still doubtlessly worthwhile to provide some theoretical predictions that can be tested against these latest and future observational results.       The purpose of this paper  is thus to investigate extensively physical properties of GCs in clusters of galaxies based on high-resolution cosmological simulations that can follow both hierarchical growth of clusters through merging of smaller subhalos and dynamical evolution of old GCs. We particularly investigate global ($\\sim$ Mpc scale) density distributions and kinematics both for ICGCs and for GCs that are within galaxy-scale halos in clusters (referred to as ``GGCs'' for convenience). Because of some uncertainty in the modeling of the formation processes of young, metal-rich GCs (e.g., Bekki et al. 2002), we investigate physical properties of {\\it cluster GCs that were formed at high redshift $z>6$}. Accordingly, the present study is regarded as  only the first step toward better understanding the nature of GCs in clusters of galaxies: The results of our simulations may well be compared with observations on ``blue GCs'' in  a more reasonable way.    The plan of the paper is as follows: In the next section, we describe our  numerical models for dynamical evolution of GCs in forming clusters of galaxies. In \\S 3, we present the numerical results mainly on radial distributions  and kinematics of GCs in clusters with different masses. In \\S 4, we discuss (1) the origin of very high $S_{\\rm N}$ of giant Es in the central regions of some clusters, (2) the origin of UCDs, and (3) the importance of the epoch of the truncation of GC formation by reionization in spatial distributions of cluster GCs. These points were not discussed at all in YB. We summarize our  conclusions in \\S 5.         Throughout this paper, we use ICGCs rather than ICG (West et al. 1995) in order to distinguish between intragroup GCs (IGGCs) and intracluster GCs (ICGCs). GCs within any galaxy-scale halos in a cluster-scale halo are referred to as GGCs so that they can be distinguished from ICGCs (Note that these GGCs are {\\it not} the Galactic GCs). GCs that were formed within subhalos at $z>6$ yet are not within any virialized halos at $z=0$ can be regarded as ``intergalactic'' (van den Bergh 1958) or ``intercluster'' (``intergroup'') GCs (YB). Intergalactic, intergroup, and intercluster GCs are simply referred to as INGCs, because these three GC populations would be difficult to be distinguished with one another in observations. For clarity and convenience, the meanings of these acronym are given in Table 1. ",
        "conclusions": "We have investigated radial density profiles and kinematics of GCs located in clusters with the total masses ($M_{\\rm CL}$) larger than $10^{14} {\\rm M}_{\\odot}$ based on high-resolution cosmological N-body simulations with a formation model of old GCs. Although YB has already described some preliminary results for only one cluster model, this paper first discussed the dependences of physical properties of ICGCs on $M_{\\rm CL}$ based on the results on 14 cluster models. We summarize our principle result as follows.  (1) GCs located outside the tidal radii of galaxy-scale halos in clusters of galaxies are formed owing to tidal stripping of GCs initially within smaller halos during hierarchical growth of clusters via halo merging. These intracluster GCs (ICGCs) comprise 20-40 \\% of all GCs in clusters with $1.0 \\times 10^{14} {\\rm M}_{\\odot} \\le M_{\\rm CL} \\le 6.5 \\times 10^{14} {\\rm M}_{\\odot}$ and the number fraction ($f_{\\rm ICGC}$) does not depend on $M_{\\rm CL}$ for the above cluster mass range.   (2) The projected radial density profiles (${\\Sigma}_{\\rm GC}$) of all GCs in clusters with different $M_{\\rm CL}$ can be diverse. If ${\\Sigma}_{\\rm GC} (R) \\propto R^{\\alpha}$, $\\alpha$ ranges from $\\approx -1.5$ to $\\approx -2.5$ for GCs in clusters with the above mass range. ${\\Sigma}_{\\rm GC}$ is more likely to be significantly flatter in ICGCs than in GCs  located in virialized galaxy-scale halos (i.e., GGCs) in clusters. There are no remarkable difference in kinematical properties (e.g., $V_{\\rm r}$ distributions) between GGCs and ICGCs.  (3) Two-dimensional distributions of ICGCs are inhomogeneous, asymmetric, and somewhat elongated, in particular, in the outer parts of cluster for most models. The simulated inhomogeneous distributions of ICGCs suggest that observations based on wide-field imaging of GCs in clusters are doubtlessly worthwhile to understand real radial distributions of ICGCs.  (4) Although total number of GCs within the central 0.05 Mpc ($N_{\\rm GC,0.05}$) and 0.2 Mpc ($N_{\\rm GC,0.2}$) are diverse in different clusters, they can depend weakly on $M_{\\rm CL}$ in such a way that both $N_{\\rm GC,0.05}$ and $N_{\\rm GC,0.2}$ are likely to be larger for clusters with larger $M_{\\rm CL}$. This result implies that central giant ellipticals (Es) in more massive clusters are likely to show a larger number of  old, metal-poor GCs compared with those  in less massive clusters. This result also provides a new clue as to why  some central Es in clusters show high $S_{\\rm N}$.  (5) ${\\Sigma}_{\\rm GC}$ of all GCs in the central 200 kpc in clusters is also diverse with the slopes $\\alpha$ ranging from $\\approx -1.0$ to $\\approx -2.0$. The derived flat distributions of GCs in the central region of clusters suggest that the origin of the observed flat distributions of GCs around central giant Es in clusters (e.g., M87 in the Virgo) can be understood in terms of hierarchical growth processes of the cluster central regions via merging of halos with old GCs.    (6) The distributions of GCs depends on $z_{\\rm trun}$, after which GC formation is assumed to be truncated by some physical processes at high redshifts. For example, the projected radial GC  distributions are more compact and steeper for the cluster models with higher $z_{\\rm trun}$. ICGC fraction can be smaller for the models with higher $z_{\\rm trun}$. This result suggests that if the suppression mechanism of GC formation is due to reionization (Santos 2003; Bekki 2005; Moore et al. 2005), the GC distributions in clusters can provide some useful constrains on the epoch of reionization.  (7) If ultra-compact dwarf galaxies (UCDs), which can be regarded as very massive INGCs, originate from  nuclei of nucleated dwarfs (dE,Ns) that were virialized very high $z$ ($z>10$), the origin of the observed compact spatial distributions can be more clearly understood.   (8) About 1\\% of all GCs formed before $z>6$ are not located within any virialized galaxy-scale, group-scale, and  cluster-scale halos and thus can be regarded as INGCs that are freely drifting in intergroup and intercluster regions. This result is already discussed extensively  by YB. Physical properties (e.g., metallicity distributions and luminosity function) of these intergroup and intercluster GCs can be  quite different from those of GCs within galaxies.   (9) The mean metallicity of ICGCs in a cluster can be smaller than that of GCs within the cluster member galaxy-scale halos by $\\sim 0.3$ in [Fe/H]. Metallicity distribution functions (MDFs)  of ICGCs show peak values around [Fe/H] $\\sim -1.6$ with a higher fraction of metal-poor GCs (with [Fe/H] $<-1.6$). The MDFs do not have remarkable  bimodality that is observationally suggested to be common features of MDFs in bright galaxies.  (10) Radial density profiles of ICGCs and galactic GCs in clusters of galaxies can be significantly steeper than those of the background dark matter halos in the central regions ($R<200$kpc) of the clusters with different  $M_{\\rm CL}$. This is due mainly to the truncation of GC formation at high $z$ ($>6$) in the simulations. The simulated GC profiles can be overall consistent with latest observations on radial density profiles of GCs and UCDs in the Fornax cluster of galaxies."
    },
    "0608/astro-ph0608524_arXiv.txt": {
        "abstract": "The universe provides numerous extremely interesting astrophysical sources of synchrotron X radiation. The \\chandra\\ X-ray Observatory and other X-ray missions provide powerful probes of these and other cosmic X-ray sources. \\chandra\\ is the X-ray component of NASA's Great Observatory Program which also includes the Hubble Space telescope, the Spitzer Infrared Telescope Facility, and the now defunct Compton Gamma-Ray Observatory. The \\chandra\\ X-Ray Observatory provides the best angular resolution (sub-arcsecond) of any previous, current, or planned (for the foreseeable near future) space-based X-ray instrumentation. We present here a brief overview of the technical capability of this X-Ray observatory and some of the remarkable discoveries involving cosmic synchrotron sources. ",
        "introduction": "In this section we present an overview of the technical capability of the Observatory. For more details, please see Weisskopf et al. (2003)\\footnote{See also http://asc.harvard.edu/proposer/POG/index.html}.  \\subsection{The Orbit and the Spacecraft}  The \\chandra\\ X-Ray Observatory began its journey into space on July 23, 1999 using the Space Shuttle Columbia for the initial ascent. An upper stage and integral propulsion then placed the observatory into an elliptical orbit with an initial apogee of 140,000 km (1/3 of the distance to the moon) and an initial perigee 10,000 km (about two earth radii). In this orbit, the satellite (Figure~\\ref{f:spacecraft}\\footnote{Pictures that are publicly available at the \\chandra\\ web site at http://chandra.harvard.edu have credits labeled ``Courtesy of NASA''}) is above the radiation belts for more than 75\\% of the 63.5-hour orbital period, allowing for uninterrupted observations up to more than 2 days.  \\begin{figure} \\begin{center} \\epsfysize=7.0cm \\epsfbox{spacecraft.ps} \\caption{Drawing of the \\chandra\\ X-Ray Observatory with major components labeled. The spacecraft has a 4800-kg mass without propellant, 13-m length and 19.5-m wingspan. The separation between the telescope's optical node and the detector focal plane is 10 m. Courtesy NASA \\label{f:spacecraft}} \\end{center} \\end{figure}  \\subsection{The Optics}  The heart of the Observatory is its telescope made of four concentric, precision-figured, superpolished, grazing-incidence Wolter (type-1) telescopes. The design is similar to that used on the previous X-Ray observatories --- Einstein and \\rosat\\ ---  but of much higher quality, larger diameter, and longer focal length. The telescope's on-axis point spread function has a full-width at half-maximum of about 0.2 arcsec and a half-power diameter less than 1 arcsec. There is a relatively mild dependence on energy, resulting from diffractive scattering by the optic's low ($< 3$-\\AA) surface microroughness.  \\subsection{The Gratings}  Just behind the telescope and along the path to the focal plane are 2 objective transmission gratings (OTGs) --- the Low-Energy Transmission Grating (LETG) and the High-Energy Transmission Grating (HETG). Positioning mechanisms are used to insert either OTG (or none) into the converging beam when commanded.  The Space Research Organization of the Netherlands (SRON) and the Max-Planck-Instit\\\"ut f\\\"ur extraterrestrische Physik (MPE) provided the LETG. With free-standing gold bars of about 991-nm period, the LETG allows for high-resolution spectroscopy ($E/\\Delta E$ $>1000$) between 175 and 80 \\AA\\/ (0.07 -- 0.15 keV) in first order and moderate resolving power at shorter wavelengths. The LETG~wavelength range accessible with the HRC-S as the detector is 175 -- 1.2 \\AA~(0.07 -- 10 keV).  MIT provided the HETG. The HETG uses 2 types of grating~--- the Medium-Energy Gratings (MEG), which, when inserted are behind the X-ray telescope's 2 outermost shells, and the High-Energy Gratings (HEG), the 2 innermost shells. The HETG uses polyimide-supported gold bars of 400-nm and 200-nm periods and provides high-resolution spectroscopy from 0.4 to 4 keV (MEG, 30 to 3 \\AA) and from 0.8 to 8 keV (HEG, 15 to 1.5 \\AA) with ACIS-S as the detector.  \\subsection{The Focal Plane Cameras \\label{s:cameras}}  The Pennsylvania State University (PSU) and working with the Massachusetts Institute of Technology (MIT) designed and fabricated the Advanced CCD Imaging Spectrometer (ACIS), with CCDs produced by MIT's Lincoln Laboratory. Made of a 2-by-2 array of large-format, front-illuminated (FI), 2.5-cm-square CCDs, ACIS-I provides high-resolution spectrometric imaging over a 17-arcmin-square field of view. ACIS-S, a 6-by-1 array of 4 FI CCDs and two back-illuminated (BI) CCDs, mounted along the dispersion direction of \\chandra's transmission gratings serves as the primary read-out detector for the HETG. In addition, using the one BI CCD which can be placed at the aimpoint of the telescope, ACIS-S also provides high-resolution spectrometric imaging extending to lower energies but over a smaller (8-arcmin-square) field than ACIS-I. The spatial resolution for imaging with ACIS~is limited by the physical size of the CCD~pixels (24.0 ${\\mu}$m square $\\approx$0.492 arcsec). Each CCD has 1M pixels.  The Smithsonian Astrophysical Observatory (SAO), provided the other focal-plane imager --- the High Resolution Camera (HRC). Made of a single 10-cm-square microchannel plate (MCP), the HRC-I provides high-resolution imaging over a 30-arcmin-square field of view. The HRC-S, 3 rectangular MCP segments (3-cm $\\times$ 10-cm each) mounted end-to-end along the grating dispersion direction, serves as the primary read-out detector for the LETG. A mechanical translation stage is moves the appropriate detector (ACIS-I, ACIS-S, HRC-I, HRC-S) to the focal point. ",
        "conclusions": "Designed for three years of operation, the \\chandra\\ X-ray Observatory has now been operating for more than seven years. The Observatory's capability for high-resolution imaging has enabled detailed studies of the structure of extended X-ray sources, including supernova remnants, astrophysical jets, and hot gas in galaxies and clusters of galaxies. In addition to mapping the structure of extended sources, the superb angular resolution permits studies of discrete sources which would otherwise be impossible. From planetary systems to deep surveys of the faintest and most distant sources, the scientific results from \\chandra\\ operation have been outstanding. Equally important have been \\chandra's unique contributions to high-resolution dispersive spectroscopy. The high spectral resolution of the \\chandra\\ gratings isolate individual lines from the myriad of spectral lines, which would overlap at lower resolution. Through all these observations, users have addressed, and continue to address, the most exciting topics in contemporary high-energy astrophysics.  I acknowledge the major contributions to the success of the Observatory made by the scientists and engineers associated with the instrument teams, the NASA Project at Marshall Space Flight Ceneter, the many contractors, and the Chandra X-Ray Center, with special thanks to its Director, Dr. H. Tananbaum. I gratefully acknowledge conversations with Drs. D. Worrall, H. Marshall, and S. O'Dell who kindly made comments to a draft of this article. Finally, I wish to acknowledge the tremendous contributions of the now-deceased Telescope Scientist, Dr. Leon Van Speybroeck."
    },
    "0608/hep-th0608219_arXiv.txt": {
        "abstract": "We study the general chaotic features of dynamics of the phantom field modelled in terms of a single scalar field conformally coupled to gravity. We demonstrate that the dynamics of the FRW model with dark energy in the form of phantom field can be regarded as a scattering process of two types: multiple chaotic and classical non-chaotic. It depends whether the spontaneously symmetry breaking takes place. In the first class of models with the spontaneously symmetry breaking the dynamics is similar to the Yang-Mills theory. We find the evidence of a fractal structure in the phase space of initial conditions. We observe similarities to the phenomenon of a multiple scattering process around the origin. In turn the class of models without the spontaneously symmetry breaking can be described as the classical non-chaotic scattering process and the methods of symbolic dynamic are also used in this case. We show that the phantom cosmology can be treated as a simple model with scattering of trajectories which character depends crucially on a sign of a square of mass. We demonstrate that there is a possibility of chaotic behavior in the flat Universe with a conformally coupled phantom field in the system considered on non-zero energy level. We obtain that the acceleration is a generic feature in the considered model without the spontaneously symmetry breaking. We observe that the effective EOS coefficient oscillates and then approach to $w=-1$. ",
        "introduction": "The observations of distant supernovae \\cite{Perlmutter:1998np,Riess:1998cb} give the evidence that our Universe is undergoing accelerated expansion in the present epoch. In principle, there are two different approaches to treat this phenomenon. In the first approach it is postulated that there is some unknown exotic matter which violates the strong energy condition $\\rho + 3p \\geq 0$, where $p$ is the pressure and $\\rho$ is the energy density of perfect fluid. This form of matter is called dark energy. In the past few years different scalar field models like quintessence and more recently the tachyonic scalar field have been conjectured for modelling the dark energy in terms of sub-negative pressure $p > - \\rho$. A scalar field with super-negative pressure $p < - \\rho$ called a phantom field can formally be obtained by switching the sign of the kinetic energy in the Lagrangian for a standard scalar field. For example in the Friedmann-Robertson-Walker (FRW) model the phantom field minimally coupled to a gravity field leads to $\\rho + p = - {\\dot{\\psi}}^{2}$, where $\\rho_{\\psi} = -\\frac{1}{2} \\dot{\\psi}^{2} + V(\\psi)$, $p_{\\psi} = -\\frac{1}{2} \\dot{\\psi}^{2} - V(\\psi)$, and $V(\\psi)$ is the phantom potential. Such a field was called the phantom field by Caldwell \\cite{Caldwell:1999ew} who proposed it as a possible explanation of the observed acceleration of the current Universe when $\\Omega_{\\rm{m},0} \\gtrsim 0.2$. Note that a coupling to gravity in the quintessence models was also explored \\cite{Uzan:1999ch}.  The second approach called the Cardassian expansion scenario has recently been proposed by Freese and Lewis \\cite{Freese:2002sq} as an alternative to dark energy in order to explain the current accelerated expansion of the Universe. In this scenario the Universe is flat and matter dominated but the standard FRW dynamics is modified by the presence of an additional term $\\rho^{n}$ such that $3H^{2} = \\rho_{\\rm{eff}} = \\rho + 3B \\rho^{n}$, where $H = (d\\ln{a})/dt$ is the Hubble parameter; and $a$ is the scale factor. However, let us note that this additional term can be interpreted as a phantom field modelled by the equation of state $p = p(\\rho) = [n(1+ \\gamma) - 1] \\rho$, where $\\rho = \\rho_{\\rm{m},0} a^{-3(1+ \\gamma)}$. Therefore for dust matter we obtain $p = (n-1) \\rho$, and $n<0$ leads to the phantom field. The Cardassian expansion with $n<0$ which can be interpreted as the phantom fluid effect.  Phantom fields lead to the super-accelerated expansion of the Universe, i.e. $\\rmd H/\\rmd t > 0$. The simplest models describe this field in terms of minimally coupled real scalar field with the negative kinetic energy term $-\\frac{1}{2} \\dot{\\psi}^2$ \\cite{Caldwell:1999ew,Caldwell:2003vq,Carroll:2003st}. It is interesting that phantom fields are also present in string theories \\cite{Mersini-Houghton:2001su,Bastero-Gil:2001nu,Frampton:2002tu}. and arise as a phenomenological description of quantum effects of particle production in terms of bulk viscosity \\cite{Barrow:1988yc}. Because the phantom fields violate the Lorentz invariance condition most physicists believe that the phantoms open the doors on new physics \\cite{Visser:2003yf}. The investigation of the theoretical possibility to describe dark energy in terms of phantom field was the subject of many papers \\cite{Parker:1999td,Boisseau:2000pr,Schulz:2001yx,Faraoni:2001tq,Onemli:2002hr, Hannestad:2002ur,Melchiorri:2002ux,Hao:2003ww,Lima:2003dd,Singh:2003vx, Dabrowski:2003jm,Hao:2003th,Johri:2003rh,Piao:2003ty,Alam:2003fg,Nojiri:2003vn, Elizalde:2004mq,Sami:2003xv,Gannouji:2006jm}. The review on dark energy models were presented by Copeland et al.\\ \\cite{Copeland:2006wr}.  In the paper by Dabrowski et al.\\ \\cite{Dabrowski:2006dd} it is considered the quantization of phantom field via the Wheeler-DeWitt equation in quantum cosmology. They showed that quantum effects give rise to avoiding the big-rip singularity. Also some other basic problems in cosmology like the problem of direction of time can be solved.  If we confront the phantom field model with the observation of SNIa data we obtain that it is the best candidate together with the $\\Lambda$CDM model for the substantial form of dark energy \\cite{Choudhury:2003tj,Godlowski:2005tw, Szydlowski:2006pz,Szydlowski:2006ay}.  It is interesting to investigate new dynamical effects created in the presence of phantom fields. The main goal of this paper is to model the phantom fields in terms of scalar fields with a potential rather than in terms of the barotropic equation of state. For the latter case the dynamics is regular and can be represented on the two-dimensional phase space \\cite{Szydlowski:2004jv}.  Usually phantom fields are modelled in terms of minimally coupled scalar fields The case of minimally coupled to gravity phantom fields was analyzed in the context of existence of periodic solutions \\cite{Giacomini:2006ak}. They demonstrated that the dynamics is trivial in a sense of nonexistence of periodic solutions. Faraoni used the framework of phase space for investigating the dynamics of phantom cosmology, late-time attractors, and their existence for different shapes of potential \\cite{Faraoni:2005gg}. It was showed that dynamics of the flat FRW universe can be reduced to the form of a two-dimensional dynamical system on a double sheeted phase space. It is a simple consequence of an algebraic equation for $\\dot{\\psi}$ which can be expressed as a function of the Hubble parameter $H$ and $\\psi$. In this case there is no place for chaotic behavior of trajectories in the phase space $(H,\\psi,\\dot{\\psi})$ \\cite{Faraoni:2006sr}. The exit on the inflationary epoch and bounce was showed in the flat FRW universe with two interacting phantom scalar fields \\cite{Dzhunushaliev:2006xh}. In the closed FRW cosmological model with minimally coupled scalar field there appears transient chaos which has character of the scattering process \\cite{Toporensky:2005uk}. The scattering process takes place during the bounce --- a transition from a contracting to expanding universe. In this context the language of symbolic dynamics and topological entropy was used \\cite{Cornish:1997ah,Page:1984qt,Kamenshchik:1998ix}.  The conformally coupled scalar field are also of interest. The significance of long-wavelength modes in the WKB approximation of a conformally coupled scalar field was analyzed in the inflationary scenario \\cite{Jankiewicz:2005tm}.  While in this paper we concentrate on the conformal coupled scalar fields, it is worth to mention the works about the case of the generic (non-minimal and non-conformal) coupling between a phantom field and gravitation. Faraoni \\cite{Faraoni:2000gx} discusses different motivations for choosing $\\xi \\ne 0$ (both positive and negative) from the theoretical point of view and from observational constraints. The inflation and quintessence with the non-minimal coupling were studied in the context of the formulation of some necessary conditions for acceleration of the universe \\cite{Faraoni:2000wk,Faraoni:2000nt} (see also \\cite{Bellini:2002zr}). Faraoni also pointed out that the non-minimally coupling term different from the conformal coupling value ($\\xi = 1/6$) spoils the equivalence principle of general relativity.  It is interesting that some observational constraints on the non-minimal coupling (possitive or negative) can be found from CMB observations. Tsujikawa and Gumjjudpai \\cite{Tsujikawa:2004my} studied this bounds for a special power-law potential. The dynamics with the exponential form of the potential of a scalar field was also studied \\cite{Tsujikawa:2000tm}. The two-field inflation models with negative non-minimal coupling and hybrid inflation are also interesting in the context of the large scale curvature perturbation \\cite{Tsujikawa:2002nf}. In turn the constraints on the ratio of the self-coupling and non-minimal coupling constant can be obtained during the inflation period \\cite{Noh:2000kr}.  In this paper we ask what kind of dynamics can be expected from the FRW model with phantom field. It is well known that the standard FRW model reveals some complex dynamics. The detailed studies gave us a deeper understanding of dynamical complexity and chaos in cosmological models and resulted in conclusion that complex behavior depends on the choice of a time parameterization or a lapse function in general relativity \\cite{Misner:1969hg,Belinsky:1970ew,Cornish:1995mf}. Castagnino et al. \\cite{Castagnino:1999wd} showed that dynamics of the closed FRW models with a conformally coupled massive scalar field is not chaotic if considered in the cosmological time. The same model was analyzed in the conformal time by Calzetta and Hasi \\cite{Calzetta:1992bv} who presented the existence of chaotic behavior of trajectories in the phase space. Motter and Letelier \\cite{Motter:2002jj} explained that this contradiction in the results is obtained because the system under consideration is non-integrable. Therefore we can speak about complex dynamics in terms of nonintegrability rather than deterministic chaos. The significant feature is that nonintegrability is an invariant evidence of dynamical complexity in general relativity and cosmology \\cite{Maciejewski:2000}.  We can find many analogies between models with spontaneously symmetry breaking and the Yang-Mills systems. Problems of chaotic behavior in the Yang-Mills cosmological models have been investigated by Barrow and his collaborators. Chaos from Yang-Mills can only occur in an anisotropic universe but it occurs for some arbitrarily small anisotropy (see the study of Bianchi I \\cite{Barrow:1997sb} and other Bianchi types \\cite{Jin:2004vh}). However, when a perfect fluid is added things change in an interesting way that contrast with the situation with a magnetic field \\cite{Barrow:2005df}.  For the FRW model with phantom it can be shown that there is a monotonous function along its trajectories and it is not possible to obtain the Lyapunov exponents or construct the Poincar{\\'e} sections. Therefore it is useful to study the nonintegrability of the phantom system and set it in a much stronger form by proving that the system does not possess any additional and independent of Hamiltonian first integrals, which are in the form of analytic or meromorphic functions. Of course, it is not the evidence of sensitive dependence of solution on a small change of initial conditions. However, it is the possible evidence of complexity of dynamical behavior formulated in an invariant way \\cite{Szydlowski:2004jv}.  The notion of deterministic chaos is a controversial issue in the Mixmaster models. The value of the numerically computed maximal Lyapunov exponent for those systems depends on the time parameterization used \\cite{Rugh:1990aq}. However, the existence of fractal structures in the phase space provides the coordinate independent signal of chaos in cosmology as it was shown by Cornish and Levin \\cite{Cornish:1995mf}. In particular, they found that the Bianchi IX has a form of chaotic scattering. The short scattering periods intermittent integrable motion and evolution of the system is chaotic. It is similar to a pin-ball machine.  The main goal of this paper is to show that dynamics of phantom cosmology can be treated as a scattering process. Only if the spontaneously symmetry breaking is admitted this process has chaotic character. Evidences of dynamical behavior are studied in tools of symbolic dynamics, fractal dimension, and analytically the Toda-Brumer-Duff test. ",
        "conclusions": "In cosmology the Mixmaster models are well known for their chaotic behavior \\cite{Cornish:1995mf}. In these models chaos is caused by anisotropy of space. In this paper we considered the phantom cosmological models which are homogeneous and isotropic where the chaos is the consequence of the nonlinearity of the potential function for the scalar field. We found chaos is a generic feature of the phantom cosmology with the spontaneously symmetry breaking. It assumes the form of the multiple chaotic scattering process. In the case of the absence of symmetry breaking we did not find the sensitivity over initial conditions and the system behaves as the standard scattering process.  There is no difference between the flat universes filled with minimal and conformally coupled scalar fields. As it was demonstrated by Faraoni et al. \\cite{Gunzig:2000ce,Faraoni:2006sr} there is no chaos for the former class of models because the solution corresponding $k=0$ are restricted to move in some two-dimensional submanifold of the phase space $(H, \\phi, \\dot{\\phi})$. There are no enough room for chaotic motion manifestation. The analogous result is valid for the other value of the coupling constant $\\xi$. Note that we can observe chaotic behavior in some cases of the conformally coupled scalar field in the flat universes, when we consider the system on non-zero energy level. Physically it means that the universe is filled with radiation matter.  The reason for which the flat universes with scalar field are non-chaotic lies in lack of invariant measure of chaotic behavior in the general relativity \\cite{Szydlowski:1996uf,Motter:2003jm}. In general the standard chaos indicators like Lyapunov exponents or Kolmogorov-Sinai entropy depends on time parameterization. In general relativity there is freedom of choice of the lapse function which defines reparameterization of time. The Hamiltonian is given modulo the lapse function. Therefore if we consider the Hamiltonian FRW system with scalar field on zero energy level it is possible that in the Hamiltonian constraint $a$ i $\\dot{a}$ appear only in the combination $H=\\dot{a}/a$. Hence the motion of the system takes place in the 3-dimensional phase space $h(H,\\phi,\\dot{\\phi}) = 0$ on some 2-dimensional submanifold. For any value of coupling constant and general class of potentials there is no place for chaos (even the models with the cosmological constant).  The new picture appears if the system is described on a non-zero energy level. Then the motion of the system is in the 3-dimensional submanifold of the 4-dimensional phase space. But in general case there is no possibility of choosing the lapse function in such a way that $a$ i $\\dot{a}$ contribute in the Hamiltonian constraint by the Hubble function.  Among the conformally coupled models there are in principal two types of behavior. First, chaotic scattering process takes place for case $m^2 < 0$ (the spontaneously symmetry breaking). In this case trajectories possess the property of topological transitivity which guarantees their recurrence in the phase space and the scattering process takes place around the origin. It is similar to the chaos appeared in the Yang-Mills theory. Exploring this analogy to this system we calculate a Gaussian curvature of the potential function which measure local instability of nearby trajectories. We proof that this curvature is negative. Because the configuration space is bounded in this case the mixing of trajectories is observed.  Second, in the opposite case of $m^2 > 0$ we characterize dynamics in the system in terms of symbolic dynamics. No fractal structure in the space of initial conditions was found and conclude that the scattering process has not the chaotic character.  We also found that the universe filled with conformally coupled scalar field is accelerating without the positive cosmological constant. In principle, there are two sources of this acceleration. The acceleration is driven by negative energy density and coefficient of equation of state $w_{\\phi}$ is greater than $-1/3$. Alternatively, energy density is positive and the coefficient of equation of state can be smaller than $-1/3$. The latter there is crossing of $w_{\\phi} = -1$. Of course, for the explanation of SNIa data apart from the acceleration of the Universe itself, the rate of acceleration is also crucial.  The system with phantom scalar field coupled conformally to gravity in the FRW universe can be treated as scattering process with chaotic and nonchaotic character. The generic feature of this class of systems with the spontaneously symmetry breaking is the chaotic scattering of trajectories."
    },
    "0608/astro-ph0608242_arXiv.txt": {
        "abstract": "We present results for the spectral distortions of the Cosmic Microwave Background (CMB) arising due to bound-bound transitions during the epoch of cosmological hydrogen recombination at frequencies down to $\\nu\\sim 100\\,$MHz. We extend our previous treatment of the recombination problem \\citep{Jose2006} now including the main collisional processes and following the evolution of all the hydrogen \\change{angular momentum} sub-states for up to 100 shells. We show that, due to the low baryon density of the Universe, even within the highest considered shell full statistical equilibrium (SE) is not reached and that at low frequencies the recombination spectrum is significantly different when assuming full SE for $n>2$. We also directly compare our results for the ionization history to the output of the {\\sc Recfast} code, showing that especially at low redshifts rather big differences arise. In the vicinity of the Thomson visibility function the electron fraction differs by roughly $-0.6\\%$ \\change{which affects the temperature and polarization power spectra by $\\lesssim 1\\%$}. Furthermore we shortly discuss the influence of free-free absorption and line broadening due to electron scattering \\change{on the bound-bound recombination spectrum} and the generation of CMB angular fluctuations due to scattering of photons within the high shells. ",
        "introduction": "The study of the Cosmic Microwave Background (CMB) provides one of the most powerful tools to test cosmological models and to determine the parameters describing the Universe. Nowadays, most of the CMB experiments are focused on the study of the temperature and polarization angular fluctuations, which have been measured recently with very high precision by the {\\sc Wmap} satellite \\citep{Page2006, Hinshaw2006}, permitting to achieve accuracies in the determination of some of the key cosmological parameters on the level of a few percent.  The CMB photons which are detected today were essentially last scattered towards us during the epoch of cosmological recombination, where the temperature of the Universe had become sufficiently low to permit the formation of neutral atoms \\citep{Zeldovich68,Peebles68}. During the process of cosmological recombination of the hydrogen and helium atoms there was a net generation of photons, which introduces distortions to the CMB blackbody spectrum. Except for the far Wien tail of the CMB, these distortions are expected to be very small: the distortions due to bound-bound transitions reach the level of $\\Delta I/I\\sim 10^{-7}$ at frequency $\\nu\\sim 1\\,$GHz and are smaller than $\\Delta I/I\\sim 10^{-8}$ in the range $10\\,\\text{GHz}\\leq\\nu \\leq 500\\,$GHz \\citep[e.g. see][RMCS06 hereafter]{Jose2006}. This is mainly due to the extremely large entropy of the Universe (there are $\\sim \\pot{1.6}{9}$ photons per baryon), since in the presence of undistorted blackbody radiation there cannot be much more than one photon per transition \\citep{Liubarskii83}. The free-bound emission increases the spectral distortions by additional $\\sim 30-90$\\% \\change{and a total $\\sim5$ photons per hydrogen atom are released due to hydrogen recombination in the undistorted ambient blackbody photon field of the CMB} \\citep{Chluba2006b}. However, a measurement of these tiny deviations from a pure blackbody would provide another independent way to determine some of the key cosmological parameters, such as the baryon density or the total matter content. In addition they in principle are a {\\it byproduct} of the accurate treatment of the recombination history and in particular the time dependence of the free electron fraction in the Universe, which today is so important for the theoretical prediction of the CMB temperature and polarization power spectra.    The numerical solution of the hydrogen recombination history and the associated spectral distortions of the CMB requires the integration of a stiff system of coupled ordinary differential equations, describing the evolution of the populations of the different hydrogen levels, with extremely high accuracy, simply because the amplitude of the distortions is so small. Several authors have computed the distortions due to bound-bound transitions for a multi-level hydrogen atom with different assumptions and simplifications \\citep[][RMCS06]{Dubrovich1975,Beigman1978,Rybicki93,DubroVlad95,Dubrovich2004,Kholu2005,Wong2006}. Among these the most important simplification is to assume {\\it full statistical equilibrium}\\footnote{i.e. the population of a given level $(n,l)$ is determined by $N_{nl} = (2l+1) N_n / n^2$, where $N_n$ is the total population of the shell with principle quantum number $n$.} (SE) within a given shell for $n>2$ (for a more detailed comparison of the different approached see RMCS06 and references therein). With this assumption the recombination problem becomes significantly simpler, since for each additional shell only one more equation has to be treated, and \\citet{Kholu2005} already presented results for the hydrogen recombination spectrum due bound-bound transitions in the frequency range $1\\,\\text{GHz}\\leq\\nu\\leq 1000\\,\\text{GHz}$ including up to 160 shells \\change{within} this approximation.  In our previous work (RMCS06) the spectral distortions of the CMB spectrum arising due to bound-bound transitions and the 2s two-photon decay during the epoch of cosmological hydrogen recombination were obtained in the frequency range $1\\,\\text{GHz}\\leq\\nu\\leq 3500\\,\\text{GHz}$. There the complete set of ordinary differential equations was solved for the cosmological hydrogen recombination problem including up to 30 shells, with {\\it all} the energetically degenerate angular momentum sub-levels treated separately. It was shown that assuming SE for shells $n>2$ leads to significant differences in the cosmological hydrogen recombination spectrum for the Balmer, Paschen and Brackett series. Furthermore, it was argued that collisions probably can only become important for shells well above $n\\sim 25-30$, so that assuming SE for the lower shells is not justified, but a detailed examination of this problem has not been performed until now.  In this paper, we present detailed results for the hydrogen recombination spectrum due to bound-bound transition at low frequencies down to $\\nu\\sim 100\\,$MHz (see Fig.~\\ref{fig:DI_results}). To this end, we have solved the complete set of equations for a multi-level hydrogen atom with 100 shells, i.e. we are following the populations of 5050 separate hydrogen states. Furthermore, we include the main collisional processes into our calculations (for more details see Sect.~\\ref{sec:coll_details}) and discuss deviations from full SE within the high shells. We also examine the influence of the free-free process (Sect.~\\ref{sec:freefree}) at low frequencies, the line broadening due to electron scattering (Sect.~\\ref{sec:el_sc}) and the generation of CMB angular fluctuations by scattering of photons within the high shells (Sect.~\\ref{sec:more}). For the treatment of the multi-level hydrogen atom we basically follow the same procedure as described in RMCS06, with some modification which are explained in Sect.~2.  Several authors have worked on percent level corrections to the ionization history arising from different previously neglected physical processes \\citep{Dubrovich2005, Chluba2006, Kholu2006, Novosyadlyj2006}. It was pointed out that these corrections lead to differences in the theoretical predictions for the CMB temperature and polarization power spectra on a similar level and therefore should be taken into account when analyzing future high precision data. Recently, \\citet{Lewis2006} have made some first steps towards quantifying the possible impact of percent level correction to the ionization history on the estimation of cosmological parameters, again showing that in the era of precision cosmology more efforts should be made to improve the calculation of the recombination history.  In RMCS06 it was also pointed out that there are corrections to the ionization history on the level of percent arising due to the detailed treatment of the angular momentum sub-states of the hydrogen atom. With computations including up to 100 shells, one reaches the point to directly compare the obtained ionization history with the output of the {\\sc Recfast} code, which was developed to represents the computations based on a 300-level hydrogen atom \\citep{SeagerRecfast1999}\\change{, but under the assumption of SE for $n>2$.} We shall discuss our results for the obtained electron recombination history as compared to the {\\sc Recfast} output in Sect.~\\ref{sec:el_hist} and show that especially at low redshifts rather big differences arise. However, still more work has to be done to reach percent level accuracy \\change{for the ionization history}.  \\change{For all the computation presented in this work the following values of the cosmological parameters were adopted: $\\Omega_{\\rm b}=0.0444$, $\\Omega_{\\rm tot}=1$, $\\Omega_{\\rm m}=0.2678$, $\\Omega_{\\Lambda}=0.7322$, $Y_{\\rm p}=0.24$, $h=0.71$ and $T_0=2.725\\,$K. } ",
        "conclusions": "We presented results for the spectral distortions of the CMB arising during the epoch of cosmological hydrogen recombination due to bound-bound transition with observing frequencies down to $\\nu\\sim 100\\,$MHz and confirm that the relative spectral distortions, $\\Delta I_\\nu/B_\\nu$, are rapidly growing towards low frequencies. We show that even within the highest considered shell full statistical equilibrium is not reached and that at low frequencies the recombination spectrum is significantly different when full SE for shells with $n>2$ is assumed. We also conclude that in cosmological hydrogen recombination calculations $l$-changing collisions should be included for shells with $n\\gtrsim 30-40$, but $n$-changing collisions can be neglected for shells with $n\\lesssim 100$.  In Sect.~\\ref{sec:el_hist} we have presented our results for the ionization fraction as compared to the {\\sc Recfast} output. \\change{Including a treatment of the populations of all the individual angular momentum sub-states slows recombination down, leaving a $\\gtrsim 5-10\\%$ larger residual electron fraction at low redshifts. Although the overall effect on the temperature and polarization power spectra due to the changes in the recombination history is $\\lesssim 1\\%$ these additional electrons may be important in computations of the first stars and structures \\citep{Barkana2001, Bromm2002}, for which especially the formation of molecular hydrogen plays a crucial role.} The discussion has shown that representing the full multi-level calculation with percent level accuracy in the full range of redshifts probably requires some more general 'fudge-function' instead of a single 'fudge-factor'. Also, one still has to push to a larger number of shells and very likely include more physics. For extensive calculations of the CMB temperature and polarization power spectra it will become important to develop a new simplified and fast scheme which accurately models the physics during the recombination epoch."
    },
    "0608/astro-ph0608074_arXiv.txt": {
        "abstract": "The Pierre Auger Observatory is a hybrid detector for ultrahigh energy cosmic rays (UHECR) with energies above 10$^{18.5}$~eV. Currently the first part of the Observatory nears completion in the southern hemisphere in Argentina. One detection technique uses over 1600 water Cherenkov tanks at ground where samples of secondary particles of extensive air showers (EAS) are detected. The second technique is a calorimetric measurement of the energy deposited by EAS in the atmosphere. Charged secondary particles of EAS lose part of their energy in the atmosphere via ionization. The deposited energy is converted into excitation of molecules of the air and afterwards partly emitted as fluorescence light mainly from nitrogen in the wavelength region between 300 and 400~nm. This light is observed with 24 fluorescence telescopes in 4 stations placed at the boundary of the surface array. This setup provides a combined measurement of the longitudinal shower development and the lateral particle distribution at ground of the same event. Details on the fluorescence technique and the necessary atmospheric monitoring will be presented, as well as first physics results on UHECR. ",
        "introduction": "Ultrahigh energy cosmic rays (UHECR) are a continuous particle flux arriving at the Earth's atmosphere with strongly decreasing rate at highest energies. At energies above 10$^{18}$~eV, their origin is expected to be extragalactic. However, even after several years of measurements, the major questions remain unanswered: where and how do these particles obtain such high energies and what is their elemental composition. For resolving the enigma about UHECR, experiments are designed for measuring the primary energy and type of cosmic rays as well as their arrival directions.  Entering the Earth's atmosphere, cosmic rays initiate extensive air showers (EAS). The primary particle of an EAS, which is the actual cosmic ray particle, interacts with a nucleus of air molecules. This hadronic interaction yields new secondary particles which will interact themselves with nuclei of air molecules. A growing cascade of secondary particles develops in the Earth's atmosphere, which is called an extensive air shower. After an initial phase of particle multiplication, the number of charged particles in the shower falls due to energy losses. For inclined showers with moderate primary energy, only some of the overwhelming amount of particles reaches the Earth's surface. These particles can be observed by particle detectors at ground. For UHECR with nearly vertical incidence, the maximum of shower development is reached close to ground or even ``below'' ground.  Another detection technique arises from the energy losses of EAS in the atmosphere. The energy deposited in the air excites molecules of the atmosphere where particularly the excitation of nitrogen molecules is of interest. The de-excitation of nitrogen molecules is partly proceeding through the channel of fluorescence emission. So a light track of the air shower in the UV-wavelength region between 300 and 400 nm can be detected by adequate telescopes. This enables a calorimetric measurement of the deposited energy of EAS which can be used to obtain the primary energy of the cosmic ray. ",
        "conclusions": ""
    },
    "0608/astro-ph0608597_arXiv.txt": {
        "abstract": "Since the discovery of short-period exoplanets a decade ago, photometric surveys have been recognized as a feasible method to detect transiting hot Jupiters. Many transit surveys are now under way, with instruments ranging from $10$-cm cameras to the Hubble Space Telescope. However, the results of these surveys have been much below the expected capacity, estimated in the dozens of detections per year.  One of the reasons is the presence of systematics (``red noise'') in photometric time series. In general, yield predictions assume uncorrelated noise (``white noise''). In this paper, we show that the effect of red noise on the detection threshold and the expected yields cannot be neglected in typical ground-based surveys. We develop a simple method to determine the effect of red noise on photometric planetary transit detections. This method can be applied to determine detection thresholds for transit surveys. We show that the detection threshold in the presence of systematics can be much higher than with the assumption of white noise, and obeys a different dependence on magnitude, orbital period and the parameters of the survey. Our method can also be used to estimate the significance level of a planetary transit candidate (to select promising candidates for spectroscopic follow-up).    We apply our method to the OGLE planetary transit search, and show that it provides a reliable description of the actual detection threshold with real correlated noise.  We point out in what way the presence of red noise could be at least partly responsible for the dearth of transiting planet detections from existing surveys, and examine some possible adaptations in survey planning and strategy. Finally, we estimate the photometric stability necessary to the detection of transiting ``hot Neptunes''. ",
        "introduction": "Many photometric surveys for transiting exoplanets are now under way, with a wide variety of instrumentation and observation strategies -- from monitoring large areas of the sky with small telescopes to deeper surveys on star clusters or the Galactic disc with $1$--$4$m telescopes\\footnote{see e.g. \\citet{alo04,bak04,bra06,bru03,hid05,hoo05,mal03,rau04, uda02a} for a description of some planet transit surveys; \\citet{cha06} for a recent review of the results.}. Altogether, thousands of telescope nights have been invested in these surveys, monitoring hundreds of thousands of target stars in the Solar neighbourhood and in the Galactic disc.  However, even after years of operation, the results of these surveys failed to meet the expectations, with only a slow trickle of detections instead of the expected bounty.  Part of the mismatch between expectations and actual performance can be attributed to the fact that these surveys were often assuming that the first known transiting extrasolar planet and for long the only one, HD209458$b$ \\citep{cha00}, was typical of hot Jupiters. Its radius of $\\sim1.3$~R$_J$ creates a $2\\%$ transit signal on a solar-type star. Subsequently discovered transiting gas giants showed the large radius of HD209458 to be an exception rather than the rule \\citep{alo04,pon04}. The mode of hot Jupiter radii is probably near $1.1$~R$_J$ or lower \\citep{gau05}, producing correspondingly shallower eclipses.  Still, even assuming smaller gas giants, transit surveys keep predicting many more detections than they actually yield. Understanding this mismatch is essential both for interpreting the results in terms of planetary statistical properties, and in order to improve the planning and strategy of the surveys. Indeed, if the statistical properties of hot Jupiters were not known from radial velocity surveys, it is very likely that drastically different conclusions would have been drawn from the results of the transit surveys. Upper limits would have been put to the abundance of hot Jupiters one or two orders of magnitude below those derived from Doppler surveys.     The ingredients in the prediction simulations seem solid: the abundance of hot Jupiters is relatively well known from radial velocity surveys, the statistics of targets are obtained from well-tested models of Galactic stellar populations, and the rest comes from simple orbital mechanics. The setting of the transit detection threshold is often considered a minor component in the simulations. The threshold is generally modelled as a minimum signal-to-noise ratio (``S/N'') of the transit detections assuming uncorrelated noise in the photometric data.  In this paper we show that, contrary to these assumptions, the correlation of photometric data at the millimagnitude level cannot be neglected when determining the detectability of planetary transits, and that taking the correlation into account can strongly affect the detection threshold, and consequently the estimates of the potential of photometric surveys in terms of planet detetion.   We develop a simple method to assess the significance of detected transit candidates in the presence of such ``red noise''. We propose a method that is robust and gives realistic results, while at the same time remains simple to use and to apply to any ground-based transit survey.        Ongoing transit surveys have shown that detecting transiting signals was not the end of the story. For transit depths in the range of a few percents, by far the largest number of detections are due to eclipsing binaries, either small transiting M dwarfs \\citep[e.g. ][]{pon05} or eclipsing binaries diluted by the light of an unresolved companion \\citep[e.g. ][]{man05}. The identification of true transiting planets among all identified transit candidates requires a considerable investment in spectroscopic follow-up observations. For transit depths below $\\sim 2\\%$, the odds become more favourable to transiting planets \\citep{bro03}, but since such transit depths are near the detection threshold of most ground-based surveys, false positive detections start being a source of contamination. False positives require even more follow-up observations than eclipsing binaries. Therefore the assignment of reliable significance levels to transit detections near the detection threshold, is another motivation for a robust method to assess the significance of transit candidates in the presence of systematics or ``red noise''.  A third motivation is to derive realistic estimates for the uncertainty on planetary parameters derived from transit lightcurves, taking into account the effect of systematics in the photometry.  In Section~\\ref{Sec2} we present our method to compute the significance level of a transit detection. In Section 3 we examine the implications for transit surveys. Section 4 summarizes the results of this paper and points out some interesting consequences. ",
        "conclusions": "This study shows that in ground-based surveys for planetary transits, the systematics in the photometry have a very important effect on the capacity to detect planetary transits, and that estimates of detection significance and detection threshold based on the assumption of white Gaussian noise are not appropriate. We propose a simple formalism to assess the significance level of detected transits, and to predict the detection threshold of an observing campaign.  We find that the uncertainty on the depth $d$ of a transit candidate is \\begin{equation} \\sigma_d^2= \\frac{1}{n^2} \\sum_{k=1}^{N_{tr}} n_k^2 \\V (n_k) \\ , \\end{equation} where $N_{tr}$ is the number of transits sampled, $n_k$ the number of data points in the $k^{\\mathrm th}$ transit. $\\V$ is a function that can be calculated through the curve itself from the variance of the average flux over transit-length intervals outside the transit.  Consequently the significance (signal-to-noise) of a transit signal candidate in the presence of real photometric noise can be calculated by: \\begin{equation} \\Sr^2= d^2 \\frac{n^2}{\\sum_{k=1}^{N_{tr}} n_k^2 \\V (n_k)} \\ . \\end{equation}    Introducing the $\\sigma_r$ factor to parametrize the red noise, the threshold for detection is: \\[ \\frac{d^2 n^2}{ \\sum_{k=1}^{N_{tr}} n_k^2 \\left( \\frac{\\sigma_w^2}{n_k} + \\sigma_r^2\\right)} > S_T^2 \\ , \\] where $\\sigma_w$ is the uncorrelated (white) noise in the photometry. Typical values of $\\sigma_r$ are of a few millimagnitudes (3 mmag in the OGLE survey), and typical values of $S_T$ are $9$ or higher in realistic settings.  In terms of physical parameters the threshold is: \\[ \\Sr^2  = \\alpha (r/R)^4 \\left[ c \\pi \\frac{\\sigma_w^2 P^{2/3} M^{1/3}}{\\beta N R} + \\frac{\\sigma_r^2 P }{N \\delta_t}\\right]^{-1} > S_T^2 \\ . \\]    Our approach bears similarities with that of \\citet{jen02}, \\citet{kov02} and \\citet{sir03}, in accounting for the systematics in the noise by studying the properties of the photometric time series over longer time intervals. We tried to devise a method that is easy to apply, does not require long computation, and can be extrapolated to ongoing or future surveys with minimal extra assumptions.   In general, accounting for the effect of systematics results in much lower yield estimates for ground-based transit surveys. We believe that these estimates are more realistic than white-noise estimates. They are also in much better agreement with the low actual detection rate. These more realistic estimates can help in the design of transit surveys. Including the red noise modifies the dependence of the yield estimates on the survey parameters. Thus, survey design can be affected not only quantitatively by lowering the predicted yields, but also qualitatively.  The presence of systematics has several important effects: (1) it reduces the premium placed on the brightest end of a survey for detection, (2) it steepens the period dependence of detectability and (3) it reduces the importance of denser time sampling.  These elements may for instance weigh a survey strategy towards including more fields with a longer separation between individual measurements, and spending more time on a given field.   In the case of surveys targeting stellar systems, red noise modifies the conclusions of \\citet{gau05} that if the detection threshold reaches a hot Jupiter transit in front of a target of any magnitude, such transits will be detected for targets of all magnitudes. Red noise will hamper the detection for the brightest magnitudes.  Our study also underlines the crucial importance of beating down the hour-timescale systematics in the photometry of transit surveys. This, of course, is well realized by all groups conducting such surveys, and considerable effort has been invested in removing the systematics. For instance, for the OGLE survey, two schemes have been developed to remove the systematics by using the correlation among all the lightcurves \\citep{kru03,tam05}.  However, for a given facility there may be a lower limit below which it will be practically impossible to reduce $\\sigma_r$. Since transit surveys need as large a time coverage as they can get, they are compelled to use all clear nights whether perfectly photometric or not. At the millimagnitude level, interactions between airmass, colour, absorption, tracking, flatfielding and seeing all have the potential to introduce systematic trends. For instance, systematic residuals can be caused by the interaction of seeing with a close unseen companion, or by the tracking drift moving part of an object across pixels with flatfield errors.  Sub-millimagnitude accuracy has been reported in some occasions, for instance by \\citet{har05}. This can be attained with a very good sampling of the PSF on a large number of pixels, as possible with a large telescope on small fields. However, the reported results concern only a few nights, probably of above-average quality. It remains to be seen if such accuracy can be maintained over a period of time long enough  for an efficient planetary transit survey. Moreover, the use of a large telescope on a small field implies that the detected transit candidates will be too faint for spectroscopic confirmation, limiting their usefulness. As shown by the existing surveys, most transit candidates are actually eclipsing binaries, so that spectroscopic follow-up is indispensable in order to confirm the planetary nature of the transiting companion.      An obvious solution to the problem of photometric systematics is to move into space, as was indeed done soon after the discovery of 51\\,Peg with {\\it HST} in the globular cluster 47\\,Tuc \\citep{gil00}. More recently, Sahu et al. (programme GO-9750) have carried out an HST survey for transits in the Galactic plane. The problem, however, is that the detected transit candidates are again too faint for spectroscopic follow-up, so that planets cannot be told apart from stellar eclipses.  In the coming years, two space-born planetary transit missions are scheduled, {\\it CoRoT} and {\\it Kepler}. They will be free from systematics caused by the atmosphere. The approach of this paper, however, also applies to space-born data. The transit detection threshold for {\\it CoRoT} and {\\it Kepler} will also be set by the characteristics of the red noise, both instrumental and due to stellar variability, because transit-length intervals will be sufficient for the white noise to average out.  The most relevant factor will be the stability of the hour-timescale average of the measured fluxes.  In summary, this work offers a practical way, both conceptually and computationally, to empirically deal with the complex problem of covariant noise in extrasolar planetary transit surveys, both ground-based and space-born."
    },
    "0608/astro-ph0608304_arXiv.txt": {
        "abstract": "On the basis of a large scale 'adiabatic', namely  non-radiative and non-dissipative, cosmological smooth particle hydrodynamic simulation we compare the entropy profiles of the gas and the dark matter (DM) in galaxy clusters. We employ the quantity $K_{\\rm g}= 3k_BT_{\\rm g}\\rho{_{\\rm g}^{-2/3}}/(\\mu m_p) = \\sigma^2_{\\rm g}\\rho{_{\\rm g}^{-2/3}}$ as measure for the {\\it entropy} of the  intra-cluster gas. By analogy the DM entropy is defined as $K_{\\rm DM} = \\sigma{^2_{\\rm DM}}\\rho{_{\\rm DM}^{-2/3}}$ ($\\sigma{^2_{\\rm DM}}$ is the 3D velocity dispersion of the DM). The DM entropy is related to the DM phase space density by $K_{\\rm DM}\\propto Q{_{\\rm DM}^{-2/3}}$. In accord with other studies the radial DM phase space density profile follows a power law behaviour, $Q_{\\rm DM} \\propto r^{-1.82}$, which corresponds to $K_{\\rm DM}\\propto r^{1.21}$. The simulated intra-cluster gas  has a flat entropy core within $(0.8 \\pm 0.4) R_{\\rm s}$, where \\rs\\ is the NFW scale radius. The outer profile follows the DM behaviour, $K_{\\rm g}\\propto r^{1.21}$, in close agreement with X-ray observations. Upon scaling the DM and gas densities by their mean cosmological values we find that outside the entropy core a constant ratio of $K_{\\rm g}/K_{\\rm DM}=0.71\\pm0.18$ prevails. By extending the definition of the gas temperature to include also the bulk kinetic energy the ratio of the DM and gas extended entropy is  found to be unity for $r\\gtrsim 0.8 R_{\\rm s}$. The constant ratio of the gas thermal entropy to that of the DM implies that observations of the intra-cluster gas can provide an almost direct probe of the DM. ",
        "introduction": "Clusters of galaxies represent a composite of baryonic and dark matter (DM). The gas is commonly described by its thermodynamic variables, such as density, temperature, pressure and also entropy. The entropy provides a unique gauge for monitoring the thermal history of the gas. In the standard model of structure formation the primordial gas is isentropic. Ignoring radiative and dissipative processes and other non-gravitational sources or sinks of energy the only mechanism by which the gas entropy can change is {\\it via} shock waves. According to the paradigm of hierarchical structure formation clusters form by merging and accretion of smaller building blocks, which usually contain baryons and DM. The collapse and merging of substructures lead to emergence of shock waves in the forming intra-cluster gas (ICM) and it ceases to be isentropic. These shocks can be seen close to the virial radius when the gas is accreted onto the cluster (see e.g. ~\\citealt{Borgani05, Voit05, Voit05b}). However, the study of \\cite{Pfrommer06} shows that a non negligible amount of energy is dissipated by shocks further inside the cluster potential well. In any case the resulting entropy structure reflects the contribution of irreversible processes in the formation of clusters. The main problem addressed here is to what extent the gas entropy resembles the structure of the DM halo and whether the classical thermodynamic description of the gas can be extended and applied to the DM.  The entropy of the ICM can be deduced from X-ray observations. As has become customary in this field, we refer to $K\\propto T_{\\rm g}\\rho_{\\rm g}^{-2/3}$ as the 'entropy', where $T_{\\rm g}$ is temperature and $\\rho_{\\rm g}$ is the density of the gas. More precisely, we define $K_{\\rm g}\\equiv 3k_BT_{\\rm g}\\rho{_{\\rm g}^{-2/3}}/(\\mu m_p) = \\sigma^2_{\\rm g}\\rho{_{\\rm g}^{-2/3}}$, where $k_B$ is the Boltzmann constant and $\\sigma_{\\rm g}$ is the internal 3D velocity dispersion of the gas. It is the latter expression upon which the direct comparison between the gas and the dark matter is based. Using {\\it XMM-Newton} observations of 10 relaxed clusters \\cite{Pratt06} find the radial profile of $K$ follows a power law of $r^\\beta$ with $\\beta \\approx 1.1$ (see also \\citealt{Donahue06}). Apart from the central regions the slope obtained in adiabatic simulations agree with the observed values (e.g.~\\citealt{Voit05}). The observations, however, show an excess in entropy which is thought to originate from non gravitational processes (e.g. \\citealt{Finoguenov03, Borgani05}). Note, the actual entropy of an ideal gas is $S=\\ln(K^{3/2})+ const.$, but here we follow the common convention and associate the entropy with $K$.  The concept of entropy has also been applied to collisionless self-gravitating systems  in its Boltzmann-Gibbs formulation (e.g.~\\citealt{LyndenBell67,LyndenBell68}) and in an analogy with the thermodynamic entropy of a perfect gas \\citep{White87,Navarro95,Eke98}. It is the latter approach that we follow here. In  analogy with classical thermodynamics and the ideal gas equation of state we write  the local specific entropy $S_{\\rm DM}=\\ln(\\sigma{_{\\rm DM}^3} / \\rho_{\\rm DM})$, where $\\sigma_{\\rm DM}$ is the DM velocity dispersion and an additive constant is neglected. A similar approach is presented by \\cite{Eke98}.  Following the convention used in the studies of clusters of galaxies we define $K_{\\rm DM}=\\sigma{^2_{\\rm DM}}\\rho{^{-2/3}_{\\rm DM}}$. Note that the DM specific entropy  is the log of the inverse of the so-called phase space density $Q_{\\rm DM}=\\rho_{\\rm DM} \\sigma_{\\rm DM}^{-3}$. Now, the analysis of simulated DM halos by \\cite{Taylor01} has shown that the radial phase space density profile closely follows a power law, $Q_{\\rm DM}(r)\\propto r^{-1.875}$. Using a set of high-resolution N-body simulations \\cite{Ascasibar04} confirm this power-law behaviour for an extend range of haloe masses, in addition they find coincidence of the zero-points if the profiles are conveniently rescaled by virial quantities. Changing variables from $Q$ to $K$ the \\cite{Taylor01} power law is given by \\begin{equation} K_{\\rm DM}(r)=Q{_{\\rm DM}^{-2/3}}(r) \\propto r^{1.25}. \\end{equation} The close similarity of the (appropriately converted) \\cite{Taylor01} slope and the slope derived from X-ray observations by \\cite{Pratt06} has motivated us to compare the DM and gaseous entropy profiles of clusters of galaxies extracted from a large scale 'adiabatic' Smoothed Particle Hydrodynamics (SPH) simulation. In this context adiabatic is defined as non-radiative with no non-gravitational sources or sinks of energy. In such kind of simulations the entropy of each particle remains constant as long as it does not pass through a shock front \\citep{Springel05}. The only dissipative process is introduced by the numerical viscosity which damps unphysical oscillations behind shock fronts and thus transfers kinetic into thermal energy in accordance with the Rankine-Hugoniot jump conditions.  This paper is organised as follows. \\S\\ref{sec:sim} introduces the simulation and the cluster ensemble and ~\\S\\ref{sec:profiles} compares the entropy profiles of the gaseous and the DM component of the clusters. \\S \\ref{sec:core} focuses on the emerging gas entropy cores and discusses their physical properties.  And a discussion is given in ~\\S\\ref{sec:discu}. ",
        "conclusions": "\\label{sec:discu} The main findings of the paper are: The DM phase space density power law of \\cite{Taylor01} is reproduced by the simulated clusters extracted from an adiabatic gas dynamical simulation. The gas entropy profile is in accord with previous simulations \\citep{Voit05b}. In particular the  entropy profile has an inner flat core and an outer profile of $K_{\\rm g}\\propto r^{1.21}$ is recovered. This stands in at least qualitative agreement with X-ray observations. The entropy core coincides with the region of the temperature inversion of the DM and its radius  is  characterised by $R_{\\rm core} /  R_{\\rm s} = 0.8 \\pm 0.4$.  However, entropy loss due to radiative processes, which is not included in our simulations, alters the thermodynamics of the central ICM, see e.g. \\cite{Kravtsov05}. Therefore, it is no contradiction to our findings that \\cite{Piffaretti05} don't detect entropy cores within 13 nearby cooling flow clusters. The most important result of the paper is that outside the entropy core the gas and DM entropies follow one another very closely. Upon appropriate scaling of the gas and DM densities by their cosmological mean values we find that outside the entropy core a constant ration of $K_{\\rm g}/K_{\\rm DM}=0.71\\pm0.18$ prevails. By extending the definition of the gas entropy to include also the kinetic energy of the turbulent motion the above ratio becomes unity, and it follows that in the outer part of the cluster the gas and DM entropies coincide and are virtually the same.  The constant ratio of the gas to the DM entropies outside the entropy core has interesting observational ramifications. To the extent that non-adiabatic physics, in the sense described above, does not affect the gas structure outside the entropy core observations of the ICM can shed light on the DM in clusters. Observational determination of the entropy of the ICM can be easily translated into the entropy of the DM.  Given the power law nature of the DM entropy, the observationally determined outer DM entropy can be extrapolated inwards. This, together with independent mapping of the mass distribution can yield interesting insight into the distribution and dynamics of DM in clusters of galaxies. In this context it is worth mentioning that \\cite{Ikebe04} described a formalism how the DM ``temperature'' can be derived from X-ray observations and applied this new technique to {\\it XMM-Newton} data of A1795.   The analysis of \\S~\\ref{sec:core} shows that the entropy core radius, measured in \\rs\\ units, is independent of the cluster NFW parameters. This seems to imply that the emergence of the entropy core is related to processes occurring within \\rs. The common understanding of halo formation is that the structure within \\rs\\ is determined exclusively in the violent merger phases the cluster undergoes \\citep{Wechsler02, RomanoDiaz06}. The quiescent accretion phases that follows the violent mergers do not alter the inner structure but leads to a continuous growth of virial radius. This suggests that the entropy core emerges as a result of processes occurring during the violent phases, in which the cluster is not in NFW equilibrium. Another alternative is that the gas component, unlike the DM, undergoes some secular evolution of the inner region in the quiescent phases. It might also be the case that both processes play an important role. In any case it is clear that only shock waves can form the entropy core and the understanding of the shock pattern is critical here.   N-body simulations show that clusters have experienced merging events in their past. The merging events of comparable sized substructures are violent enough to send bow shocks into the intra-substructure gas which subsequently  propagate to the outskirts of the  clusters \\citep{Nagai03}. This shocks raise the entropy of the gas. \\cite{Pfrommer06} develop a formalism for the identification and accurate estimation of the strength of structure formation shocks during cosmological SPH simulations. They find that most of the energy is dissipated in weak internal shocks with Mach numbers $\\mathcal{M}=2$ which are predominantly central flow shocks or merger shock waves traversing halo centres. The study by \\cite{Navarro95} utilises cosmological simulations to compare the evolution of DM and gas of clusters between $z=1$ and $z=0$. They find that the DM evolves at roughly constant phase space density, whereas shocks increase the temperature of the gas with relatively little change in density. \\cite{Ascasibar04} investigating pure N-body simulations  report the remarkable stability of the phase space density profile of the DM independent of the dynamical state. \\cite{Pearce94} explore simple head-on mergers between model galaxy clusters containing a mixture of gas and dark matter, testing the coupling of the gas to the underlying collisionless material. After the first passage of the dark matter cores the gas forms a pancake-like structure about the centre of mass. While separating for a second time dark matter cores leave their gas behind. The central gas is no longer confined to a deep potential well so it expands. This process has been nicely confirmed by the recent observations of \\cite{Clowe06}. They find an obvious displacement of the gas measured by X-ray observations compared to the potential revealed by weak lensing maps of the merging cluster system 1E 0657-558. At the same time gravitational attraction due to the central gas exerts a drag force on the separating DM cluster cores, causing an energy transfer from the DM to the gas (see also \\citealt{Navarro93}). This is in agreement with recent investigations by \\cite{Lin06} reporting a $\\sim10$ per cent increase in the concentration of the cluster's DM component when adiabatic gas physics is included as compared to pure DM simulations.  The core structure of the gas entropy, in contrast to the power-law profile of the DM, may therefore be established during merging events in consequence of the impenetrability of the gas as opposed to the collisionless nature of the DM.  A different point of view  might be provided by the following speculations. The DM 'temperature' profile is characterised by a temperature inversion. Now, a 'normal' thermodynamic system with a mechanism for heat transport and no energy sinks or sources cannot sustain such a structure. Heat would flow inwards, leading to an increase of the entropy where the temperature inversion prevails. The increase in entropy amounts to an increase in temperature, and thereby decrease in density so as to keep the hydrostatic equilibrium. Now, let us conjecture that the collapse of the gas and DM results in a configuration where the gas follows the DM distribution. (This  is supported by the self-similar solutions of Bertschinger 1985 and Chuzhoy \\& Nusser 2000). If the gas had a mechanism of transporting energy it would behave in the manner described above, namely the entropy   will increase  where the initial temperature inversion initially existed. Inspection of Fig.~\\ref{fig:select3} proves that this might have actually happened. The gas entropy departs from the DM and an entropy core replaces the temperature inversion. This is clearly manifested by the most relaxed cluster of Fig.~\\ref{fig:select3}, and  is seen in   other relaxed clusters (not shown here). So, what can provide the heat transport mechanism? Fig.~\\ref{fig:select3} shows that outside the entropy core about $(25-30)$ per cent of the non-gravitational energy is in the form of kinetic energy. Only a small fraction of this energy is in large scale bulk motion and therefore most of it is in small scale motion, namely turbulence. The dissipation of these motions by weak shocks is expected to be efficient at the centre and this would lead to the increase in the gas entropy. In the language of one dimensional spherically symmetric hydrodynamics one would attribute this phenomenon to turbulent viscosity and heat transport. This transport mechanism erases the temperature inversion and produces the entropy core.  The application of   the concept of maximum entropy  as a mean of calculating the equilibrium configuration   of   self gravitating systems has proven somewhat  futile (e.g.~\\citealt{LyndenBell67,LyndenBell68,White87,LimaNeto99}). In particular,   these attempts   have not predicted the NFW density profile that  faithfully reproduces the structure of DM halos. The problem that lies at the root of the shortcoming of the thermodynamic approach is the long range nature of the gravitational force and the resulting negative specific energy (see e.g.~\\citealt{LyndenBell68}). In spite of all that our results show that the concept of the  DM entropy is closely related to the structure of DM halos. Moreover, the DM entropy is very closely related to the gas entropy. Beyond the entropy core of the gas the entropies of the two components are tracing one another. In fact,  the extended gas entropy fully coincides with the DM entropy in that regime. It follows that the observed power law behaviour of the gas entropy provides an indirect support for the NFW structure of the DM. Now, this opens the door for a fundamental question. The gas entropy results from the conversion of gravitational to kinetic energy and by the dissipation of the kinetic energy into internal one by means of shock waves. The power law regime seems to be dominated by the accretion shock that is predicted by self-similar solutions \\citep{Bertschinger85,Chuzhoy00}. In the DM sector  the gravitational energy is converted to random motions  by means of phase mixing and the time variation of the  gravitational potential. The latter process happens by violent relaxation, whereby the gravitational potential changes rapidly, or by means of adiabatic changes as envisaged in the secondary infall model \\citep{Gunn77, Fillmore84, Bertschinger85, Hoffman85}. The physical processes that shape the entropy profile in the gas and the DM appear to be very different, yet the resulting structure seems to be identical. This phenomenological finding  needs to be further investigated.  Additional studies of the adiabatic formation of clusters of galaxies are needed to deepen our understanding of the role of entropy in the formation and evolution of clusters of galaxies. In particular, high resolution gas dynamical simulations need to be analysed with a special emphasise on the role of strong, large scale, shocks and of weak, small scale, ones. The relevance of these to the emergence of the entropy core needs to be further clarified. Finally, the simulations should be expanded so as to include preheating, radiative cooling and feedback processes."
    },
    "0608/astro-ph0608418_arXiv.txt": {
        "abstract": "{} {We study the hard X-ray luminosity function and absorption distribution of local ($z\\lesssim 0.1$) active galactic nuclei (AGN) and discuss the implications for AGN cosmological evolution and for the cosmic X-ray background (CXB). } {We use the INTEGRAL all-sky hard X-ray survey to perform a statistical study of a representative sample of nearby AGN. Our entire all-sky sample consists of 127~AGN, of which 91 are confidently detected ($>5\\sigma$) on the time-averaged map obtained with the IBIS/ISGRI instrument and 36 are detected only during single observations. Among the former there are 66 non-blazar AGN located at $|b|>5^\\circ$, where the survey's identification completeness is $\\sim 93$\\%, which we use for calculating the AGN luminosity function and X-ray absorption distribution. } {In broad agreement with previous studies, we find that the fraction $f_{\\rm a}$ of obscured ($\\log\\nh>22$) objects is much higher ($\\sim 70$\\%) among the low-luminosity AGN ($\\lhx<10^{43.6}$~erg~s$^{-1}$) than among the high-luminosity ones ($\\lhx>10^{43.6}$~erg~s$^{-1}$), $f_{\\rm a}\\sim 25$\\%, where $\\lhx$ is the luminosity in the 17--60~keV energy band. We also find that locally the fraction of Compton-thick AGN is less than 20\\% unless there is a significant population of AGN that are so strongly obscured that their observed hard X-ray luminosities fall below $\\sim 10^{40}$--$10^{41}$~erg~s$^{-1}$, the effective limit of our survey. The constructed hard X-ray luminosity function has a canonical, smoothly broken power-law shape in the range $40<\\log\\lhx<45.5$  with a characteristic luminosity of $\\log\\lb=43.40\\pm 0.28$. The estimated local luminosity density due to AGN with $\\log\\lhx>40$ is $(1.4\\pm 0.3)\\times 10^{39}$~erg~s$^{-1}$~Mpc$^{-3}$ (17--60~keV). We demonstrate that the spectral shape and amplitude of the CXB are consistent with the simple scenario in which the $\\nh$ distribution of AGN (for a given $\\lhx/\\lb(z)$ ratio) has not changed significantly since $z\\sim 1.5$, while the AGN luminosity function has experienced pure luminosity evolution. } {} ",
        "introduction": "\\label{s:intro}  Recent deep extragalatic X-ray surveys (see \\citealt{brahas05} for a review) have significantly advanced our understanding of the cosmological evolution of active galactic nuclei (AGN) and the growth of massive black holes. In particular, these surveys have revealed the phenomenon of AGN ``downsizing'' , i.e. gradual transition from the dominance of luminous quasars at high redshift ($z\\gtrsim 2$) to that of lower luminosity AGN at $z\\lesssim 1$. Another important result is the discovery of numerous type~2 (i.e. with significant X-ray absorption) AGN at medium and high redshifts. Despite these successes, the census of AGN in the Universe may still be significantly incomplete, in particular because deep X-ray surveys, performed effectively at energies below $\\sim 8$~keV, are biased against detecting heavily obscured sources. Furthemore, since deep extragalatic X-ray surveys usually cover small areas on the sky, they do not probe the bright end of the AGN luminosity distribution at low redshift ($z\\lesssim 0.3$). It is in these two overlapping areas -- heavily obscured AGN and luminous nearby AGN -- where deep X-ray surveys could be efficiently complemented by all-sky hard X-ray (above $\\sim 10$~keV) surveys.  Recently, a serendipitous all-sky survey in the 3--20~keV energy band was performed with the RXTE observatory (the RXTE Slew Survey, or XSS), in which $\\sim 100$ AGN were detected at $|b|>10^\\circ$ \\citep{revetal04}. This sample was used to construct an X-ray luminosity function of local ($z\\lesssim 0.1$) AGN and to study their absorption distribution \\citep{sazrev04}. However, despite its harder X-ray range compared to traditional X-ray surveys, the XSS was still biased against finding sources with absorption columns $\\nh\\gtrsim 10^{23}$~cm$^{-2}$.  A comparable or better (for heavily obscured sources and/or in crowded sky regions) capability for studying bright ($\\gtrsim 1$~mCrab) hard X-ray sources is now provided by the IBIS telescope \\citep{ubeetal03} of the INTEGRAL observatory \\citep{winetal03}, owing to its good sensitivity above 20~keV, large field of view ($\\sim 28^\\circ \\times 28^\\circ$), and good ($\\sim 10^\\prime$) angular resolution. Since its launch in October 2002, INTEGRAL had observed $\\sim 60$\\% of the sky by the spring of 2005, mostly around the Galactic plane and the Galactic center. This motivated us to initiate a series of pointed INTEGRAL observations (Proposal 0320108) aimed at covering the remaining ``empty'' (extragalactic) sky regions. This campaign has now been completed and practically full coverage of the sky has been achieved after all public and our proprietary INTEGRAL observations were added.  Based on the all-sky hard X-ray map obtained and the follow-up source identifications, we compiled a catalog of detected sources \\citep{krietal06}. In the present paper we use the INTEGRAL all-sky survey to study such key statistical properties of the local AGN population as the hard X-ray luminosity function and absorption distribution. We also discuss the implications for the cosmological evolution of AGN and for the cosmic X-ray background. We note that there have been previous attempts to use INTEGRAL for studying the statistics and demography of nearby AGN \\citep{krietal05,basetal06,becetal06}. The significant advantage of our present study is that it is based on an all-sky, largely serendipitious survey, which allowed us to significantly alleviate the previously significant bias associated with INTEGRAL AGN pointings. We also note that a similar all-sky hard X-ray survey has since recently been carried out by the Swift observatory \\citep{maretal05}. ",
        "conclusions": ""
    },
    "0608/astro-ph0608132_arXiv.txt": {
        "abstract": "We discuss the reference time $t_0$ of afterglow light curves in the context of the standard internal-external shock model. The decay index of early afterglow is very sensitive to the reference time one chooses. In order to understand the nature of early afterglow, it is essential to take a correct reference time. Our simple analytic model provides a framework to understand special relativistic effects involved in early afterglow phase. We evaluate light curves of reverse shock emission as well as those of forward shock emission, based on full hydrodynamic calculations. We show that the reference time does not shift significantly even in the thick shell case. For external shock emission components, measuring times from the beginning of the prompt emission is a good approximation and it does not cause an early steep decay. In the thin shell case, the energy transfer time from fireball ejecta to ambient medium typically extends to thousands of seconds. This might be related to the shallow decay phases observed in early X-ray afterglow at least for some bursts. ",
        "introduction": "It is well known that gamma-ray burst (GRB) afterglow decays as a power-law $L\\propto(t-t_0)^{-\\alpha}$.  The temporal decay index $\\alpha$, together with the spectral index, provides us precious information about GRB jets and their environment. In the pre-Swift era afterglow observations start typically a few hours after a burst. In such a late phase, the decay index is insensitive to the choice of the reference time $t_0$, and GRB trigger time is often used in afterglow modelings.  The multi-wavelength observatory Swift was launched in Nov 2004. Thanks to its fast pointing capabilities Swift is disclosing the early afterglow phase. One of unexpected finds by Swift is that early X-ray afterglows show a canonical behavior, where light curves include three components: (1) a steep decay component, (2) a shallow decay component and (3) a ``normal'' decay component. On top of this canonical behavior, many events have superimposed X-ray flares (Zhang et al. 2006a; Nousek et al. 2006; Chincarini et al 2005; O'Brien et al. 2006). The transition from the early steep decay to the shallow decay typically occurs at several hundred seconds, and the timescale is comparable to the duration of rather long GRBs. When discussing the early afterglow and its connection to the prompt emission component, the decay index is very sensitive to the reference time $t_0$ one chooses. Correctly choosing $t_0$ is therefore essential to derive the right index as well as to interpret each component in the canonical light curve (Piro et al. 2005; Tagliaferri et al. 2005; Quimby et al. 2006).  Tagliaferri et al. (2005) investigated the first two bursts GRB 050126 and GRB 050219a which have an X-ray light curve well sampled by the X-Ray Telescope on board Swift. They sought for a possible delay of the afterglow onset by fitting the early X-ray light curves (the components 1 and 2 which we have discussed above) with a single power-law model. In both cases, the decaying light  curves can be fitted if the onset of the afterglow is shifted to $t_0 \\sim 100$ s after the burst trigger with a single power-law. However, while in the case of GRB 050126 the light curve does not allow us to clearly state whether a broken power-law modeling is better than a single power-law model, for GRB 050219a a broken power-law definitively provides a better fit.  In the standard GRB model the time shift between the GRB trigger and the reference time $t_0$ is expected to be ``small''. The early steep decay should not be an artifact due to a wrong choice of $t_0$. Lazzati and Begelman (2006) studied forward shock emission, based on a simple energy injection model. Their numerical light curves show that measuring times from the beginning of the prompt phase is a good approximation. The early steepening and X-ray flares are likely to be  produced by another mechanism  (e.g. internal shocks; Burrows et al. 2005; Falcone et al. 2006; Zhang et al. 2006a; Nousek et al. 2006; Ioka, Kobayashi \\& Zhang 2005; Fan \\& Wei 2005). Recently long-lasting soft emission is reported in a short burst GRB 050724 (Barthelmy et al. 2005). Such a soft component was hinted at in the sum of multiple short BATSTE GRBs, and it might be the onset of short burst afterglow (Lazzati, Ramirez-Ruiz \\& Ghisellini 2001). It is therefore of interest to quantitatively examine the $t_0$ issue.  In this paper, we study the physics and timescales involved in early afterglow stage, and give more direct and clear arguments for the afterglow slopes. In \\S 2 we study a simple analytic model. In \\S 3 we evaluate light curves of reverse shock emission as well as that of forward shock emission, based on full hydrodynamic calculations. In \\S 4 we address how inhomogeneity of a fireball affects an early afterglow light curve. Conclusions and discussion are given in \\S 5. ",
        "conclusions": "We have studied the reference time $t_0$ for the afterglow modeling. Although measuring times from the beginning of the prompt emission (GRB trigger) might cause a slight overestimate of the early afterglow slope in the thick shell case. This choice of $t_0$ gives a reasonable approximation, and it does not induce a very steep decay ($\\alpha\\sim 3-5$) like the early steep decay or X-ray flares  in the canonical Swift X-ray light curve.  The leading model to explain the rapid decay and flares in early X-ray afterglow is the internal shock emission. A clear, testable prediction of this model is that the temporal decay index $\\alpha$ of the tail part should be related to the spectral index $\\beta$ by an equation $\\alpha=2+\\beta$ (Kumar \\& Panaitescu 2000). When evaluating the emission decay $L\\propto (t-t_0)^{-\\alpha}$ in the internal shock model, an important difference is that the GRB trigger time is no longer special. The reference time $t_0$ should correspond to the onset of each particular spike in the prompt emission or in afterglow (Kobayashi et al. 1997; Zhang et al. 2006a; Nousek et al. 2006; Fan \\& Wei 2005). Every time when the central engine is re-stared to eject sub-shells, the reference time $t_0$ should be re-set to the reactivation time of the engine. Although O'Brien et al. (2006) have found that $\\alpha$ appears to be largely independent of $\\beta$ when the BAT trigger time is used as $t_0$, Liang et al (2006) have shown that the relation $\\alpha=2+\\beta$ is more or less satisfied in most cases if $t_0$ is set near the beginning of rising segment of the last pulse of the prompt emission or a corresponding X-ray flare, and if the underlying forward shock emission component is subtracted. Swift observations support the internal shock model.  The self-consistent internal shock interpretation should be more favorable than the beginning-of-the-afterglow interpretation. The latter can not explain multiple X-ray flares in a single event. Such behavior is observed in many Swift bursts (Burrows et al. 2005; Falcone et al. 2006; Romano et al. 2006; O'Brien et al 2006). Usually X-ray flux already begins to decay before X-ray flares appear, and it suggests that the onset of afterglow is prior to the flares. If a large amount of energy is impulsively injected to a fireball during the deceleration, $t_0$ might be re-set to the injection time. However, an afterglow baseline also should shift after a flare (energy injection). This clearly contradicts with observations in which after a flare peak, afteglow decays back to its pre-flare flux level. We can not explain X-ray flares by the shift of $t_0$ associated with large energy injections.  Swift discovered that a large fraction of X-ray afterglows have a slow decay phase, and it is suggested that energy injection into a blast wave takes place several hundred seconds after the burst. This implies that right after the burst the kinetic energy of a blast wave is very low and in turn the efficiency of internal shock process is extremely high (Zhang et al. 2006a; Nousek et al. 2006; Ioka et al. 2005; Granot, K\\\"onigl \\& Piran 2006; Zhang et al.2006b; however see also Fan \\& Piran 2006). The round-off forward shock peak in the thin shell case might be a good candidate for the shallow decay phase. However, if we interpret the observed shallow decay as the round-shape of the deceleration phase, the model light curve is shallower than some of the observed ones. It may be because the observed curve is the combination of this round phase and the rapid decay from the GRB tail emission (high latitude). Equalizing the deceleration time and the shallow phase timescale we obtain the initial Lorentz factor $\\Gamma_0 \\sim 110 \\zeta^{3/8} E_{53}^{1/8} \\rho_{1,-1}^{-1/8}t_{d,3}^{-3/8}$ where the time dilation effect is taken into account $\\zeta=(1+z)/3$.  For a wind environment, we can discuss the $t_0$ issue in a very similar way. The Lorentz factor of a shocked shell is constant in both of the thin and thick shell cases during the shock crossing, while the BM blast wave decelerates as $\\Gamma \\propto R^{-1/2}$ (Kobayashi \\& Zhang 2003). The separations are comparable $S_F \\sim 2S_B$ at the deceleration time. Measuring times from the beginning of the prompt phase should be a good approximation for events in a wind environment also. Since most bursts in a wind environment fall in the thick shell case (Kobayashi, \\Mesz \\& Zhang 2004), and since the deceleration time is comparable to GRB duration in the thick shell case, another process (e.g. refreshed shocks) rather than the afterglow peak is necessary to explain the shallow decay phase.  We thank Luigi Piro for useful discussion. This work is supported by NASA NNG05GB67G and NNG06GH62G."
    },
    "0608/astro-ph0608126_arXiv.txt": {
        "abstract": "Observations of low mean metallicity of damped Lyman-alpha (DLA) quasar absorbers at all redshifts studied appear to contradict the predictions for the global mean interstellar metallicity in galaxies from cosmic chemical evolution models. On the other hand, a number of metal-rich sub-DLA systems have been identified recently, and the fraction of metal-rich sub-DLAs appears to be considerably larger than that of metal-rich DLAs, especially at $z < 1.5$. In view of this, here we investigate the evolution of metallicity in sub-DLAs. We find that the mean Zn metallicity of the observed sub-DLAs may be higher than that of the observed DLAs, especially at low redshifts, reaching a near-solar level at $z \\la 1$. This trend does not appear to be an artifact of sample selection, the use of Zn, the use of $N_{\\rm H I}$-weighting, or observational sensitivity. While a bias against very low metallicity could be present in the sub-DLA sample in some situations, this cannot explain the difference between the DLA and sub-DLA metallicities at low $z$. The primary reason for the difference between the DLAs and sub-DLAs appears to be the dearth of metal-rich DLAs. We estimate the sub-DLA contribution to the total metal budget using measures of their metallicity and comoving gas density. These calculations suggest that at $z \\la 1$, the contribution of sub-DLAs to the total metal budget may be several times that of DLAs. At higher redshifts also, there are indications that the sub-DLAs may contribute significantly to the cosmic metal budget. ",
        "introduction": "As the cycle of stellar birth and death progresses in galaxies, the interstellar gas is expected to become increasingly metal-rich. A fundamental prediction of most chemical evolution models is that the mass-weighted mean interstellar metallicity of galaxies should rise from a low value at $z > 2$ to a near-solar value at $z=0$  (e.g., Pei et al. 1999; Somerville et al. 2001; Cora et al. 2003). Indeed, the mass-weighted mean metallicity of local galaxies is close to solar (e.g., Kulkarni \\& Fall 2002; Fukugita \\& Peebles 2004). How was this present level of metallicity reached? A direct way to study this enrichment history is through the absorption lines superposed by galaxies on the spectra of background quasars, in particular the DLAs (defined to have neutral hydrogen column densities log $N_{\\rm HI} \\ge 20.3$) and the sub-DLAs (defined to have $19. 0 \\le$ log $N_{\\rm HI} < 20.3$; P\\'eroux et al. 2003). These absorbers contain sufficient neutral gas mass to account for a significant fraction of the visible stellar mass in present-day galaxies (e.g., Storrie-Lombardi \\& Wolfe 2000; P\\'eroux et al. 2003, 2005; Prochaska et al. 2005). Indeed, hydrodynamical simulations naturally produce DLAs from cool gas accumulating at centers of cold dark matter (CDM) halos and predict DLAs to be closely associated  with star forming regions (e.g., Haehnelt et al. 2000; Cen et al. 2003; Nagamine et al. 2004).  Previous studies of element abundances in high-$N_{\\rm H I}$ quasar absorbers have focussed primarily on DLAs, especially those at high redshift (e.g., Prochaska et al. 2003b). Recently, we have started to expand the low-$z$ DLA Zn data and tripled the $z < 1.5$ sample so far (Khare et al. 2004; Kulkarni et al. 2005; Meiring et al. 2006a, 2006b; P\\'eroux et al. 2006b). These studies find that the DLA mean metallicity stays substantially subsolar even at low redshifts and increases relatively slowly  (see also Prochaska et al. 2003a). Furthermore, based on deep emission-line imaging studies (e.g. Kulkarni et al. 2006 and references therein), the star formation rates (SFRs) of a large fraction of DLAs appear to fall far below the predictions from the cosmic star formation history (inferred from galaxy imaging surveys such as the Hubble Deep Field, e.g., Madau et al. 1998). A missing metals problem has also been noted in high-$z$ DLAs from SFR estimates based on C II* absorption (Wolfe et al. 2003). These results could potentially present important challenges to our understanding of galaxy evolution.  Why do DLA data disagree with the mean metallicity in galaxies predicted from the global star formation history? Smoothed particle hydrodynamics (SPH) simulations (Nagamine et al. 2004) indicate that true DLA metallicities could be 1/3 solar at $z=2.5$ and higher at lower $z$. Has the DLA metal content been underestimated, or is there a larger quantity of metals in a hitherto unexplored absorber population? While dust selection or metallicity gradients can lead to a flat metallicity-redshift relation, the extent of these effects is still uncertain (e.g., Boiss\\'e et al. 1998; Vladilo \\& P\\'eroux 2005; Ellison et al. 2005; Akerman et al. 2005; Chen et al. 2005; Zwaan et al. 2005). It has been suggested recently that dust obscuration is not significant and that metal-rich DLAs are inherently rare (Meiring et al. 2006a; Khare et al. 2006, hereafter K06; Herbert-Fort et al. 2006). In view of this, it is important to explore other possible sites for the missing metals.  Over the past two years, we have been investigating the metallicities of sub-DLAs at $z < 1.5$ (Khare et al. 2004; P\\'eroux et al. 2006a, 2006b; Meiring et al. 2006b). These studies have led to the discovery of five solar or supersolar metallicity sub-DLAs at $0.7 < z < 1.5$. Before these studies began, only a single  metal-rich sub-DLA was identified (Pettini et al 2000). Indeed, five of the nine sub-DLAs with Zn data at $z < 1$ and six of the 17 sub-DLAs at $z < 1.5$ have solar or supersolar metallicity. Despite the small number of measurements (caused by the focus on DLAs in the past), the large fraction of metal-rich sub-DLAs suggests that  log $N_{\\rm H I} \\approx 20.0$ may be the optimum range for revealing metal-rich systems, since such $N_{\\rm H I}$ is high enough to have a significant metal column density and low enough to make dust attenuation less important (Lauroesch et al. 1996; Vladilo \\& P\\'eroux 2005; P\\'eroux et al. 2005, 2006a). To investigate this further, here we examine the evolution of the observed amount of metals in neutral gas in sub-DLAs and their contribution to the cosmic metal budget. ",
        "conclusions": "While the sub-DLA samples are still small and not all the observed sub-DLAs are metal-rich, the fraction of observed metal-rich systems is larger for sub-DLAs than for DLAs. Our calculations suggest that the mean metallicity of the observed sub-DLAs is higher than that of the observed DLAs, reaching a near-solar level as expected in cosmic chemical evolution models. At $z < 1$, we find that the contribution of the observed sub-DLAs to the total metal budget may be several times higher than that of the observed DLAs. At higher $z$, the observed sub-DLAs may contribute $\\ge 32 \\%$ of the DLA contribution to the total metal budget. Furthermore, the results of Prochaska et al. (2005) on the relative number of sub-DLAs and DLAs may triple the contribution of sub-DLAs compared to the results based on P\\'eroux et al. (2005). In general, our results support and extend the suggestions by P\\'eroux et al. (2006a), York et al. (2006), and Prochaska et al. (2006) that sub-DLAs may contribute significantly to the overall global metallicity.  One may wonder whether ionization effects are likely to be responsible for the difference we see between the global metallicities for the DLA and sub-DLA samples. We believe that this is not a significant effect for our samples. First, the $N_{\\rm H I}$ of the sub-DLAs in our sample are relatively high ($86 \\%$ systems with log $N_{\\rm H I} > 19.7$ and $73 \\%$ with log $N_{\\rm H I} > 19.9$ in our primary sample.) Second, photoionization calculations presented in high-resolution studies of sub-DLAs (Dessauges-Zavadsky et al. 2003; Meiring et al. 2006b) show that while sub-DLAs have some ionized gas, the ionization corrections in abundance estimates are not significant, and would increase [Zn/H] if applied. Vladilo et al. (2001), Dessauges-Zavadsky et al. (2003), York et al. (2006), and Meiring et al. (2006b) all suggest that there is no trend in $N_{\\rm Al \\, III}$/$N_{\\rm Al \\, II}$ vs. $N_{\\rm H I}$ in sub-DLAs and that the ratio is smaller than 1. Therefore, using $N_{\\rm Zn II}/ N_{\\rm H I}$ in sub-DLAs for a measure of $N_{\\rm Zn}/N_{\\rm H}$ does not introduce a systematic bias.  It is important to establish whether the observed DLA and sub-DLA samples reflect their true properties or differential dust selection effects. If dust obscuration is significant and depends primarily on the metal column density (e.g., Vladilo et al. 2006), then metal-rich DLAs may be more affected by dust selection than metal-rich sub-DLAs (Vladilo \\& P\\'eroux 2005). Such an effect could make DLAs appear less metal-rich than sub-DLAs; but, this alone would not explain the large difference between DLAs and sub-DLAs seen in Fig. 1, because so far there is no evidence for the necessary large number of high extinction DLAs (K06). K06 use this result and the observed mass-metalicity relationship for galaxies (e.g., Tremonti et al. 2004; Savaglio et al. 2005; Erb et al. 2006) to conclude that sub-DLAs may arise in massive galaxies and DLAs in less massive galaxies, consistent with the relatively stronger evolution of metals in sub-DLAs compared to DLAs that we find here. That conclusion is also consistent with trends in gas dispersion with metalicity of quasar absorbers (Ledoux et al 2006, Meiring et al. 2006b). [But, see an opposing opinion on this last point by Bouch\\'e et al. (2006b), using a less direct argument.]  The relatively smaller sample size for sub-DLAs  compared to DLAs and the conflicting results on gas velocity dispersion just referred to emphasize the need for much larger samples of sub-DLAs and better velocity dispersion measurements for the absorbers at a range of redshifts, to confirm the results reached here. Large telescopes are essential to obtain adequate samples of excellent abundance and velocity dispersion measurements. The sub-DLA sample needs to be made at least as large as the DLA sample. It will also be important to obtain $N_{\\rm H I}$ and [Zn/H] for absorbers toward UV-faint quasars, to confirm that there are not a large number of highly reddened quasars with foreground DLAs. More extensive observations with the HST Cosmic Origins Spectrograph to measure H I at $z < 1.5$ and Zn II  at $z < 0.55$, and with ground-based spectrographs to measure Zn II at $0.55 < z < 1.5$ will be very important to confirm the findings presented here ."
    },
    "0608/astro-ph0608310_arXiv.txt": {
        "abstract": "Suzaku deep observations have discovered two highly significant nonthermal X-ray sources, Suzaku~J1804$-$2142 (Src~1) and Suzaku~J1804$-$2140 (Src~2), positionally coincident with the unidentified TeV $\\gamma$-ray source, HESS~J1804$-$216. The X-ray sources are not time variable and show no counterpart in other wavebands, except for the TeV source. Src~1 is unresolved at Suzaku spatial resolution, whereas Src~2 is extended or composed of multiple sources. The X-ray spectra are highly absorbed, hard, and featureless, and are well fitted by absorbed power-law models with best-fit photon indices and absorption columns of $-0.3_{-0.5}^{+0.5}$ and $0.2_{-0.2}^{+2.0}\\times 10^{22}$~cm$^{-2}$ for Src~1, and $1.7_{-1.0}^{+1.4}$ and $1.1_{-0.6}^{+1.0}\\times 10^{23}$~cm$^{-2}$ for Src~2. The measured X-ray absorption to the latter source is significantly larger than the total Galactic neutral hydrogen column in that direction. The unabsorbed 2--10~keV band luminosities are $7.5\\times 10^{32}(d/{\\rm 5~kpc})^2$~ergs~s$^{-1}$ (Src 1) and $1.3\\times 10^{33}(d/{\\rm 5~kpc})^2$~ergs~s$^{-1}$ (Src 2), where $d$ is the source distance. Among the handful of TeV sources with known X-ray counterparts, HESS J1804$-$216 has the largest ratio of TeV $\\gamma$-ray to hard X-ray fluxes. We discuss the nature of the emission and propose the Suzaku sources as plausible counterparts to the TeV source, although further observations are necessary to confirm this. ",
        "introduction": "The origin and acceleration mechanism of cosmic rays has remained a long-standing problem since their discovery nearly 100 years ago. The most widely accepted acceleration mechanism, at least for  cosmic rays up to $10^{15.5}$~eV (the knee energy), is diffusive shock acceleration at the rims of young supernova remnants (SNRs) (e.g., \\cite{bell1978,blandford1978}). This hypothesis was greatly strengthened by the discovery of localized regions of synchrotron X-ray emission from the rims of SNR SN~1006 \\citep{koyama1995}, which indicated that SNRs were able to accelerate electrons up to $\\sim$TeV energies. Since then, several other SNRs have been identified as electron accelerators (e.g., \\cite{koyama1997,slane2001,bamba2000,bamba2001}) and numerous studies have indicated that the acceleration process is very efficient (e.g., \\cite{vink2003,bamba2003,bamba2005,voelk2005}). Many papers regard RX~J1713$-$3946 as a proto-type of cosmic-ray accelerating SNRs, since it is bright in both X-ray and TeV bands \\citep{uchiyama2003}. To date the best evidence for the acceleration of hadronic cosmic rays comes from observations of the Tycho SNR \\citep{warren2005}.  Recently, a TeV $\\gamma$-ray telescope, High Energy Stereoscopic System (HESS), has discovered some new sources in the Galactic plane \\citep{aharonian2005,aharonian2006}. Some of these sources are extended on scales of $\\sim \\timeform{0.1D}$, suggesting Galactic counterparts like pulsar wind nebulae (PWNe), SNRs, star-forming regions, and so on. However, their nature is still unclear: mainly because many of the new sources have no counterparts in other wavebands. They have been sometimes referred to as ``dark particle accelerators''. Deep follow-up observations are crucial to elucidate their nature and hard X-ray observations  are a valuable tool for this, since above a few keV we do not have to contend with absorption by the Galactic plane. The X-ray satellite Suzaku has large effective area and low background \\citep{mitsuda2006}, making it an excellent choice for such follow-up studies.   HESS~J1804$-$216 is one of the brightest of the unidentified (unID) TeV sources \\citep{aharonian2005}. Its size ($\\timeform{0.2D}$) and photon index ($2.72\\pm0.06$) makes it one of the largest and softest sources. \\citet{aharonian2006} proposed two possible counterparts: the radio SNR G8.7$-$0.1 (W30) and a young Vela-like pulsar PSR~J1803$-$2137, although neither is a definitive identification. The HESS source overlaps only the southwestern portion of G8.7$-$0.1, which is about $3/4^\\circ$ in diameter. The distance is estimated to be $\\sim$5~kpc based on associating the SNR with nearby  H~II regions \\citep{kassim1990,finley1994}. From CO observations, the surrounding region is known to be associated with molecular gas where massive star formation takes place \\citep{ojeda-may2002}. From the observed dispersion measure, PSR~J1803$-$2137 is located at a distance of $\\sim$5.3~kpc \\citep{kassim1990}. Both candidates were detected by ROSAT \\citep{finley1994}; the soft X-ray emission of G8.7$-$0.1 is only apparent along the northeastern radio rim. Very recently, \\citet{brogan2006} reported the discovery of a new SNR (G8.31$-$0.9) in this region. Swift independently discovered some X-ray sources in this region recently \\citep{landi2006}, but the statistics was quite limited due to the short observation. It is interesting to note that the AGASA and SUGAR groups have reported an excess of EeV cosmic rays from the general direction (i.e., within $\\sim$20$^\\circ$) toward HESS~J1804$-$216 \\citep{hayashida1999,bellido2001}, although this result has not been confirmed by more sensitive observations by the Pierre Auger Observatory \\citep{letessier2005}. The evidence that HESS~J1804$-$216 is a proton accelerator therefore remains ambiguous \\citep{fatuzzo2006}.  In this paper, we report the first results of Suzaku follow-up observations of HESS~1804$-$216. \\S\\ref{sec:obs} introduces the observations, our analysis is explained in \\S\\ref{sec:analysis} and we discuss our findings in \\S\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  The new hard X-ray sources we have successfully detected with Suzaku are now the nearest potential counterparts to HESS~J1804$-$216. Although the morphology of TeV $\\gamma$-ray emission and hard X-rays are quite different, this is also the situation for the radio SNR G8.7$-$0.1 and the soft X-ray emission.  None of the possible counterparts is a morphological match to the HESS source. Nevertheless the new X-ray sources provide interesting new information on the nature of HESS~J1804$-$216 and the environment. An independent report \\citep{landi2006} notices new X-ray sources in this region using Swift. The position and the flux of the source nearest to HESS~J1804$-$216 is consistent with those of Src~2, indicating that we successfully confirmed their detection. We estimated the chance coincidence probability of unrelated X-ray sources for HESS~J1804$-$216. With the $\\log N$-$\\log S$ relation of X-ray sources in the 2.0--10.0~keV band \\citep{sugizaki2001,yamauchi2002,ebisawa2005}, we found that the expected number of X-ray sources is only 4--9$\\times 10^{-3}$ in the elliptical region with radii of 0.016 and 0.018 degrees (the position uncertainty of the HESS source; \\cite{aharonian2006}). Therefore, we concluded that they are physically associated with HESS~J1804$-$216. In the following we discuss Src~1 and Src~2 separately.  The best-fit absorbing column for Src~1, $0.2_{-0.2}^{+2.0}\\times 10^{22}$~cm$^{-2}$, is consistent with the Galactic neutral hydrogen column in this direction of $1.5\\times 10^{22}$~cm$^{-2}$ \\citep{dickey1990}. The unabsorbed luminosity is $7.5\\times 10^{32}(d/{\\rm 5~kpc})^2$~ergs~s$^{-1}$ in the 2.0--10.0~keV band, where $d$ is the source distance.  We quote results for a nominal value of 5 kpc although in fact we have no independent  constraints on the distances to the new X-ray sources. Src~1 has a very flat spectrum, flatter than typical for young pulsars and their associated wind nebulae (PWN) (typical $\\Gamma$ values of 1.5 to 2.5; \\cite{possenti2002}) or shell-like SNRs emitting synchrotron X-rays (typical $\\Gamma$ values of 2.3 to 3). \\citet{uchiyama2002} have reported a region near SNR G78.2+2.1 (Gamma-Cygni) with a photon index of 0.8--1.5, that might be comparable to Src~1. High mass X-ray binaries also have hard photon indices ($\\Gamma$ of 0.5--2.5; e.g. \\cite{muno2003}). So Src~1 might be this type of source.  In contrast Src~2 has a best-fit absorbing column that is about an order-of-magnitude higher than the expected Galactic column. This is consistent with Src~2 being embedded in dense gas, although there are no cataloged molecular clouds in this direction. The X-ray and radio counterpart to HESS~J1813$-$178 \\citep{brogan2005} also appears to show evidence for excess absorption. The unabsorbed 2--10 keV band luminosity of Src~2, $1.3\\times 10^{33}(d/{\\rm 5~kpc})^2$~ergs~s$^{-1}$, best fit photon index (see Table~\\ref{xis_spec_para}), and extended nature are all consistent with either SNR origin: a PWN or shell-like synchrotron emission.  Their origin is still unknown, but there is some interesting scenario for their nature; old SNRs colliding with molecular clouds. \\citet{yamazaki2006} argue that an old SNR shock (age of $10^5$ yr) interacting with a giant molecular cloud can emit hard nonthermal X-rays and strong TeV $\\gamma$-rays from shock accelerated protons through $\\pi^0$ decay (see also {\\cite{fatuzzo2006}). They can explain large TeV $\\gamma$-ray to X-ray flux ratios (factors of $\\sim$100). Src~2 might represent the X-ray component to the HESS source under this scenario, with the large excess absorption as a key piece of evidence.  Comparing images in other wavelengths also suggests the SNR origin. Soft X-ray emission and HESS source are both encompassed by the radio contours of G8.7$-$0.1, suggesting a common origin, although they do not overlap, On the other hand, the heavily absorbed Suzaku-discovered hard X-ray sources are located right at the position of peak HESS emission. The anti-correlation of HESS and ROSAT emission is likely a result of a dense region near the SNR which absorbs the soft X-rays, while at the same time enhancing the production of TeV emission from $\\pi^0$ decay. The larger size of the TeV source compared to the hard X-ray source might be due to the much larger diffusion length of protons (which provide the TeV $\\gamma$-rays) compared to primary electrons (which emit the hard X-rays).  HESS~J1804$-$216 has the large ratio of TeV $\\gamma$-ray to hard X-ray fluxes, which should be an important key to solve its nature. Table~\\ref{ratio} summarizes X-ray and TeV $\\gamma$-rays fluxes for a sample of TeV $\\gamma$-ray emitters. Generally speaking, PWNe have rather small values of $F_{TeV}/F_X$, whereas SNRs tend to be brighter in TeV $\\gamma$-rays. For HESS~J1804$-$216, we assumed three counterparts; Src~1, 2, and the upper limit of the diffuse emission. We extracted the diffuse emission from the entire region of XIS FOV, whereas the TeV emission comes from wider region ($\\sigma$ = 0.2 degree; \\cite{aharonian2006}). Therefore, we re-normalized the flux by assuming that X-ray comes uniformly from the same region of TeV $\\gamma$-rays. The flux in the TeV band is the integrated one from the entire region of the HESS source. The ratio of HESS~J1804$-$216 is significantly larger than that of PWNe \\citep{willingale2001,aharonian2004a}, and known shell-like SNRs emitting synchrotron X-rays \\citep[for example]{ozaki1998,aharonian2005c}. In fact, the $F_{TeV}/F_X$ value is the largest among the identified objects, even if we include the upper-limit to the total hard X-ray diffuse emission from the entire XIS field of view The value is consistent with that of HESS~J1303$-$631, one of the new class of  ``dark particle accelerators'' (\\cite{mukhergee2005,aharonian2005e}, see Table~\\ref{ratio}).  Under the assumption that the TeV $\\gamma$-rays are emitted via inverse Compton on relativistic electrons, higher values of $F_{TeV}/F_X$ suggest weaker magnetic fields, all other things being equal. Therefore, the large $F_{TeV}/F_{X}$ of HESS~J1804$-$216 suggest that the magnetic field in this source is weaker than typical in traditional PWNe and SNRs, or inverse Compton origin is unlikely."
    },
    "0608/astro-ph0608060_arXiv.txt": {
        "abstract": "We discuss spectroscopy and infrared photometry for a complete sample of $\\sim$800 galaxies in close pairs objectively selected from the CfA2 redshift survey. We use 2MASS to compare near infrared color-color diagrams for our sample with the Nearby Field Galaxy Sample and with a set of IRAS flux-limited pairs from Surace et al. We construct a basic statistical model to explore the physical sources of the substantial differences among these samples. The model explains the spread of near infrared colors and is consistent with a picture where central star formation is triggered by the galaxy-galaxy interaction before a merger occurs. For 160 galaxies we report new, deep JHK photometry within our spectroscopic aperture and we use the combined spectroscopic and photometric data to explore the physical conditions in the central bursts. We find a set of objects with H-K $\\geq$ 0.45 and with a large F$_{FIR}$/F$_H$. We interpret the very red H-K colors as evidence for 600---1000K dust within compact star-forming regions, perhaps similar to super-star clusters identified in individual well-studied interacting galaxies. The galaxies in our sample are candidate ``hidden'' bursts or, possibly, ``hidden'' AGN. Over the entire pair sample, both spectroscopic and photometric data show that the specific star formation rate   decreases with the projected separation of the pair. The data suggest that the near infrared color-color diagram is also a function of the projected separation; all of the objects with central near infrared colors indicative of bursts of star formation lie at small projected separation. ",
        "introduction": "Larson \\& Tinsley (1978) first recognized the close connection between galaxy-galaxy interactions and star formation. They suggested that this relationship is a fundamental ingredient of galaxy formation and evolution. Studies of individual interacting systems and of statistical samples of close pairs in the local universe have confirmed and refined the initial results (e.g. Kennicutt \\& Keel 1984; Kennicutt et al. 1987; Jones \\& Stein 1989; Sekiguchi \\& Wolstencroft 1992; Keel 1993; Liu \\& Kennicutt 1995a,b; Keel 1996; Donzelli \\& Pastoriza 1997; Barton, Geller, \\& Kenyon 2000 (BGK hereafter); Barton Gillespie, Geller, \\& Kenyon (2003; BGK2 hereafter); Lambas et al. 2003; Allam et al. 2004; Alonso et al. 2004; Nikolic et al. 2004). Examinations of the universe at high redshift underscore the importance of interactions in molding the evolution of galaxies (e.g. Patton et al. 2002; Conselice et al. 2004; Lin et al. 2004; Papovich et al. 2005; Bell et al. 2006).  The connection between interacting pairs of galaxies and star formation is now well established observationally and there is a promising correspondence between the predictions of simulations and the properties of well-defined samples of close pairs. The data are consistent with a burst of star formation triggered by a close galaxy-galaxy interaction; the burst continues and ages as the galaxies move apart (Mihos et al. 1991; Mihos \\& Hernquist 1996; BGK, BGK2, Lambas et al. 2003; Nikolic et al. 2004). The primary feature of the data is an anti-correlation between the projected separation of a pair on the sky and measures of the normalized star formation rate. This anti-correlation suggests a dependence of the burst strength and probably burst age on the separation of the galaxy pair (BGK2).  Here we focus on  near infrared observations of close pairs of galaxies in the full BGK sample with attention to the distinctive features of infrared emission possibly associated with the interaction. The literature contains detailed near infrared observations of famous individual interacting galaxy pairs (e.g. Lopez-Sanchez et al. 2004) and a few optical-infrared imaging surveys of small objectively selected samples of close pairs. Bushouse \\& Werner (1990) imaged a sample of 22 interacting galaxies. In their sample, the nuclei are redder than the outer regions by $\\sim 0.4$ and 0.9 mag in J-K and R-K, respectively. They attribute the color gradient mostly to an increase in the amount of dust in the nuclear regions and argue that the infrared colors are unaffected by dust emission. Cutri and McAlary (1985) carried out small aperture photometry of a larger sample of pairs. They concluded that a larger fraction of interacting than non-interacting galaxies have J-K and H-K colors outside the normal range. Giuricin et al. (1993) used a compendium of data from the literature to investigate the effect of interactions on the near infrared properties of spiral galaxies. In contrast with Bushouse and Werner (1990), they concluded that apparently interacting galaxies have H-K and K-L excesses indicative of thermal emission from hot dust probably related to star formation induced by the interaction. We use our spectroscopy and matched small aperture photometry along with 2MASS photometry to revisit these issues in a much larger dataset of nearly 800 galaxies in close pairs and n-tuples.  In the far infrared, IRAS observations still yield the largest datasets for consideration of the properties of galaxy-galaxy interactions as a function of observable descriptors of the pair (see, for example, Goto 2005). Telesco, Wolstencroft \\& Done (1988) compiled IRAS observations of pairs selected from the catalog of Arp \\& Madore (1987) and resolved by IRAS. In their set of 93 pairs of comparable luminosity, they find that pairs with the highest far infrared color temperature have the smallest projected separation on the sky. They concluded that interactions measurably enhance the intensity or efficiency of star formation. Bushouse et al. (1988) reached similar conclusions based on IRAS observations of a sample of pairs selected on the basis of optical morphology.  Recently Surace, Sanders \\& Mazzarella (2004; SSM hereafter) used HIRES image reconstruction to resolve the IRAS emission at 12, 25, 60, and 100 $\\mu$m for 106 interacting galaxy systems where the galaxies are separated by less than three average galaxy diameters. The sample has a flux limit at 60 $\\mu$m of 5.24 Jy. We use their sample to provide a context for discussion of the generally unresolved IRAS detections in our sample. Neither the SSM or BGK samples contain any ultraluminous infrared galaxies.   The complexity of the underlying sources of emission in the near infrared has probably been a disincentive to investigating pairs in this spectral range. With current data it remains difficult to disentangle the effects of reddening and thermal dust emission at several temperatures. Photometry at L together with spatially resolved photometry of large samples at longer wavelengths are necessary for clearer understanding of these issues. Currently available data are, however, adequate to demonstrate the promise of the infrared for probing the galaxy-galaxy merger process and the links between galaxy-galaxy interactions and star formation.   Here we compile the largest sample to date of near infrared colors of galaxy pairs. The data we consider include 2MASS JHK$_s$ photometry for 791 galaxies in the BGK sample and our own deeper JHK small aperture photometry for a subsample of 160 BGK galaxies. We also use  IRAS detections and upper limits for the entire BGK sample as a marker of the relationship between galaxy-galaxy interactions, star formation, and thermal emission from hot dust in these systems. We use our optical spectra and infrared colors measured through the same aperture to investigate basic properties of the central regions of these galaxies.  Sec. 2 is a description of the infrared photometry and a review of the optical photometry and spectroscopy. The data include resolved 2MASS photometry for 44 systems unresolved in the public catalogs and for 24 galaxies not detected or confused with nearby stars or galaxies. In Section 3 we investigate the infrared properties of the sample. We discuss the near infrared color-color diagram and investigate the relationships among the near and far infrared colors, the Balmer decrement, and the normalized star formation rate. In Section 4 we show that the star formation rate as indicated by the normalized far infrared flux increases as the projected pair separation decreases as expected if the interaction triggers star formation. There is also a suggestive increase in the  spread of near infrared color as the projected separation of the pair decreases. This increased  spread is evidence for triggered central bursts which may be either very blue or very  heavily obscured and reddened.  We conclude in section 5. ",
        "conclusions": ""
    },
    "0608/astro-ph0608256_arXiv.txt": {
        "abstract": "Hyperfine structure lines of highly-charged ions may open a new window in observations of hot rarefied astrophysical plasmas. In this paper we discuss spectral lines of isotopes and ions abundant at temperatures $10^5-10^7$~K, characteristic for warm-hot intergalactic medium, hot interstellar medium, starburst galaxies, their superwinds and young supernova remnants. Observations of these lines will allow to study bulk and turbulent motions of the observed target and will broaden the information about the gas ionization state, chemical and isotopic composition.  The most prospective is the line of the major nitrogen isotope having wavelength $\\lambda=5.65$~mm~\\citep{SCh84}. Wavelength of this line is well-suited for observation of objects at $z\\approx0.15-0.6$ when it is redshifted to $6.5-9$~mm spectral band widely-used in ground-based radio observations, and, for example, for $z\\ge 1.3$, when the line can be observed in 1.3~cm band and at lower frequencies. Modern and future radio telescopes and interferometers are able to observe the absorption by $^{14}$N~VII in the warm-hot intergalactic medium at redshifts above $z\\approx0.15$ in spectra of brightest mm-band sources. Sub-millimeter emission lines of several most abundant isotopes having hyperfine splitting might also be detected in spectra of young supernova remnants.  \\bigskip PACS: 32.10.Fn, 98.62.Ra, 98.38.-j, 98.58.-w, 32.30.Bv  Keywords: hyperfine structure, intergalactic gas, supernova remnants, radio lines ",
        "introduction": "\\label{SecIntro}  In the temperature range $10^5-10^7$~K the most abundant ion having hyperfine splitting of the ground state is the H-like ion of the major nitrogen isotope $^{14}$N~VII with line in radio range at around 5.64~mm~\\citep{SCh84}. This wavelength of transition between its hyperfine structure (HFS) components has been estimated with high precision to be $\\lambda=5.6519(11)$~mm~\\citep{Shabaev95}. Wavelength of this line is well suited for observation of objects at $z\\approx0.15-0.6$ when it is redshifted to $6.5-9$~mm spectral band widely-used in ground-based radio observations, and, for example, for $z\\ge 1.3$, when the line can be observed in 1.3~cm band and at lower frequencies. Unfortunately, for our Galaxy and its vicinity with $z<0.1$ it is difficult to observe this line due to atmospheric absorption. Possibly, this is the main reason why there have been no attempts to observe it from the ground after this line was reported by~\\citet{SCh84} as a promising candidate for detection of hot rarefied plasmas. Till now astrophysical plasma at such temperatures has been studied only by rocket-based instruments and space ultraviolet and soft X-ray missions such as \\textit{ROSAT}, % \\textit{FUSE}, % \\textit{Suzaku}, % \\textit{Chandra} % and \\textit{XMM-Newton}. %  One of important and interesting predictions of large-scale structure simulations using hydrodynamical approach is the existence of rarefied intergalactic gas heated to temperatures $T\\approx10^5-10^7$~K. This warm-hot intergalactic medium (WHIM, e.g., \\cite{CenWhere1} and \\cite{CroftWHIM,Dave01}) contains dominant fraction of barions in the present Universe (according to the cosmic census of barions made by \\cite{FukugitaBudget}), but is practically unobserved till now. Computations by~\\cite{Gnedin,CenWhere1} show that heavy element abundances in it rise sharply in regions of higher temperature and density reaching values close to the solar ones.  There were many attempts to observe the WHIM at $z>0$, but till now only far-ultraviolet absorption lines of lithium-like oxygen ($\\lambda\\lambda=1032,1038$~\\AA) have been detected by \\textit{Hubble Space Telescope} % for $z>0.12$ and by \\textit{FUSE} spacecraft for $z<0.15$ (e.g.,~\\cite{DanforthOVI}). Despite many attempts using \\textit{Chandra} and \\textit{Suzaku}~\\citep{SuzakuA2218} orbital telescopes, O~VII and O~VIII soft X-ray lines for $z>0$ are evading detection till now (for example, reported case of detection in~\\cite{Nicastro} was recently criticized by~\\cite{NicastroAnti1} and~\\cite{NicastroAnti2}). Realization of planned proposals considering the micro-calorimeter X-ray sky surveys (e.g.,~\\cite{SMEX}) will certainly highly increase the probability to detect soft X-ray emission and absorption lines from the WHIM.  It is obvious that existence of ground-based methods of detection of this gas will bring a lot of additional strength in such efforts. We are proposing radio observations of hyperfine structure lines of highly-charged ions in sub-millimeter to centimeter bands as a way to get additional information about the velocity field, mass, temperature and chemical abundance distribution of the warm-hot intergalactic medium.  It is important to mention that spectral resolution of radio detectors is better than even that of micro-calorimeters. Therefore these lines might permit to look for turbulence and bulk motions in the objects of interest~\\citep{KKHP-SK}. In addition, radio methods do not only probe the absorption lines in spectra of brightest radio sources in millimeter spectral band, but also the corresponding emission of the gas.  None of the lines we are discussing was detected so far neither in astrophysical objects, nor in physical laboratory conditions. Simplest estimates using well-known relativistic corrections and approximations~\\citep{Sobelman1} give precision of the HFS line wavelength on the order of few percents. The list of hyperfine structure lines of interest for study of hot plasmas, published by~\\cite{SCh84}, attracted attention of atomic physicists. \\cite{ZhangSampson97,ZhangSampson00} made computations of electronic excitation of HFS levels in plasmas using relativistic distorted-wave method, accounting for resonance effects. \\cite{Shabaev95,Shabaev97} calculated more precise transition wavelengths using a combination of configuration interaction Hartree-Fock method and the $\\frac1Z$ perturbation theory.  Observations of supernova remnants brightest in soft X-ray band (according to \\textit{ROSAT} All-Sky Survey~\\citep{RASS-BSC} and recent \\textit{XMM-Newton} studies) indicate the most prospective objects to search for HFS line emission. Knowledge of the studied object radial velocity from optical observations will allow to measure wavelengths of the lines with precision of at least 10~km/s (corresponding to wavelength uncertainty of 30~ppm) that is better than the theoretical estimate precision (200~ppm or more). Knowledge of exact wavelengths will also be very important for study and identification of absorption lines from WHIM filaments. Additionally, one must not forget that such theoretical computations should always be confirmed by experimental measurements.  Among such brightest targets we made estimates of brightness of particular HFS lines arising in Cygnus Loop, North Polar Spur, Vela XYZ, N157B and Cas~A supernova remnants (SNR). Our estimates show that in all these objects (except for too hot Cas~A) the line of $^{14}$N~VII is the brightest. Unfortunately, it is subject to atmospheric obscuration.  At Chajnantor plain height ($h\\approx5100$~m, where \\textit{APEX} telescope is operating and \\textit{Atacama Cosmology Telescope} and \\textit{ALMA} interferometer are being built) the air specific attenuation is already significantly less than on the sea level (0.2 vs.\\ 1.1 dB/km at 53~GHz,~\\cite{Liebe60GHz}) and atmospheric transmission of $^{14}$N~VII line may be as high as 30--50\\%~\\citep{SchwabHogg}. This opens a possibility of direct detection of nitrogen radio line from bright supernova remnants of the southern skies and the Cygnus Loop. High-altitude radio telescopes, naturally, are able to observe this line also at much lower redshifts than the instruments at the sea level.  There are many well-studied star-forming galaxies at redshifts $z=0.2-0.5$ which should have a lot of relatively young supernova remnants with gas in interesting range of temperatures. Such SNRs and galaxies as a whole would be extremely interesting objects to study by means of hyperfine structure lines using existing (such as \\textit{Green Bank Telescope} % and \\textit{VLA}) % and future (such as \\textit{Square Kilometer Array} % and \\textit{ALMA}) radio telescopes and interferometers. Hyperfine lines might be bright also in strong outflows of the hot gas from star-forming galaxies (e.g.,~\\cite{GalOutflows}). It is commonly believed that such galactic winds is one of the ways of the intergalactic medium heavy element enrichment.  \\bigskip Hyperfine line observations in principle permit to separate contributions of heliospheric charge-exchange emission and of the Local Hot Bubble, that is now constituting a hardly solvable problem for soft X-ray measurements (e.g.,~\\cite{MBM12Chandra}). This becomes possible due to HFS line emissivity dependence on the plasma density (see below), and the fact that the solar wind is much denser than the gas of the Local Bubble.  The Galaxy is essentially transparent in millimeter band. Therefore the HFS lines allow studying objects that are strongly obscured in visible and soft X-ray bands. Line emission will mostly probe young supernova remnants, but the total absorption column in HFS lines is a measure of hot ($T\\approx10^6$~K) interstellar gas in the Galaxy, thus being complimentary to H~I 21~cm line probing neutral interstellar medium (ISM).  In case of observations of young supernova remnants where the bulk of emission is coming from the enriched material, it becomes possible to measure the isotopic composition by comparing the HFS line intensity (sensitive to just one isotope) with X-ray lines of the same element (produced by all isotopes).  \\bigskip The structure of the paper is the following. In the first part we are presenting wavelengths and other data on the most abundant isotopes in the temperature range $10^5-10^7$~K having hyperfine structure lines in sub-millimeter, millimeter and centimeter bands. We are discussing sub-millimeter lines in connection with \\textit{ALMA} array being able to observe in this spectral band. We are also computing the HFS level population as a function of electron density and radiation temperature for different transitions. In the second part we are presenting the results of the emission line differential brightness temperature computations from some of the brightest objects in our Galaxy and its surroundings. In the last Section we are presenting our estimates of HFS absorption line optical depth arising in WHIM and hot ISM. ",
        "conclusions": "\\label{SecResults}  We discuss the feasibility of the emission and absorption hyperfine structure line observations from astrophysical objects in temperature range of $10^5-10^7$~K. We find that thanks to $^{14}$N high isotopic abundance in interstellar and intergalactic gas the line of $^{14}$N~VII is the most prospective to be observed in absorption in spectra of brightest mm-band extragalactic radio sources with $z>0.15$. Typical optical depth predicted for WHIM filaments is about $10^{-4}$ that is within the reach of the existing instruments. Other HFS lines in absorption will only be detectable with the next generation radio telescopes.  We estimate the $^{14}$N~VII brightness temperature in emission from several Galactic supernova remnants to constitute up to 700~$\\mu$K. Brightness temperature in other lines of $^{25}$Mg, $^{27}$Al, $^{29}$Si and $^{33}$S reaches up to 100~$\\mu$K in oxygen-rich SNR Cas~A, the level on the limit of sensitivity of the existing instruments and accessible to future observatories.  Observations of hyperfine structure lines might provide additional information on isotopic composition of supernova remnants, help to discover obscured Galactic SNRs, allow studying star-forming galaxies and superwinds, discriminating between heliospheric and Local Bubble contribution to diffuse soft X-ray background. They might also allow to measure hyperfine splitting experimentally with high precision, thus letting to choose between theoretical highly-charged heavy ion hyperfine splitting computation models and to precise isotopic nuclear magnetic moments."
    },
    "0608/astro-ph0608706_arXiv.txt": {
        "abstract": "We review the status of indirect Dark Matter searches, focusing in particular on the connection with gamma-ray Astrophysics, and on the prospects for detection with the upcoming space telescope GLAST and Air Cherenkov Telescopes such as CANGAROO, HESS, MAGIC and VERITAS. After a brief introduction where we review the fundamental motivations for indirect searches, we tackle the question of whether it is possible to obtain strong enough evidence from astrophysical observations, to claim discovery of Dark Matter particles. To this purpose, we discuss some recent conflicting claims that have generated some confusion in the field, and present new strategies that may provide the long-awaited smoking-gun for Dark Matter. ",
        "introduction": "\\label{intro}  Dark Matter (DM) is one of the pillars of Modern Cosmology, and cosmological observations provide the most robust (if not the only) {\\it evidence} for physics beyond the Standard Model. The discovery of Dark Matter particles may thus represent at the same time the discovery of new physics, and the identification of one of the most fundamental consituents of the Universe (for reviews see e.g. Refs.~\\cite{Bergstrom:2000pn,Bertone:2004pz}. Of course the best case scenario is that DM particles are found at accelerators, in which case astrophysical observations can be used to prove that the discovered particles have the appropriate cosmological abundance, and to further constrain their properties~\\cite{Hooper:2006wv}.  But what if DM particles are {\\it not} observed at accelerators, what if it takes many years to obtain conclusive answers from, say, the analysis of events at the Large Hadron Collider (LHC)? Is it possible to obtain conclusive enough evidence from astrophysical observations {\\it only}, to actually claim ``discovery'' of DM particles?  DM searches can be broadly divided in three different categories. First, one could search for new particles at accelerators. Much attention has in fact been devoted to the prospects for detecting new physics, and in particular signals of Supersymmetry (SUSY) with accelerators, such as Tevatron and the LHC. If new particles are discovered, it might be even possible, starting from the mass spectrum and particle properties of the new theory, to obtain a tentative estimate of the relic density of the new particles, thus allowing a tentative identification of the DM particle (see e.g. Refs.~\\cite{Tovey:2002jc,Baltz:2006fm} and references therein).  The prospects for discovery at accelerators have been worked out also for an alternative candidate, i.e. the lightest Kaluza--Klein particle (LKP)~\\cite{Servant:2002aq} , in theories with Universal Extra Dimensions (UED)~\\cite{Appelquist:2000nn}. It might be possible already at LHC to discriminate between SUSY and UED, thus to obtain a first hint on the nature of DM~\\cite{Datta:2005zs}. Alternatively, one might search for DM particles {\\it directly}, i.e. my measuring the energy recoil of a nucleus hit by a DM particle streaming through the Earth. More than 20 direct DM detection experiments are either now operating or are currently in development. We refer the interested reader to Refs.~\\cite{Bergstrom:2000pn,Bertone:2004pz} for more details.  Finally, DM can be searched for {\\it indirectly}, i.e. through the detection of its annihilation products such as gamma-rays, neutrinos and anti-matter (e.g. Ref.~\\cite{Bertone:2004pz} and references therein). Although indirect searches are inevitably affected by large astrophysical uncertainties, it is nonetheless possible to obtain useful constraints on the properties of DM, and possibly, as we argue below, to obtain strong enough evidence to claim discovery.  The paper is organized as follows: we tackle some fundamental questions about the motivations for indirect DM searches in Sec.~\\ref{sec:faq}, whose title could be  ``Everything you always wanted to know about DM and never dared to ask''. The experienced reader can skip this section and move directly to Sec.~\\ref{sec:claims}, where we provide a critical discussion of some recent conflicting claims of ``discovery'' of DM, while Sec.~\\ref{sec:smoking} is devoted to the most recently proposed strategies that could eventually provide the smoking-gun for DM annihilations. ",
        "conclusions": "To establish the connection between gamma-ray astrophysics and DM searches, we have first reviewed the fundamental motivations for indirect DM searches. We tried to convey the fundamental message, especially to non-experts, that despite the ``exotic'' nature of the particles under consideration, indirect signals of DM annihilation are a ``natural'' expectation in the standard WIMP paradigm, and that it is certainly worth pursuing further indirect DM searches. We have also provided a critical assessment of some recent ``conflicting claims'' of discovery, recently appeared in literature. Without taking sides in the disputes over specific DM candidates, we have stressed the difficulty of making unambiguous claims, especially in view of the large uncertainties associated with the mass scale, the distribution and the nature of DM. We have finally shown that despite these conflicting claims, the ``status quo'' shouldn't discourage further studies, especially in view of the numerous new strategies that have been recently proposed, and that may provide unambiguous evidence for DM. Among them we have discussed the peculiar power-spectrum of the gamma-ray background from DM annihilations; the multi-wavelength, multi-messenger approach; and the so-called ``mini-spike'' scenario, where IMBHs act as ``boosters'' of the DM annihilation signal."
    },
    "0608/astro-ph0608195_arXiv.txt": {
        "abstract": "By performing magnetohydrodynamical (MHD) simulations, we investigate the mass loss of intermediate- and low-mass stars from main sequence (MS) to red giant branch (RGB) phases. \\Alfven waves, which are excited by % the surface convections travel outwardly and dissipate by nonlinear processes to accelerate and heat the stellar winds. We dynamically treat these processes in open magnetic field regions from the photospheres to $\\simeq$25 stellar radii. When the stars evolve to slightly blueward positions of the dividing line (Linsky \\& Haisch), the steady hot corona with temperature, $T \\approx 10^6$ K, suddenly disappears. Instead, many hot ($\\sim 10^6$ K) and warm ($\\gtrsim 10^5$ K) bubbles are formed in cool ($T\\lesssim 2\\times 10^4$ K) chromospheric winds because of thermal instability; the red giant wind is not a steady stream but structured outflow. As a result, the mass loss rates, $\\dot{M}$, largely vary in time by 3-4 orders or magnitude in the RGB stars. Supported by magnetic pressure, the density of hot bubbles can be kept low to reduce the radiative cooling and to maintain the high temperature long time. Even in the stars redward of the dividing line, hot bubbles intermittently exist, and they can be sources of ultraviolet/soft X-ray emissions from hybrid stars. Nearly static regions are formed above the photospheres of the RGB stars, and the stellar winds are effectively accelerated from several stellar radii. Then, the wind velocity is much smaller than the surface escape speed, because it is regulated by the slower escape speed at that location. We finally derive an equation that determines $\\dot{M}$ from the energetics of the simulated wave-driven winds in a {\\it forward} manner. The relation explains $\\dot{M}$ from MS to RGB, and it can play a complementary role to the Reimers' formula, which is mainly for more luminous gaints. ",
        "introduction": "Solar-type stars have X-ray emitting coronae with the temperatures, $T\\gtrsim 10^6$ K, and hot stellar winds stream out with the asymptotic terminal velocities, $v_{\\infty}=100-1000$ km s$^{-1}$, comparable to the escape velocities at the surfaces. The total mass lost by the wind is negligible during the main sequence (MS) phase in comparison with the stellar mass, whereas the wind may play a role in the loss of the angular momentums. Along with the stellar evolution to red giant branch (RGB), the mass loss rate, $\\dot{M}$'s, increases by 4-10 orders of magnitude \\citep{js91}, while $v_{\\infty}$ drops to $10-100$ km s$^{-1}$, which is smaller than the surface escape velocity.  In addition, the temperature of the wind also decreases, which is probably drastic across the ``dividing line'' (DL hereafter; Linsky \\& Haisch 1979) discussed below. The mass loss during the RGB phase becomes more important, and it may control the later evolution of a star itself. For instance, some stars evolve to a peculiar group of blue horizontal branch stars, and it is inferred that it is related to the mass loss during the RGB phase \\citep{ydk97}. In spite of its importance, however, the acceleration mechanism of red giant winds is not well understood.     It was reported that cool single RGB stars redward of the DL (near spectral type K1 III in Hertzsprung-Russel (HR)-diagram; Figure \\ref{fig:hr}) showed an evidence of neither ultraviolet (UV) radiation of C$_{\\rm IV}$ ($T\\sim 10^5$ K) \\citep{lh79} nor soft X-ray ($T\\sim 10^6$ K) \\citep{ayr81}. \\citet{sm80} further pointed out that the stellar winds appear cool and massive after stars cross the DL. However, \\citet{hdr80} observed late-type super-giants in the noncoronal zone that exhibit both UV emission and blue-shifted chromospheric absorption features, indicating that hot gas and cool outflow coexist. Such hybrid stars were also detected % by later observations (e.g. Ayres et al. 1998), whereas the activity levels of UV/soft X-ray are much lower than in the coronal zone. These observations show that the simple DL picture roughly holds while the reality seems a bit more complicated.   RGB stars are generally slow rotators. The radiation luminosities of intermediate- and low-mass stars near the DL are not still large, either. Therefore, the centrifugal force and radiation pressure are not sufficient in accelerating the stellar winds (e.g. Judge \\& Stencel 1991). Instead, the surface convective layer is expected to be the main origin of driving the winds and UV/X-ray activities. Magnetic fields are probably generated by turbulent dynamo in the convection zone. The energy of the surface turbulences is lifted up through the magnetic fields, and the dissipation of the energy in upper regions leads to the heating of the atmosphere and the acceleration of the wind. This picture is essentially the same as what takes place in the solar corona and wind. Difference from the sun is that the surface gravity is much smaller in RGB stars. In a na\\\"{i}ve way of thinking, the transition from the steady coronae to the cool winds (with some hybrid activities) is possibly explained by this gravity effect even though other conditions are similar; the cool atmosphere comes to flow out before heated up to coronal temperature ($T\\sim 10^6$ K), as the gravity confinement weakens.   Many pioneering works have tried to reveal nature of red giant winds (e.g. Hartmann \\& MacGregor 1980; Holzer, Fl\\aa, \\& Leer 1983; Charbboneau \\& MacGregor 1995). However, the consecutive processes from the surface to the wind explained above have not been studied self-consistently; previous model calculations % started from locations far above the photosphere in order to avoid difficulties arising from the huge density contrast. Moreover, they assume mechanical or phenomenological ``wave'' energies with somewhat ad hoc prescription of the dissipation.  In this paper, we perform dynamical simulations of stellar winds from the photosphere to $\\simeq 25\\; R_{\\star}$ in open magnetic field regions, where $R_{\\star}$ is stellar radius. We give perturbations at the photosphere originating from the surface convection. These fluctuations excite waves that propagate upwardly. Among various modes the \\Alfven wave travels to a sufficiently large distance to contribute to the acceleration and heating of the stellar wind. We simulate the propagation and dissipation of the waves by solving nonlinear MHD equations, and treat the heating and acceleration of the gas in a fully self-consistent manner without conventional heating functions (\\S\\ref{sec:dwv}). Wind properties, such as $\\dot{M}$ and $v_{\\infty}$, are determined from the surface conditions without ad hoc parameters. Another advantage of dynamical simulations is that we do not have to care about the critical (sonic) point(s), which is often problematic when constructing a transonic wind solution under the steady-state condition. A stable wind structure is automatically selected whether it is transonic or subsonic.   We already carried out simulations in the solar corona and wind \\citep[][hereafter SI05; SI06]{si05,si06}. We showed that the coronal heating and solar wind acceleration in the open field regions are natural consequences of the photospheric perturbations; the \\Alfven waves generated from the surface effectively dissipate in the corona by nonlinear mechanism to sufficiently heat and accelerate the solar wind in spite of the energy loss via radiative cooling, thermal conduction, and wave reflection. The aim of this paper is to study the evolution of the stellar winds from MS to RGB stages by extending this work to cool RGB stars. We investigate the stellar winds of a 1$M_{\\odot}$ star from MS ({\\it i.e.} present sun) to RGB across the DL, where $M_{\\odot}$ is the solar mass. For comparison we also study a more massive RGB star with $3M_{\\odot}$ both blueward and redward of the DL.  \\S\\ref{sec:mod} presents our simulation method. In \\S\\ref{sec:esw} we show the results of the simulations, focusing on the evolution and time-dependent nature of the stellar winds. We discuss comparisons with related works as well as the uncertainties and limitations of our simulations in \\S\\ref{sec:dis}. ",
        "conclusions": "\\label{sec:dis} \\subsection{Magnetic Fields} \\label{sec:unc} When performing the simulations, we have to set up in advance the surface fluctuations and the magnetic flux tubes. While we can reasonably estimate the amplitude and period (spectrum) of the surface fluctuations from the conditions of the convection zones (\\S \\ref{sec:phfl}), there are large uncertainties in the properties of the magnetic fields.   We only consider the open magnetic field regions because our aim is to study the stellar winds. We assume the same magnetic field strength ($B_{r,0}=240$ G) at the surface and expansion factor ($f_{\\rm tot}=240$) in the different stars. Although it is unclear how in reality the open magnetic flux tubes are, we had better consider effects of different flux tubes (see \\citet{fvj06} for configurations of flux tubes).  We have actually studied the effects of the strength and configuration of magnetic flux tubes on the solar wind by both numerical simulations (SI06) and analytical method \\citep{suz06}. We showed that $B_{r,0}/f_{\\rm tot}$ is a more important parameter in controling solar wind properties than individual $B_{r,0}$ and $f_{\\rm tot}$, if nonlinear \\Alfven waves dominantly work in the acceleration. This is also consistent with recent observation by \\citet{kfh05}. For example, flux tubes with different $B_{r,0}$ and $f_{\\rm tot}$ but the same ratio $B_{r,0}/f_{\\rm tot}$ give similar wind properties if the photospheric amplitudes are the same. An input energy flux is $\\sim \\rho \\langle \\delta v_0 \\rangle^2 v_{A,0} \\propto B_{r,0}$ at the surface, and it is diluted according to $1/fr^2$ in an expanding flux tube. The acceleration of the stellar wind mainly occurs in the outer region where the super-radial expansion of the flux tube already finishes. Therefore, $B_{r,0}/f_{\\rm tot}$, which is proportional to $S_c$ ($\\approx$energy flux) of the input \\Alfven wave {\\it measured at the outer region}, becomes an important control parameter in the \\Alfven wave-driven wind. Moreover, a mass flux, {\\it i.e.} $\\dot{M}$, is determined almost solely by $\\langle \\delta v_0 \\rangle^2$, % and very weakly dependent on $B_{r,0}/f_{\\rm tot}$ (SI06). Then, $B_{r,0}/f_{\\rm tot}$ plays a role only in tuning the terminal velocity of the solar wind.   This argument is directly applicable to red giant winds. $\\dot{M}$'s of the stellar winds (the top panel of Figure \\ref{fig:chp}) are determined with little uncertainty since they are mainly determined by $\\langle \\delta v_0 \\rangle^2$ of the surface convections. However, when we discuss the wind terminal speeds, $B_{r,0}/f_{\\rm tot}$ needs to be taken carefully; wind speed becomes faster in a flux tube with larger $B_{r,0}/f_{\\rm tot}$.  $B_{r,0}/f_{\\rm tot}(=240/240=1\\; {\\rm G})$ corresponds to the total magnetic flux over the whole open field regions divided by the entire stellar surface. To discuss the average field strength we have to take into account closed regions, which occupy a sizable fraction of the surface. If 10\\% of the surface is occupied by open field regions, the average strength becomes 10 times, namely 10G. This value is quite reasonable from an energetics point of view. The magnetic energy density at the surface is $B_{r,0}^2/8\\pi \\simeq 4$ erg cm$^{-3}$, which is considerably smaller than the kinetic energy of the surface turbulence, $\\rho \\langle \\delta v_0^2 \\rangle \\sim 10^{3} - 10^4$ erg cm$^{-3}$, of the simulated stars. It is sufficient only if $\\sim 0.1$ \\% of the turbulent energy is converted to the magnetic energy by motions of charged particles ({\\it i.e.} current) in the surface convection.  Related to this, \\citet{ros91,ros95} tried to interpret the DL by the change of magnetic topologies. The stellar rotation becomes slow with the evolution so that the generation of magnetic fields by large-scale dynamo is switched-off. Magnetic fields are generated only by (small-scale) turbulent convection in the surface layers in evolved stars. As a result, % closed fields, which have dominantly covered the surface in unevolved stars, become open in evolved stars redward of the DL. Therefore, hot plasma, which has been well confined by closed loops, would flow out.  If this scenario is correct\\footnote{It should be noted that \\citet{ahb03} recently proposed a different scenario for the DL: even though closed loops, which are sources of X-rays, are distributed in stars redward of the DL, these X-rays are obscured by larger column density of above chromospheric material. This density effect is a reasonable explanation of the DL.}, surface field strength, $B_{r,0}$, becomes smaller in evolved stars. However, we expect that the field strength in the outer regions ($\\propto B_{r,0}/f_{\\rm tot}$) is less affected, because an decrease of $B_{r,0}$ firstly reduces a fraction of closed structures rather than strength of open fields in the outer region;  % even if $B_{r,0}$ decreases, $B_0/f_{\\rm tot}$ stays more or less constant before all the field lines become open. Because $B_{r,0}/f_{\\rm tot}$ is more important than $B_{r,0}$ as stated above, we anticipate that our results of the stellar winds are affected little even if the magnetic field topologies change with the stellar evolution.       \\subsection{Previous Wind Calculations} \\citet{hm80} investigated stellar winds from late-type stars with explicitly taking into account the effects of \\Alfven waves. They calculated several stars with different wave energies and dissipation parameters under the steady-state condition from assumed inner boundaries. They concluded that the dissipation of \\Alfven waves resulted in coronal heating in higher gravity stars (MS and sub-giant) and mass loss in lower gravity stars (RGB). Our analysis of the time-average wind structures basically confirm this tendency, % whereas our simulations have revealed the various important time-dependent aspects (\\S\\ref{sec:tdp}).  Most of previous works claim that it is difficult to accelerate slow and dense red giant winds by \\Alfven waves \\citep{hfl,cm95}\\footnote{Although these works focus on more massive RGB stars than ours, the result will not change if less massive stars are considered. } The main problem is that a sizable amount of the input wave energy remains in the supersonic region, which leads to faster wind. There is also a similar problem in the concept of a supersonic transition locus, which \\citet{mul78} introduced to explain the onset of massive wind. The main difference of our results from the previous works is that the large subsonic regions are formed in our simulations due to the nearly static regions above the surfaces. Therefore, most of the input \\Alfven wave energies are deposited in the subsonic regions, which results in the slow winds, although the dissipation is not much faster than in the previous model calculations.  We speculate that our dynamical wind solutions % correspond to ``inner solutions'' introduced by \\citet{hfl}. They pointed out that there are two transonic solutions in some cases; one is the inner solution in which the wind is accelerated very gradually and the wind speed is slower, although they mainly studied the other one, the outer solution, in which the acceleration starts just above the surface. Our dynamical simulations automatically select stable transonic wind solutions. Then, our results might show that the inner solution is the stable branch of \\Alfven wave-driven winds.  We would like to remark on an issue of time-dependent winds. Our simulations show that some portions of the RGB atmospheres transiently accrete ($v_r<0$), which lasts longer than the assumed longest wave periods ({\\it e.g.} see the dashed line in the top panel of Figure \\ref{fig:mpw}). This phenomenon may be an example of the hysterisis transition between subsonic accretion and supersonic wind discovered by \\citet{vel94,dvl98}.    \\subsection{Limitations} \\label{sec:lim} Here we discuss processes and effects which need to be studied in more detailed calculations.  \\subsubsection{Heating of Low Chromosphere: Sound Waves?} We do not input longitudinal fluctuations at the photospheres, which would directly drive sound waves, because they do not contribute to the heating and acceleration of the stellar winds ({\\it e.g.} Judge \\& Carpenter 1998) (In our simulations sound waves are generated from \\Alfven waves in the upper regions ; \\S \\ref{sec:dwv}). However, the sound waves directly from the surface might be important in the heating of low chromosphere (e.g. Carlsson \\& Stein 1997). We switch off the radiative cooling in the inner regions if $T< 3000$ K and $\\rho > 5\\times 10^{-17}$ g cm$^{-3}$ by the technical reason (\\S\\ref{sec:siml}). By taking into account the longitudinal perturbations at the photospheres, we may carry out a more self-consistent treatment of the low chromospheres without switching off the radiative cooling artificially.   \\subsubsection{Multi-dimensional Effects -Wave Generation-} \\label{sec:clp} In our simulations we only consider the wave generation at the photospheres in the 1D flux tubes. However, interactions between flux tubes may be important in generation of waves, though 2D or 3D modeling is required to study such processes. For example, reconnection events between closed loops could be sources of small flare-like events, and these events probably excite compressive waves at locations above the photosphere \\citep{str99}.  The magnetic reconnection between a closed loop and an open field line also possibly excites \\Alfven waves traveling upwardly \\citep{miy05}. These waves would also play a role in the acceleration of red giant winds. Further detailed analysis is necessary to estimate how such waves produced by activities involving closed loops are important with respect to the energetics.  \\subsubsection{Multi-dimensional Effects -Wave Propagation-} We treat the propagation of \\Alfven waves in the 1D flux tubes. In other words, \\Alfven waves are assumed to travel in the same way in neighboring flux tubes. However, if \\Alfven speeds are slightly different in neighboring field lines, the waves become out of phase even if they are generated in the same way at the surface. Then, these \\Alfven waves dissipate by phase mixing through resistivity and viscosity \\citep{hp83}.  \\Alfven waves in stellar winds become more or less turbulent-like. Then, \\Alfvenic turbulence cascades to higher frequency in the transverse direction with respect to underlying magnetic fields \\citep{gs95}. Therefore, \\Alfven waves might dissipate by different channels in more realistic 3D calculations, while in our 1D simulations the \\Alfven waves dissipate mainly through the decay instability (\\S\\ref{sec:dwv}).    \\subsubsection{Collisionless Effects} We assume the MHD approximation, namely a mean free path with respect to Coulomb collisions, $l_{\\rm mfp} = 9.38\\times 10^7\\;{\\rm cm} \\frac{(T/10^6\\;{\\rm K})^2}{n/10^8\\;{\\rm cm^{-3}}},$ is smaller than a typical scale. Here, we can take the minimum wavelength, $\\lambda_{\\rm min}\\simeq v_{\\rm A}\\tau_{\\rm min}$, of the \\Alfven waves as a typical scale, hence, \\begin{equation} \\lambda_{\\rm min} > l_{\\rm mfp} \\end{equation} is the requirement for the MHD condition. Otherwise we have to consider collisionless effects of plasma. This condition breaks in $r\\gtrsim 2\\;R_{\\star}$ of the MS star (Model I) and in $\\gtrsim 5\\;R_{\\star}$ of the sub-giant star (Model II), while in the RGB stars (Models III - VI) it holds in the entire simulation regions because the densities are high and the temperatures are low.  In the collisionless regime, compressive (slow/fast MHD) waves suffer collisionless damping % \\citep{bar66, shl06}. Then, the dissipation rate of the slow MHD waves, which are nonlinearly generated from the \\Alfven waves, may be changed. This might further modify the mode conversion rate from the \\Alfven waves. Then, the dissipation of the \\Alfven waves might be indirectly affected in Models I and II.  \\subsubsection{Atoms, Molecules, and Dusts} \\label{sec:pion} In the RGB atmospheres, the temperatures sometimes become $10^3 - 10^4$ K as shown in the snap-shot results (Figures \\ref{fig:mpw} \\& \\ref{fig:tdpd1M}). Although even in such low temperatures a fraction of the gas is expected to be ionized by the UV/X-ray emissions from hot bubbles, a sizable amount of neutral atoms remain, or even molecules can also exist if $T\\lesssim 3000$ K \\citep{tsu73}.  The snap-shot results also show that the temperatures partly decreases further to $\\sim 1000$ K, and in such circumstances dusts are also formed.  The existence of atoms and molecules indicate that the gas is not fully ionized but partially ionized. The MHD approximation implicitly assumes that the coupling between ions and neutrals is sufficient so that the gas can be treated as one fluid. Otherwise, the friction between these two components plays a role in the dissipation of \\Alfven waves. We can check the condition of ion-neutral coupling by comparing wave frequency and a ion-neutral collision rate \\citep{drd83}. We have found that the coupling is sufficient with respect to the low-frequency waves we are dealing with; we do not have to worry about the effect of weakly ionized gas. However, if high frequency \\Alfven waves are generated by turbulent cascade, these waves dissipate by the ion-neutral damping more effectively than in fully-ionized gas \\citep{hfl}.   Molecules can be used as diagnostics to understand physical conditions of RGB star atmospheres. Observations of CO show that quasi-static cool ($T\\sim 1000 -2000$ K) regions containing molecules are in the atmospheres of RGB stars, typically at $\\sim 2R_{\\star}$ \\citep{tsu88, tsu97}. Our simulations of the RGB stars also show the quasi-static regions in the atmospheres and the temperatures partly becomes low enough to form molecules. The cooling by molecules drive thermal instability at $T=2000 - 3500$ K \\citep{mns87}, similarly to the ions in $T\\gtrsim 10^5$ K. This will be important in controlling thermal properties of the atmospheres and winds of RGB stars.    Radiation pressure on dusts is not efficient in driving winds of RGB stars \\citep{js91,vj06}, although it could be important in asymptotic giant branch stars. However, \\Alfven waves may dissipate by the cyclotron resonance with charged dusts \\citep{vj06}, which might also give a correction to the dissipation of the \\Alfven waves. To study this effect, we have to carefully consider the formation processes of dusts in intermittently existing cool ($T\\sim 1000$ K) gas."
    },
    "0608/astro-ph0608476_arXiv.txt": {
        "abstract": "The recent discovery of high frequency oscillations during giant flares from the Soft Gamma Repeaters SGR 1806-20 and SGR 1900+14 may be the first direct detection of vibrations in a neutron star crust.  If this interpretation is correct it offers a novel means of testing the neutron star equation of state, crustal breaking strain, and magnetic field configuration. We review the observational data on the magnetar oscillations, including new timing analysis of the SGR 1806-20 giant flare using data from the {\\it Ramaty High Energy Solar Spectroscopic Imager} and the {\\it Rossi X-ray Timing Explorer}.  We discuss the implications for the study of neutron star structure and crust thickness, and outline areas for future investigation. ",
        "introduction": "\\label{intro} The Soft Gamma Repeaters (SGRs) are thought to be magnetars, neutron stars with magnetic fields in excess of $10^{14}$G \\citep{duncan1992, thompson1995}. Decay of the strong field powers regular gamma-ray flaring activity that culminates, on rare occasions, in a giant flare with a peak luminosity in the range $10^{44}$- $10^{46}$ erg~s$^{-1}$.  The three giant flares detected to date consist of a short, spectrally hard initial peak, followed by a softer decaying tail that lasts for several hundred seconds.  Pulsations with periods of a few seconds are visible in the tail and reveal the neutron star spin.  Their presence is thought to be due to a fireball of ejected plasma, trapped near the stellar surface by the strong magnetic field \\citep{thompson1995}.  Powering the giant flares requires a catastrophic global reconfiguration of the magnetic field.  The coupling between the field and the charged particles in the neutron star crust means that this is likely to be associated with large-scale crust fracturing \\citep{flowers1977, thompson1995, thompson2001, schwartz2005}.  This in turn should excite global seismic vibrations:  on Earth seismologists regularly observe global modes after large earthquakes (see for example \\citealt{park2005}).  In the SGR case, the coupling of field and crust should cause the modes to modulate the X-ray lightcurve.  Various different types of oscillation are possible, but theory suggests that the easiest to excite and observe should be the toroidal shear modes of the crust \\citep{blaes1989}.  The particular harmonics excited will depend on the size, shape and speed of the fracture.  The observed mode frequencies depend on the neutron star mass and radius (via gravitational redshift and the influence of the equation of state on crust structure), crustal composition, and magnetic field strength and configuration \\citep{hansen1980, schumaker1983, mcdermott1988, strohmayer1991, duncan1998, messios2001, piro2005}.  Detection and identification of crustal modes would therefore probe all of these areas of neutron star physics.  In 2004, SGR 1806-20 emitted the most powerful flare ever recorded.  In a landmark paper, \\cite{israel2005} reported the detection of Quasi-Periodic Oscillations (QPOs) in the 18-93 Hz range in the decaying tail of the flare.  Prompted by this result, \\cite{strohmayer2005} re-analysed the 1998 giant flare from SGR 1900+14 and found QPOs in the range 28-155 Hz. In both cases tentative identifications can be made with a sequence of toroidal modes.  More in-depth analysis of the SGR 1806-20 flare has since revealed additional QPOs, including the first evidence for a higher radial overtone \\citep{watts2006, strohmayer2006}.  We review all of the observational results in Section \\ref{sec:1}.   In Section \\ref{sec:2} we discuss the results in the context of the toroidal mode models, and show how they can be used to constrain neutron star parameters including the crust thickness.  We conclude in Sections \\ref{sec:3} and \\ref{sec:4} with a discussion of outstanding issues. ",
        "conclusions": "\\label{sec:4}  The discovery of high frequency oscillations during giant flares from the Soft Gamma Repeaters SGR 1806-20 and SGR 1900+14 may be the first direct detection of vibrations in a neutron star crust.  If this interpretation is correct it offers a new means of testing the neutron star equation of state, crustal breaking strain, and magnetic field configuration.  In particular, if the mode interpretation is secure, it allows us to make the first direct estimate of the thickness of a neutron star crust.  This is particularly impressive when one considers the fact that all of the data acquired so far have been obtained by chance using satellites that were observing other objects.  If a rapid-slew instrument such as SWIFT could be configured to record high time resolution data throughout the tails of giant flares from known SGRs, the potential for additional discoveries is immense."
    },
    "0608/astro-ph0608640_arXiv.txt": {
        "abstract": "The Palomar High-precision Astrometric Search for Exoplanet Systems (PHASES) has monitored 37 sub-arcsecond binary systems to determine whether their Keplerian orbits are perturbed by faint astrometric companions to either star.  Software has been developed to evaluate the regions in a companion mass-period phase space in which the PHASES observations can exclude the possibility of face-on orbit perturbations.  We present results for 8 systems for which astrometric companions with masses as small as those of giant planets can be excluded. ",
        "introduction": "A technique has been developed to obtain high precision (10-20 $\\microas$) astrometry of close stellar pairs \\citep[separation less than one arcsecond;][]{LaneMute2004a} using long-baseline infrared interferometry at the Palomar Testbed Interferometer \\citep[PTI;][]{col99}. These observations provide precision visual orbits of the binaries and allow detection of tertiary components orbiting either the primary or secondary due to the reflex motion of the subsystem center-of-light.  These measurements were made at PTI as part of the Palomar High-precision Astrometric Search for Exoplanet Systems (PHASES) program. PTI is located on Palomar Mountain near San Diego, CA. It was developed by the Jet Propulsion Laboratory, California Institute of Technology for NASA, as a testbed for interferometric techniques applicable to the Keck Interferometer and other missions such as the Space Interferometry Mission (SIM).  It operates in the J ($1.2 \\mu{\\rm m}$), H ($1.6 \\mu{\\rm m}$), and K ($2.2 \\mu{\\rm m}$) bands, and combines starlight from two out of three available 40-cm apertures. The apertures form a triangle with one 110 and two 87 meter baselines.  PHASES observations use two of the three available baselines at PTI: the NS (110 meter) and SW (87 meter) baselines.  PTI has previously contributed to exoplanet search efforts through high spatial resolution visibility ``imaging'' of the 51 Pegasi system\\citep{boden1998}, which placed limits on the luminosity of its companion and support the conclusion that it is substellar, and likely planetary, in nature.  Here the initial detection limits of a different effort at PTI are presented, this time based on astrometric measurements of binary systems. ",
        "conclusions": "PHASES observations are able to exclude tertiary companions with masses as small as a few Jupiter masses in several binary systems.  The phase-space searching algorithm correctly identifies the natures of two previously studied triple star systems and finds the correct periods for the short period pairs. While the number of systems that can currently be examined is too small to make strong conclusions about the frequency of giant planets in close binary systems, by the end of the PHASES program enough measurements will be made on all target systems to address this important question.  Systematic effects currently limit the observed precision; if overcome, companions as small as a fifth of a Jupiter mass might be detected in favorable systems such as HR 7162."
    },
    "0608/astro-ph0608689_arXiv.txt": {
        "abstract": "The aim of this article is to draw attention to the importance of the electric current loss in the energy output of radio pulsars. We remind that even the losses attributed to the magneto-dipole radiation of a pulsar in vacuum can be written as a result of an Ampere force action of the electric currens flowing over the neutron star surface~\\citep{Mich,BeskGurIst-93}. It is this force that is responsible for the transfer of angular momentum of a neutron star to an outgoing magneto-dipole wave. If a pulsar is surrounded by plasma, and there is no longitudinal current in its magnetosphere, there is no energy loss~\\citep{BeskGurIst-93,MPS-99}. It is the longitudinal current closing within the pulsar polar cap that exerts the retardation torque acting on the neutron star. This torque can be determined if the structure of longitudinal current is known. Here we remind of the solution by Beskin, Gurevitch \\& Istomin (1993) and discuss the validity of such an assumption. The behavior of the recently observed \"part-time job\" pulsar B1931+24 can be naturally explained within the model of current loss while the magneto-dipole model faces difficulties. ",
        "introduction": "\\label{sec:1}  The first idea to explain the energy loss of radio pulsars was to consider the magneto-dipole radiation~\\citep{Pacini-67}. Indeed, the magneto-dipole formula gives for the radiation power \\begin{equation} W_{{\\mathrm md}} =\\frac{1}{6}\\frac{B_{0}^{2}\\Omega^{4}R^{6}}{c^{3}}\\sin^{2}{\\chi}, \\label{MD-0} \\end{equation} where $\\chi$ is the angle between rotational and magnetic axis, $R$ is a neutron star radius $\\sim 10$ km, and $\\Omega$ is a pulsar angular velocity. This formula explains pulsar activity and observed energy loss for expected large magnetic field near the surface $B_{0}\\sim 10^{12}$ Gs.  Let us recall that the physical reason of such energy loss is the action of the torque exerted on the pulsar by the Ampere force of the electric currens flowing over the neutron star surface \\citep{Ist-05}. The electric and magnetic fields in the outgoing magneto-dipole wave in vacuum can be found by solving the wave equations $\\nabla^2\\mathbf{B}+{\\Omega^{2}}/{c^{2}}\\mathbf{B}=0$ and $\\nabla^2\\mathbf{E}+{\\Omega^{2}}/{c^{2}}\\mathbf{E}=0$ with the boundary conditions stated as the fields corresponding components $\\mathbf{E}_{\\mathrm t}$ and $\\mathbf{B}_{\\mathrm n}$ being continuous through the neutron star surface. Inside the star one can consider magnetic field as homogeneous, and find the corresponding electric field using the frozen-in condition. As a result, the full solution will give us the discontinuity of $\\{\\mathbf{B}_{\\mathrm t}\\}$ and $\\{\\mathbf{E}_{\\mathrm n}\\}$ attributed to the surface charge $\\sigma_{\\mathrm s}$ and the surface current $\\displaystyle\\mathbf{J}_{\\mathrm s}=\\frac{c}{4\\pi}[\\mathbf{n},\\{\\mathbf{B}\\}]$. The Ampere force exerts the torque \\begin{equation} \\mathbf{K}=\\frac{1}{c}\\int{[\\mathbf{r},[\\mathbf{J}_{\\mathrm s},\\mathbf{B}]]dS} \\end{equation} on the neutron star. The energy loss of a pulsar due to this torque is equal to (\\ref{MD-0}). Thus, it is the surface current that is responsible for the angular momentum transform from a neutron star to an outgoing magneto-dipole wave~\\citep{Mich,BeskGurIst-93}.  A pulsar in vacuum loses its rotational energy due to angular momentum transform to the electromagnetic wave at the rate given by (\\ref{MD-0}). However, this is not so if the pulsar magnetosphere is filled with plasma and there is no longitudinal current in the magnetosphere. As was shown by \\cite{BeskGurIst-93} and \\cite{MPS-99}, in this case the Poynting flux through the light cylinder is equal to zero. Indeed, as the ideal conductivity condition is applicable not only inside the neutron star but outside as well there is no magnetic field discontinuity at the star surface. Consequently, there is no Ampere force acting on a pulsar and, hence, there is no energy loss. For zero longitudinal current the light cylinder is a natural boundary of the pulsar magnetosphere. ",
        "conclusions": "As we have seen, the current loss connected with the longitudinal current flowing in the magnetosphere plays the main role in pulsar dynamics, and recent observations of \"part-time job\" pulsar supports this point. This evolution includes not only a neutron star retardation but also the sufficient change in the angle between magnetic and spin axis. We have seen as well that the model of current loss depends crucially on the distribution of the electric current and its value in the inner gap. For current loss model with $j_{\\rm GJ}$ the inclination angle grows with time so a pulsar tends to be an orthogonal rotator. In this case the energy loss is to be $\\Omega R/c$ times smaller than for the aligned rotator. As a consequence, the light surface must be located in the very vicinity of the light cylinder. On the other hand, to realize the homogeneous MHD outflow up to infinity for the orthogonal rotator the current density in the open field line region is to be much larger than $j_{\\rm GJ}$."
    },
    "0608/astro-ph0608555_arXiv.txt": {
        "abstract": "We study the optical--UV/X--ray spectral energy distribution of GRB 060218 during the prompt phase and during what seems to be the afterglow phase. The results are puzzling, since if the opt--UV and the X--ray emission belong to a single \\bbp, then its luminosity is too large, and this \\bb cannot be interpreted as the signature of the shock breakout of the supernova. There are also serious problems in associating the emission expected by the supernova shock breakout with either the opt--UV or the X--ray emission. In the former case we derive too small ejecta velocities; in the latter case, on the contrary, the required velocity is too large, corresponding to the large radius of a \\bb required to peak close to the UV band. We then present what we think is the most conservative alternative explanation, namely a synchrotron spectrum, self--absorbed in the opt--UV and extending up to the X--ray band, where we observe the emission of the most energetic electrons, which are responsible for the exponential roll--over of the spectrum. The obtained fit can explain the entire spectrum except the \\bb observed in the X--rays, which must be a separate component. The puzzling feature of this interpretation is that the same model is required to explain the spectrum also at later times, up to $10^5$ s, because the opt--UV emission remains constant in shape and also (approximately) in normalisation. In this case the observed X--ray flux is produced by self--Compton emission. Thus the prompt emission phase should last for $\\sim 10^5$ s or more. Finally, we show that the \\bb observed in X--rays, up to 7000 seconds, can be  photospheric emission from the cocoon or stellar material, energized by the GRB jet at radii comparable to the stellar radius (i.e. $10^{10}$--$10^{11}$ cm), not very far from where this material becomes transparent (e.g. $10^{12}$ cm). ",
        "introduction": "The burst exploded February 18, 2006 is a low redshift burst ($z=0.033$, Mirabal et al. 2004), associated to the Type Ibc supernova SN2006aj (e.g. Modjaz et al. 2006, Mazzali et al. 2006). Due to its long duration (more than 3000 s) GRB 060218 could be followed simultaneously with the BAT, XRT and UVOT instruments onboard the SWIFT satellite (Gehrels et al. 2004). The 0.3--10 keV X--ray light curve (followed by XRT starting 157 s after the BAT trigger) presents three main phases (Campana et al. 2006, hereafter C06): phase X1: a smooth/long ($\\sim 3\\times 10^3$ s) bump peaking at $\\sim 10^3$ s whose time--integrated spectrum shows a non--thermal component (power-law with an exponential cutoff) peaking (in $\\nu F_\\nu$) in the X-ray band ($E_{\\rm peak}\\sim 5$ keV) and a \\bb which comprises about 20\\% of the total 0.3--10 keV flux and dominates the soft X-ray energy band ($kT\\sim 0.13$ keV); X2: a steep power--law (or exponential) time--decay up to 10$^4$ s still showing a slightly softer ($kT\\sim 0.1$ keV) \\bb component (comprising about 50\\% of the total flux) together with a softer non--thermal component; X3: a shallower flux decay ($\\propto t^{-1.2}$) starting at 10$^4$ s, and lasting up to several days, with a very soft (energy spectral index $\\alpha\\sim 2.3$; Cusumano et al. 2006) non--thermal component. The optical--UV light curve presents 2 phases: in phase UV1 there is a slow increase of the flux, peaking at $\\sim 8\\times10^4$ s, followed by a fast decay up to $\\sim 1.5\\times 10^5$ s; in phase UV2 there is a second bump peaking at $\\sim$10 days showing the typical spectral signatures of the underlying supernova (Ferrero et al. 2006; Mirabal et al. 2006; Modjaz et al. 2006; Sollerman et al. 2006) and suggesting photospheric expansion velocities of 2$\\times 10^4$ km s$^{-1}$ (e.g. Pian et al. 2006). Finally, in the radio band, the flux between 2 and 22 days shows a typical power law decay ($\\propto t^{-0.8}$, Soderberg et al. 2006).  Despite the available wealth of information, some of the observed properties of GRB 060218 are not yet understood. In fact, although the radio flux could be due to external shocks, the radio spectrum at 5 days is inconsistent with the strong X--ray emission at the same epoch (Soderberg et al. 2006). Moreover, the typical late time X--ray light curve decay (phase X3) is hardly reconcilable with the extremely soft spectrum in the framework of the external shock model for GRB afterglows. This suggests that the late time X--ray emission might be produced by a continued activity of the central engine (i.e. ``central engine afterglow'' -- Fan, Piran \\& Xu 2006). ",
        "conclusions": "We have studied the broad band opt--UV to X--ray SED at three epochs which characterise different flux evolution phases observed (C06) both in the X--ray and opt--UV light curves up to $10^{5}$ s. The spectrum shows a soft ($kT\\sim$ 0.1--0.18 keV) X--ray \\bb component (coupled with a typical non--thermal component) together with a Rayleigh--Jeans tail in the opt--UV band. Both these components are almost steady in flux and slope, while the X--ray \\bb is undetected at late times (10$^4$--10$^5$ s). They have been interpreted (e.g. C06) as the shock breakout of the accompanying (nearly simultaneous) SN2006aj. However, if the X--ray and opt--UV emission belong to the same \\bb emission (which we also exclude by direct spectral fitting) its energetics is even larger than the total kinetic energy of the SN ejecta as derived from its late time spectroscopy (e.g. Mazzali et al. 2006; Pian et al. 2006), and cannot be cured by relativistic beaming, since the SN shock is not relativistic. Alternatively, only the X--ray or the opt--UV thermal component might be the shock breakout signal. The former possibility seems to be excluded both by detailed simulations (Li 2006) and because the SN expansion velocity, derived from the X--ray \\bb fits, is too small compared to that derived from optical spectroscopy at late times (Pian et al. 2006). On the other hand, a shock breakout in the opt--UV band would imply a too large SN expansion velocity to be consistent with the observed opt--UV early spectrum.  If the SN shock breakout went undetected both in the opt--UV and in the X--ray band we are left with two puzzling quests: (a) which is the nature of the observed early time opt--UV emission (i.e. before the SN radioactive decay emission sets in at 10$^{5}$ s) and (b) what is the origin of the X--ray thermal component?  We explored the possibility that the SED is produced by synchrotron self--Compton emission. A simple synchrotron fit to these data is satisfactory, with the opt--UV data corresponding to the self--absorbed synchrotron spectrum. Although the choice of the input parameter is not unique, the spectra at different times can be modelled by changing the slope of the electron distribution. A remarkable result of these fits is that also the late time (10$^{4}$--10$^{5}$ s) opt--UV emission should be described by the same model (i.e. with the same parameters of the prompt, but with steeper still slopes of the electron distribution). The X--ray flux is in this case first order self--Compton emission and naturally accounts for the unusually soft spectrum observed (which is hardly reconcilable with the normal flux time decay within the context of the standard afterglow model -- Fan, Piran \\& Xu 2006).  Finally, we considered the possibility that the \\bb observed in the X--ray band is a separate component. It is consistent to be photospheric emission from the cocoon or some stellar material energized by the GRB jet. The origin of the \\bb photons should correspond to dissipation occurring not much below the photospheric radius of this material."
    },
    "0608/astro-ph0608249_arXiv.txt": {
        "abstract": "We present {\\it Spitzer} imaging of the metal-deficient (Z\\,$\\simeq$30\\%\\,\\zsun) Local Group dwarf galaxy NGC\\,6822.  On spatial scales of $\\sim$130 pc, we study the nature of IR, \\halpha, \\HI, and radio continuum emission.  Nebular emission strength correlates with IR surface brightness; however, roughly half of the IR emission is associated with diffuse regions not luminous at \\halpha\\ (as found in previous studies).  The global ratio of dust to \\HI\\ gas in the ISM, while uncertain at the factor of $\\sim$2 level, is $\\sim$25 times lower than the global values derived for spiral galaxies using similar modeling techniques; localized ratios of dust to \\HI\\ gas are about a factor of five higher than the global value in NGC\\,6822. There are strong variations (factors of $\\sim$ 10) in the relative ratios of \\halpha\\ and IR flux throughout the central disk; the low dust content of NGC\\,6822 is likely responsible for the different \\halpha/IR ratios compared to those found in more metal-rich environments.  The \\halpha\\ and IR emission is associated with high-column density ($\\gsim$10$^{21}$ cm$^{-2}$) neutral gas. Increases in IR surface brightness appear to be affected by both increased radiation field strength and increased local gas density. Individual regions and the galaxy as a whole fall within the observed scatter of recent high-resolution studies of the radio-far IR correlation in nearby spiral galaxies; this is likely the result of depleted radio and far-IR emission strengths in the ISM of this dwarf galaxy. ",
        "introduction": "\\label{S1}  Nearby dwarf galaxies provide a unique opportunity to resolve the interaction of the multiple phases of the interstellar medium (ISM).  Star formation in these systems is more susceptible to small-scale processes (e.g., turbulence, feedback) than to the galaxy-wide perturbations (i.e., spiral density waves) that are common in more massive galaxies.  With their typically low nebular metallicities, star-forming dwarf galaxies serve as fiducial examples for comparison with the high-redshift star-forming systems predicted in the cold dark matter paradigm \\citep[e.g.,][]{babul96,ellis97,mateo98}.  \\begin{figure*}[!ht] \\plotone{f1.lowres.eps} \\caption{DSS image of NGC\\,6822, overlaid with a contour of \\HI\\ coulmn density at the 2\\,$\\times$\\,10$^{20}$ cm$^{-2}$ level.  The \\HI\\ disk is much larger than the high surface brightness optical body, which was the target of our IRAC and MIPS observations.  However, \\citet{deblok06} find an extended stellar population that exceeds the size of the \\HI\\ disk (see further discussion in \\S~\\ref{S1}).} \\label{figcap1} \\end{figure*}  The detailed processes of heating and cooling of the ISM are fundamental to our understanding of the nature of galaxy evolution. An environmentally-dependent fraction of UV photons produced by massive stars will be absorbed by dust and gas and re-radiated in the far-infrared (FIR) or in nebular emission lines.  This fraction will depend on various factors, including the local dust-to-gas ratio (related to the metallicity and hence the dust content) and the porosity of the local ISM.  In cases where the dust and stars are well mixed, and where the extinction is moderate, there will exist a correlation between the fluxes in various wavelength bands, including the ultraviolet (UV), \\halpha, FIR, and radio continuum.  Nearby dwarf galaxies offer a unique opportunity to study one of the regimes where these correlations may begin to break down.  NGC\\,6822 is a Local Group (D $=$ 490\\,$\\pm$\\,40 kpc; {Mateo 1998}\\nocite{mateo98}) dwarf irregular galaxy with a nebular abundance Z $\\simeq$ 30\\% \\zsun\\ \\citep{skillman89,lee06a}.  The total optical luminosity is $\\sim$9.4\\,$\\times$\\,10$^7$ \\lsun\\ \\citep{mateo98}, while the total \\HI\\ mass is 1.34\\,$\\times$\\,10$^{8}$ \\msun\\ (see detailed discussion in \\S~\\ref{S2.2} and {de~Blok \\& Walter 2006}\\nocite{deblok06}).  The system contains distributed low-level star formation throughout the disk as well as actively star-forming regions with associated \\halpha\\ emission which were first studied by \\citet{hubble25}.  The star formation has been quiescent over the past few Gyr, without major peaks or dips in the absolute rate \\citep{wyder01,wyder03}.  \\citet{deblok06} derive a total \\halpha\\ flux of 2\\,$\\times$\\,10$^{39}$ erg\\,s$^{-1}$, corresponding to a global \\halpha-based star formation rate (SFR) of $\\sim$ 0.015 \\msun\\,yr$^{-1}$ using the calibration of \\citet{kennicutt98}.  \\begin{figure*}[!ht] \\plotone{f2.lowres.eps} \\caption{\\HI\\ (a) and IRAC (b) images of NGC\\,6822. (a) shows the total \\HI\\ column density distribution (de~Blok \\& Walter 2000), while (b) shows the underlying stellar and warm dust components. The three-color image shows the 3.6 \\mm\\ band as blue, the 4.5 \\mm\\ band as green, and the 8 \\mm\\ band as red; regions of hot dust emission, indicative of active star formation, appear as diffuse regions of red emission. The box in (a) shows the approximate field of view shown in (b).} \\label{figcap2} \\end{figure*}  NGC\\,6822 boasts an unusual \\HI\\ distribution; in the extended \\HI\\ disk (see Figure~\\ref{figcap1}) lies one of the largest known holes in the ISM of a dwarf galaxy, first discussed by \\citet{deblok00}. \\citet{deblok06} argue that this hole is likely caused by the combined effects of stellar evolution.  That study also reveals the presence of a significant blue star population in the outer \\HI\\ disk, and a lower surface brightness, spheroidal red stellar population extending beyond the \\HI\\ distribution.  The extended \\HI\\ disk may include a dwarf companion galaxy in the northwest region (note that interactions with this system may have played a part in triggering some of the recent star formation in the galaxy; see {de~Blok \\& Walter 2000}\\nocite{deblok00}).  The well-sampled rotation curve shows that the system is highly dark-matter dominated \\citep{weldrake03}.  The proximity of NGC\\,6822 means that it is one of only a few places where the small-scale ($\\sim$100 pc) structure of the ISM can be studied at multiple wavelengths.  Previous infrared studies of NGC\\,6822 have examined the heating and cooling mechanisms in the ISM.  Using {\\it InfraRed Astronomical Satellite} ({\\iras}) data, \\citet{gallagher91} found a variable \\halpha/FIR ratio within the disk, with roughly 50\\% of the FIR emission arising from regions that are luminous at \\halpha.  The remaining fraction of diffuse IR flux was attributed to a low optical depth (i.e., low dust content) in the ISM.  Thus, regions of strong IR emission correspond to comparatively dusty regions that also contain a younger stellar population.  A second \\iras\\ study by \\citet{israel96} found that three of the bright IR peaks were associated with luminous \\HII\\ regions, while two equally bright sources were unassociated with obvious sources of ionizing radiation.  These authors confirm the prominent diffuse dust continuum emission found by \\citet{gallagher91}.  Further, they find, using single-temperature blackbody fits to the far-IR data, a dust-to-gas ratio (1.4\\,$\\times$\\,10$^{-4}$) that is much lower than typical values seen in more metal-rich systems such as the Milky Way (average dust-to-gas ratios of $\\sim$0.006--0.01; {Sodroski \\etal\\ [1997]}\\nocite{sodroski97}, {Li [2004]}\\nocite{li04}).  NGC\\,6822 was observed with the {\\it Infrared Satellite Observatory} at 6.75 and 15 \\mm, but only the most luminous \\HII\\ regions were detected \\citep{dale00,hunter01}.  In this paper, we present \\spitzer\\ Infrared Array Camera (IRAC) and Multiband Imaging Photometer for \\spitzer\\ (MIPS) imaging \\citep[see][]{werner04,fazio04,rieke04} of the high optical surface brightness component (i.e., the inner portion of the \\HI\\ disk) of NGC\\,6822 (see Figure~\\ref{figcap2} for a comparison of the field of view imaged and the size of the \\HI\\ disk).  These data were obtained as part of the {\\it Spitzer Infrared Nearby Galaxies Survey} (\\sings; see {Kennicutt \\etal\\ 2003}\\nocite{kennicutt03}). As the nearest galaxy in the \\sings\\ sample, these data probe the nature of mid- and far-IR emission at unprecedented sensitivity and resolution ($\\sim$130 pc; see further discussion below). We compare these \\spitzer\\ data to sensitive \\halpha, \\HI\\ and radio continuum imaging in order to study the relation between \\HI\\ surface density and the locations and intensities of star formation and IR emission.  \\begin{figure*}[!ht] \\begin{center} \\plotone{f3.lowres.eps} \\end{center} \\caption{2MASS J (a), 3.6 and 8 \\mm\\ (b, c), 24, 70, and 160 \\mm\\ (d, e, f), \\halpha\\ (g), 20\\,cm radio continuum (h), and \\HI\\ spectral line (i) images of NGC\\,6822.  All image scales are linear, except for the \\halpha\\ image (g), which is shown in a logarithmic stretch to highlight the low-surface brightness ionized gas that permeates much of the central disk of the system.  Beam sizes are shown as boxed ellipses for the MIPS and radio data, in the bottom right corners.  Overlaid are the sizes and locations of the smaller apertures [diameter $=$ 114\\arcsec\\ $\\simeq$ 3\\,$\\times$\\,FWHM (160 \\mm) $\\simeq$ 270 pc] discussed in the text; these regions are labeled in panel (a).  Note that the regions that encompass the major IR emission complexes (regions 3, 7, 8, 11, and 14; see also Table~\\ref{t1}) have also been analyzed using apertures with four times the area. The 8 \\mm\\ image does not cover the full field of view shown in panel (c); since the 4.5 \\mm\\ band has the same orientation, flux densities are not available for region 13 in either of these bandpasses.  Note that the radio continuum image in panel (h) shows only the apertures where 20\\,cm emission is detected; numerous background sources exist within this field of view (see further discussion in \\S~\\ref{S3.3}).  The aperture sizes are large enough to cover the bulk of the extended long-wavelength PSFs, minimizing aperture correction effects (see further discussion in \\S~\\ref{S3}).} \\label{figcap3} \\end{figure*}  \\begin{deluxetable*}{lcccc}[ht] \\tablecaption{Infrared Sources in NGC\\,6822} \\tablewidth{0pt} \\tablehead{ \\colhead{Identification} &\\colhead{$\\alpha$ (J2000)} &\\colhead{$\\delta$ (J2000)} &\\colhead{Previous I.D.} &\\colhead{References}} \\startdata NGC\\,6822--1  &19:44:31.28 &$-$14:47:46.69 &I96\\,1 &3\\\\ NGC\\,6822--2  &19:44:31.92 &$-$14:44:16.70 &Hubble\\,II &1\\\\ NGC\\,6822--3  &19:44:32.96 &$-$14:42:10.71 &Hubble\\,I/III, G91\\,3, I96\\,2 &1, 2\\tablenotemark{a}, 3\\\\ NGC\\,6822--4  &19:44:37.89 &$-$14:51:37.79 &I96\\,3 &3\\\\ NGC\\,6822--5  &19:44:43.69 &$-$14:50:19.86 &I96\\,3 &3\\\\ NGC\\,6822--6  &19:44:46.18 &$-$14:46:19.89 & &\\\\ NGC\\,6822--7  &19:44:49.47 &$-$14:52:40.93 &Hubble\\,IV, G91\\,1, I96\\,4 &1, 2, 3\\\\ NGC\\,6822--8  &19:44:52.81 &$-$14:43:04.97 &Hubble\\,V, G91\\,3, I96\\,5 &1, 2, 3\\\\ NGC\\,6822--9  &19:44:55.69 &$-$14:50:17.00 & &\\\\ NGC\\,6822--10 &19:44:57.36 &$-$14:44:50.01 & &\\\\ NGC\\,6822--11 &19:44:58.59 &$-$14:47:47.02 &Hubble\\,VI/VII, G91\\,2, I96\\,6 &1, 2, 3\\\\ NGC\\,6822--12 &19:44:59.82 &$-$14:54:41.03 & &\\\\ NGC\\,6822--13 &19:45:04.99 &$-$14:57:29.07 &I96\\,7 &3\\\\ NGC\\,6822--14 &19:45:05.01 &$-$14:43:17.07 &Hubble\\,X, G91\\,3, I96\\,8 &1, 2, 3\\\\ NGC\\,6822--15 &19:45:09.35 &$-$14:48:44.10 &Hubble\\,IX, G91\\,2, I96\\,9 &1, 2, 3\\\\ NGC\\,6822--16 &19:45:16.68 &$-$14:45:24.26 & &\\\\ \\enddata \\label{t1} \\tablenotetext{a}{The sources in \\citet{gallagher91} are not defined by low-resolution (1.5\\arcmin\\,$\\times$\\,5\\arcmin) \\iras\\ imaging; the cross-identifications here correspond to only the highest-surface brightness IR peaks in the IR or \\halpha\\ images.} \\tablerefs{1 - \\citet{hubble25}; 2 - \\citet{gallagher91}; 3 - \\citet{israel96}} \\end{deluxetable*} ",
        "conclusions": "\\label{S4}  We have presented IR observations of the Local Group dwarf irregular galaxy NGC\\,6822 obtained with {\\it Spitzer} as part of \\sings.  The galaxy is highly resolved at all imaging bands; the resolution limit is driven by the FWHM of the 160 \\mm\\ PSF.  At $\\sim$130 pc physical resolution, we study the nature of FIR emission and compare to observations in the optical and radio. These sensitive data reveal a wealth of structure within the galaxy.  Our study confirms some previous results, while also offering new insights into the nature of infrared emission in the metal-poor ISM.  We discuss these topics in more detail below.  The total monochromatic FIR flux densities are found to be S$_{\\rm 24} =$ 2.51\\,$\\pm$\\,0.50 Jy; S$_{\\rm 70} =$ 53.2\\,$\\pm$\\,15 Jy; S$_{\\rm 160} =$ 136.2\\,$\\pm$\\,40 Jy.  SED model fits to these data allow comparison with previous {\\it IRAS} flux measurements; our {\\it Spitzer} flux densities are in excellent agreement with the measurements in \\citet{israel96}.  Using the relations presented in \\citet{dale02}, these global flux densities correspond to a total IR flux of $\\sim$ 5.7\\,$\\times$\\,10$^{-12}$ W\\,m$^{-2}$ ($\\sim$ 4.3\\,$\\times$\\,10$^7$ L$_{\\odot}$).  IR emission at 24, 70 and 160 \\mm\\ is found only within the central, high-\\HI\\ column density ($\\gsim$ 10$^{21}$ cm$^{-2}$) region of NGC\\,6822. The total \\HI\\ distribution is much more extended than the high-surface brightness stellar or dust emission components.  The 24 \\mm\\ surface brightness profile closely traces that of \\halpha, delineating regions of ongoing star formation.  The profiles of 70 and 160 \\mm\\ emission also trace local \\HI\\ column density maxima.  Comparing (monochromatic and total) FIR emission with a sensitive \\halpha\\ image reveals a strong morphological correlation between IR and \\halpha\\ surface brightness; here we are directly observing the re-processing of photons by dust.  Roughly 50\\% of the TIR flux is attributable to discrete FIR sources, all of which have at least a low-level \\halpha\\ counterpart; the remaining FIR emission arises from a ``diffuse'' component of the ISM (slightly stronger at 70 and 160 \\mm\\ than at 24 \\mm; see detailed discussion in \\S~\\ref{S3.1.1}) that is apparently not associated with massive star formation. \\citet{gallagher91} and \\citet{israel96} find similar results from {\\it IRAS} imaging of NGC\\,6822; we interpret this as the re-processing of non-ionizing photons in the ISM or as the escape of radiation from the star formation regions.  We find variations in the relative ratios of nebular and dust continuum emission. While differential extinction can cause some of the observed fluctuations, it cannot account for the full breadth of properties seen in the galaxy.  Using \\sings\\ observations of the starbursting dwarf galaxies IC\\,2574 and NGC\\,1705, \\citet{cannon05,cannon06} found similar variations and attributed them to the direct impact of recent star formation on the local ISM (i.e, production and re-processing of photons).  The results presented here extend this trend to more typical, quiescent dwarfs.  Evidence is mounting from spatially resolved {\\it Spitzer} observations of dwarf galaxies that these systems follow a different trend of 24 \\mm\\ vs. \\halpha\\ luminosity than what is seen in more metal-rich systems (see, e.g., {Calzetti \\etal\\ 2005}\\nocite{calzetti05}).  This result is in agreement with previous observations of dwarf galaxies \\citep[e.g.,][]{hunter89}.  Two different scenarios can account for this observational trend. The first one, physically simple and more commonly used, is related to the low dust content of low-metallicity galaxies. Here, low-metallicity galaxies simply do not have enough metals to form dust in the same fraction as more metal rich environments.  The second scenario posits that dwarfs harbor massive components of cold dust (T $\\sim$ 5--10 K) that is simply not heated by ongoing star formation. Such dust would peak past 200 \\mm\\ and remain undetected in the present \\spitzer\\ observations, contributing to part of the observed discrepancy; observational evidence exists both for and against such a component in nearby galaxies (see detailed discussion in \\S~\\ref{S3.2.2}). Clearly, sensitive observations of a large sample of dwarf galaxies between 200 and 1000 \\mm\\ are needed to address this issue.  As mentioned previously, our global flux densities for NGC\\,6822 are in good agreement with those from \\citet{israel96}.  We use sophisticated SED modeling techniques [\\citet{dale01} and \\citet{dale02}; {Li \\& Draine (2001, 2002)}\\nocite{li01,li02} and {Draine \\& Li (2006)}\\nocite{draine06}] to derive dust masses that are higher than those derived by \\citet{israel96}.  The differences in mass are thus attributable solely to our more physical modeling approach, which takes into account the variety of dust species and heating levels in each region within the disk (see detailed discussion in \\S~\\ref{S3.2.2}).  For a given TIR luminosity, our models recover a higher dust mass than simple blackbody fits (see detailed discussion in \\S~\\ref{S3.2.2}).  The treatment of multiple grain populations likely means that the increased dust masses are a better representation of the mass radiating between 10 and 160 \\mm. Using these techniques, we derive a total dust mass of $\\sim$10$^5$ \\msun\\ (this value is derived using the Li \\& Draine models).  Comparing to our high-resolution \\HI\\ imaging, we can explore the mass ratio of warm dust to \\HI\\ gas (\\dgr) in the ISM.  Given the low nebular abundance of the system (Z $\\simeq$ 30\\% \\zsun), we can thus provide constraints on a potential trend of \\dgr\\ with abundance. Metal-poor dwarf galaxies might be expected to have a lower \\dgr\\ (i.e., higher gas mass fraction) than more metal-rich galaxies such as the Milky Way \\citep[e.g., ][]{thronson86}.  \\citet{lisenfeld98} have argued that dwarf irregular galaxies have a clear relation between these two quantities in a global sense (using dust masses derived from single-temperature blackbody fits).  Here, we derive a global \\dgr\\ $\\simeq$ 8\\,$\\times$\\,10$^{-4}$.  This is $\\sim$25 times lower than the average global values found in the new \\sings\\ study of Draine \\etal\\ (2006, in prep.), wherein dust masses are calculated for a large sample of galaxies (excluding dwarfs).  On local scales ($\\sim$130 pc), the dust mass per unit \\HI\\ mass is higher (by a factor of $\\sim$5), though still lower than the global values found by Draine et\\,al.  We propose a simple model for the relative ratios of dust mass to \\HI\\ mass in the ISM.  As noted above, dwarf galaxies may contain large reservoirs of cold dust (T $\\sim$ 5--10 K) that are undetectable with \\spitzer\\ observations \\citep[e.g.][]{madden02,galliano05}.  Note, however, that the dilution of the radiation field required for this dust to remain at such low temperatures may require it to be located at large distances from the star formation regions and thus essentially decoupled from the active ISM (note that other mechanisms may also be important in dwarfs, including dust destruction by supernova shocks; see, e.g., {Bot \\etal\\ 2004}\\nocite{bot04}). Superposed on this putative cool component are regions radiating between 10 and 160 \\mm; here, photons are being absorbed and re-radiated as IR photons.  The low dust contents of metal-poor galaxies plays an important role in the microphysics of this process. Galaxies with widespread, elevated star formation will heat more of their total dust component than more quiescent systems.  When more sensitive observations are possible between 200 and 1000 \\mm, we will be able to quantify the relative masses of these components in low-mass dwarf galaxies.  Finally, it is interesting that NGC\\,6822 seems to fall near the canonical measures of the radio-FIR relation \\citep[e.g.,][]{bell03a,murphy06a}.  20\\,cm radio continuum imaging reveals six high surface brightness sources, each of which appears to be of thermal origin. The monochromatic and total IR ``q'' parameters are similar to those found for more massive galaxies.  This agreement is interesting, given that nonthermal emission dominates the radio continua of spiral disks.  The simplest explanation is that depleted radio and far-IR emission strengths conspire to produce the observed values.  These findings in NGC\\,6822 are in marked contrast to recent results for very active star-forming dwarf galaxies, where significant deviations from the canonical radio-FIR relation are found \\citep{cannon05,cannon06}.  Taken together, these data suggest that star formation intensity is one of the important parameters that governs the observed radio-FIR relation in dwarf galaxies."
    },
    "0608/astro-ph0608280_arXiv.txt": {
        "abstract": "We present the current estimates of the Galactic merger rate of double-neutron-star (DNS) systems. Using a statistical analysis method, we calculate the probability distribution function (PDF) of the rate estimates, which allows us to assign confidence intervals to the rate estimates. We calculate the Galactic DNS merger rate based on the three known systems B1913+16, B1534+12, and J0737-3039. The discovery of J0737-3039 increases the estimated DNS merger rate by a factor $\\sim$6 than what is previously known. The most likely values of DNS merger rate lie in the range $3-190$ Myr$^{-1}$ depending on different pulsar models. Motivated by a strong correlation between the peak rate estimates and a pulsar luminosity function, we calculate a {\\em global} probability distribution as a single representation of the parameter space covered by different pulsar population models. We compare the global PDF with the observed supernova Ib/c rate, which sets an upper limit on the DNS merger rate. Finally, we remark on implications of new discoveries such as of J1756$-$2251, the 4th DNS in the Galactic disk, and J1906+0746, a possible DNS system. ",
        "introduction": "Merging (or inspiralling) DNS systems are one of the prime targets for ground-based gravitational-wave (GW) interferometers such as GEO600, TAMA, VIRGO, and LIGO (the Laser Interferometer Gravitational-Wave Observatory) \\citep{ligo}. Event rates of the DNS inspiral searches by these detectors can be inferred using the rate estimates with an extrapolation out to the maximum detection distances for any detector under consideration. Before 2003, the Galactic DNS merger rate had been estimated between $\\sim10^{-7} - 10^{-5}$ yr$^{-1}$ (see Kalogera et al.\\ 2001\\nocite{knst} and references therein). At that time, there were only two systems available for empirical studies, PSRs~B1913+16 \\citep{ht75} and B1534+12 \\citep{w91}. We calculated the PDF of Galactic DNS merger rates, $P(R)$, based on these two systems (see Kim et al.\\ 2003, hereafter KKL, for further discussions). Soon after the discovery of the highly relativistic system PSR~J0737$-$3039 \\nocite{0737A}, we were able to revise $P(R)$ including PSR~J0737$-$3039 in collaboration with the observation team\\footnote{Only the millisecond component (J0737-3039A) is considered in our calculation. For instance, the current age of the system is derived from the pulsar A.}\\citep{0737A, k04}. The discovery of J0737$-$3039, resulted in a significant  increase in the estimated Galactic DNS by a factor of $\\sim$6. This implies a boost in event rates for DNS searches for GW interferometers. In this paper, we present our recent results on: (i) the PDF of Galactic DNS merger rate estimates with updated observations; (ii) the {\\em global\\/} PDF of rate estimates considering the systematic uncertainties; (iii) constraints on upper limits for rate estimates based on the observed supernova rate incorporated with our theoretical understanding of the SN-DNS relation; (iv) the approximate contribution of J1756$-$2251 to the Galactic DNS merger rates and uncertainties in our assumptions on the efficiency of the acceleration searches for Parkes multibeam pulsar survey (PMPS). Finally, we discuss implications of the most recently discovered pulsar binary J1906+0746 \\citep{j1906}. ",
        "conclusions": "We summarize our results on the Galactic merger rate of double neutron stars (DNS) in view of the recent discovery of PSR J0737-3039. We update the results based on the most recent observations. The {\\em global} PDF of rate estimates incorporates the already known systematics due to the radio pulsar luminosity function. The most likely value obtained from the global distribution is at only $\\sim$ 13 Myr$^{-1}$, and a reanalysis of the {\\em current} pulsar sample and radio luminosities is needed for a reliable assessment of the best fitting distribution. The DNS merger rate can be constrained by empirical Type Ib/c supernova rates, and we compare the supernova rate with the global PDF. Recent discoveries of relativistic systems, such as J1756$-$2251 and J1906+0746 are important for the reduction of the uncertainties associated with the empirical rate estimates. As for J1756$-$2251, we expect their contribution to the Galactic DNS merger rate is negligible, and this pulsar can be seen another sample of B1913+16-like population. It is not clear whether the companion of J1906+0746 is a neutron star or a white dwarf. This system may increase the estimated merger rate of the relevant population by a factor of 2 whether it is a DNS or another eccentric NS-WD binary."
    },
    "0608/hep-ph0608060_arXiv.txt": {
        "abstract": "\\medskip In the light of recent neutrino oscillation and non-oscillation data, we revisit the phenomenological constraints applicable to three observables sensitive to absolute neutrino masses: The effective neutrino mass in single beta decay $(m_\\beta)$; the effective Majorana neutrino mass in neutrinoless double beta decay $(m_{\\beta\\beta})$; and the sum of neutrino masses in cosmology $(\\Sigma)$. In particular, we include the constraints coming from the first Main Injector Neutrino Oscillation Search (MINOS) data and from the Wilkinson Microwave Anisotropy Probe (WMAP) three-year (3y) data, as well as other relevant cosmological data and priors. We find that the largest neutrino squared mass difference is determined with a  15\\% accuracy (at $2\\sigma$) after adding MINOS to world data. We also find upper bounds on the sum of neutrino masses $\\Sigma$ ranging from $\\sim2$~eV (WMAP-3y data only) to $\\sim0.2$~eV (all cosmological data) at $2\\sigma$, in agreement with previous studies. In addition, we discuss the connection of such bounds with those placed on the matter power spectrum normalization parameter $\\sigma_8$. We show how the partial degeneracy between $\\Sigma$ and $\\sigma_8$ in WMAP-3y data is broken by adding further cosmological data, and how the overall preference of such data for relatively high values of $\\sigma_8$ pushes the upper bound of $\\Sigma$ in the sub-eV range. Finally, for various combination of data sets, we revisit the (in)compatibility between current $\\Sigma$ and $m_{\\beta\\beta}$ constraints (and claims), and derive quantitative predictions for future single and double beta decay experiments. ",
        "introduction": "One of the greatest challenges in current neutrino physics is to establish the absolute masses $(m_1,m_2,m_3)$ of the three neutrino mass eigenstates $(\\nu_1,\\nu_2,\\nu_3)$, for which there is ample experimental evidence of mixing with the three neutrino flavor eigenstates $(\\nu_e,\\nu_\\mu,\\nu_\\tau)$ through a unitary matrix $U(\\theta_{12},\\theta_{23},\\theta_{13})$, where $\\theta_{ij}$ are the mixing angles \\cite{PDG4}.  Atmospheric, solar, reactor, and accelerator neutrino oscillation experiments constrain two squared mass differences, $\\delta m^2$ and $\\Delta m^2$ (with $\\delta m^2\\ll \\Delta m^2$), parametrized as \\begin{equation} \\label{masses} (m^2_1,\\,m^2_2,\\,m^2_3)=\\frac{m^2_1+m^2_2}{2}+ \\left( -\\frac{\\delta m^2}{2},\\,+\\frac{\\delta m^2}{2},\\,\\pm\\Delta m^2 \\right)\\ , \\end{equation} where the cases $+\\Delta m^2$  and $-\\Delta m^2$ distinguish the so-called normal hierarchy (NH) and inverted hierarchy (IH), respectively. The same data also measure $\\theta_{12}$ and $\\theta_{23}$, and place upper bounds on $\\theta_{13}$ (see \\cite{PPNP,Viss} for recent reviews).  Non-oscillation neutrino data from single $\\beta$ decay, from neutrinoless double $\\beta$ decay ($0\\nu2\\beta$), and from cosmology, add independent constraints on the absolute neutrino masses, through their sensitivity to the observables \\cite{PDG4} \\begin{eqnarray} \\label{mb} m_\\beta &=& \\left[c^2_{13}c^2_{12}m^2_1+c^2_{13}s^2_{12}m^2_2+s^2_{13}m^2_3 \\right]^\\frac{1}{2}\\ ,\\\\ \\label{mbb} m_{\\beta\\beta} &=& \\left| c^2_{13}c^2_{12}m_1+c^2_{13}s^2_{12}m_2e^{i\\phi_2}+s^2_{13}m_3 e^{i\\phi_3}\\right|\\ ,\\\\ \\label{Sigma} \\Sigma &=& m_1+m_2+m_3\\ , \\end{eqnarray} respectively (with $c_{ij}=\\cos\\theta_{ij}$ and $s_{ij}=\\sin\\theta_{ij}$, while $\\phi_{2,3}$ are unknown Majorana phases).  In a previous work \\cite{fogli}, a global analysis of world oscillation and non-oscillation data was performed, with the purpose of showing the interplay and the (in)compatibility of different data sets in constraining the $(m_\\beta,m_{\\beta\\beta},\\Sigma)$ parameter space. In this work, we revisit such constraints in the light of two recent relevant developments: the first results from the Main Injector Neutrino Oscillation Search (MINOS) \\cite{MINOS} (see Sec.~II), and the three-year (3y) results  from the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{WMAP3} (see sec.~IV), which have a direct impact on $\\Delta m^2$ and $\\Sigma$, respectively. In Sec.~II we show that, after MINOS, the $\\Delta m^2$ parameter is globally determined with an accuracy of 15\\% at $2\\sigma$. In Sec.~V we discuss the degeneracy between $\\Sigma$ and the cosmological parameter $\\sigma_8$ (which normalizes the matter power spectrum \\cite{PDG4,WMAP3}), as well as the breaking of such degeneracy with the addition of further cosmological data, which strenghten the $2\\sigma$ upper bound on $\\Sigma$ from $\\sim 2$ eV (WMAP 3y only) to $\\sim 0.2$ eV (all cosmological data). Such limits on $\\Sigma$ are in agreement with other recent analyses \\cite{WMAP3,Seljak:2006bg,goobar,fukugita,raffelt,cirelli,pastor}. Older $\\beta$ and $0\\nu2\\beta$ results, including the $0\\nu2\\beta$ signal claim of \\cite{Kl04}, are  also briefly reviewed (see Sec.~III). In Sec~V we discuss the issues of the (in)compatibility among the previous data sets  and of their possible combinations and constraints in the parameter space $(m_\\beta,m_{\\beta\\beta},\\Sigma)$. In particular, we show that the case with ``WMAP 3y only'' still allows a global combination (at the $2\\sigma$ level) with the $0\\nu2\\beta$ signal claimed in \\cite{Kl04}, while the addition of other cosmological data excludes this combination (at the same or higher significance level). Quantitative implications for future $\\beta$ and $0\\nu2\\beta$ decay searches are worked out for various combinations of current data  (Secs.~V and VI). Conclusions are given in Sec.~VII. ",
        "conclusions": "Building on previous work \\cite{fogli} and on updated experimental input, we have revisited the phenomenological constraints applicable to three observables sensitive to absolute neutrino masses: The effective neutrino mass in single beta decay $(m_\\beta)$; the effective Majorana neutrino mass in neutrinoless double beta decay $(m_{\\beta\\beta})$; and the sum of neutrino masses in cosmology $(\\Sigma)$. In particular, we have included the constraints coming from the first MINOS results \\cite{MINOS} and from the WMAP-3y data \\cite{WMAP3}, as well as other relevant cosmological data and priors. We have found that the largest neutrino mass squared difference $\\Delta m^2$ is now determined with a  15\\% accuracy (at $2\\sigma$). We have also examined the upper bounds on the sum of neutrino masses $\\Sigma$, as well as their correlations with the matter power spectrum normalization parameter $\\sigma_8$. The bounds range  from $\\Sigma \\lesssim2$~eV (WMAP-3y data only) to $\\Sigma\\lesssim 0.2$~eV (all cosmological data) at $2\\sigma$, in agreement with previous studies. Finally, for various possible combination of data sets, we have revisited the (in)compatibility between current $\\Sigma$ and $m_{\\beta\\beta}$ constraints, and have derived quantitative predictions for single and double beta-decay observables, which can be useful to evaluate the sensitivity required by present and future experiments in order to explore the currently allowed parameter space."
    },
    "0608/astro-ph0608143_arXiv.txt": {
        "abstract": "We present IRTF SpeX observations of the M/L binary system 2MASS J17072343-0558249.  SpeX imaging resolves the system into a 1$\\farcs$01$\\pm$0.17 visual binary in which both components have red near infrared colors.  Resolved low-resolution (R$\\sim$150) 0.8--2.5 $\\micron$ spectroscopy reveals strong H$_2$O, CO and FeH bands and alkali lines in the spectra of both components, characteristic of late-type M and L dwarfs.  A comparison to a sample of late-type field dwarf spectra indicates spectral types M9 and L3.  Despite the small proper motion of the system (0$\\farcs$100$\\pm$0$\\farcs$009 yr$^{-1}$), imaging observations over 2.5~yr provide strong evidence that the two components share common proper motion.  Physical association is also likely due to the small spatial volume occupied by the two components (based on spectrophotometric distances estimates of 15$\\pm$1 pc) as compared to the relatively low spatial density of low mass field stars.  The projected separation of the system is 15$\\pm$3 AU, similar to other late-type M and L binaries.  Assuming a system age of 0.5-5 Gyr, we estimate the masses of the binary components to be 0.072-0.083 and 0.064-0.077 M$_{\\sun}$, with an orbital period of roughly 150-300 yr. While this is nominally too long a baseline for astrometric mass measurements, the proximity and relatively wide angular separation of the 2MASS J1707-0558AB pair makes it an ideal system for studying the M dwarf/L dwarf transition at a fixed age and metallicity. ",
        "introduction": "Wide field photographic imaging and proper motion surveys of the past were largely incapable of producing an accurate census of nearby very low mass (VLM; M $<$ 0.1~M$_{\\sun}$) stars and brown dwarfs. Late-type dwarfs emit weakly at visual bands, but are brighter at red optical and near infrared wavelengths.  Progress in optical and infrared detector technology have paved the way for the current generation of wide field sky surveys, most notably the Deep Near Infrared Sky Survey \\citep{epc97}, the Sloan Digital Sky Survey \\citep{yor00} and the Two Micron All Sky Survey \\citep[hereafter 2MASS]{skr06}.  Searches for faint red sources in these catalogs (e.g., \\citealt{del97,kir99,fan00}) have identified hundreds of very low mass stars and brown dwarfs, greatly expanding our VLM census and revealing the new spectral classes L \\citep{kir99,mar99} and T \\citep{bur02,geb02}.  The nascent field of VLM stars has been confronted with a number of fundamental questions regarding the properties of low luminosity sources, including formation scenarios, thermal evolution, chemical compositions and dynamics of cool stellar/substellar atmospheres, and the initial mass function.  Multiple star systems are important laboratories for understanding these physical properties.  Binary systems are key in the determination of stellar masses, which can be used to calibrate theoretical evolutionary and structure models. Binary parameters probe the star formation process, and relative comparisons can be made between coeval components.  Searches for cool dwarf binaries through high resolution imaging have been conducted to explore the nature of VLM stars and brown dwarfs (e.g., \\citealt{mar99,koe99,rei01,clo02,clo03,bou03,bur03a, giz03,sie05,law06,bur06c}).  The result of these efforts has been the discovery of roughly 75 VLM binary systems with varying mass ratios and projected separations (cf.\\ Burgasser et al.\\ 2006).  The properties of the current sample of VLM stars and brown dwarfs suggests that both the multiplicity fraction and peak of the semimajor axis distribution are directly related to the primary star mass.  G to M binaries tend to peak at separations of 3-30 AU \\citep{fis92,hen93,rei97}, with binary fractions $f_{bin}$ ranging from $\\sim$65\\% \\citep{duq91} down to $\\sim$30\\% \\citep{fis92,rei04,del04}.  In contrast, nearly all VLM binaries have separations $\\rho~$$\\lesssim$ 20 AU, with a binary fraction $f_{bin}$ $\\approx$ 15\\% (e.g., \\citealt{clo03}; however, see also \\citealt{max05}).  The low frequency and preference of small separations for VLM binaries is a challenge for star formation theories, and a clear understanding of binary parameters as a function of mass, age, and metallicity provides empirical clues for understanding VLM star and brown dwarf formation processes.  Studies of VLM binaries also reveal the detailed properties of cool star and brown dwarf atmospheres.  Dust formation and evolution remains an outstanding problem across the spectral transition from M dwarfs to L dwarfs (hereafter the M/L transition), where lower temperatures and higher pressures enable some molecules to form solid condensates and descend from the photosphere, changing the overall morphology of the spectrum and the atmospheric pressure/temperature profile \\citep{lun89,tsu96,lod99,lod02,burr99,ack01,lod02b}.  In particular, M dwarfs are characterized by their strong TiO and VO absorption bands; but as temperatures descend into the L dwarf regime, refractory elements such as Fe, Mg, Ti, V, Al, and Ca are removed from the gas in the photosphere by the condensation and sedimentation process.  Theoretical models examining these processes are compared to the existing spectra of field dwarfs, which have a variety of metallicities, masses, and ages.  M/L binaries, on the other hand, can help clarify the existence of atmospheric condensation by isolating parameters such as age and metallicity, assuming coevality.  A total of 14 binaries comprised of M and L dwarf components have been discovered to date \\citep[see Burgasser et al.\\ 2006]{giz00,giz03,bou03,clo03,fre03,mar03,sie03,sie05,cha04,cha05,bil05,bur06}. This paper presents the discovery of a new M/L binary\\footnote{Reid et al.\\ (2006) have concurrently resolved this system using the Hubble Space Telescope NICMOS instrument.  Their photometric estimates of the spectral types and distance of the 2MASS 1707-0558AB system agree with our analysis.}, 2MASS J17072343-0558249AB (hereafter 2MASS 1707-0558), identified via resolved near infrared imaging and spectroscopy using the SpeX instrument \\citep{ray03} mounted at the 3 m NASA Infrared Telescope Facility (IRTF). In $\\S$~2 we describe our observations and data reduction, and present results.  In $\\S$~3 we analyze these data, determining individual spectral types for the resolved components and the photometric properties of the system.  We argue for physical association by common proper motion, similar spectrophotometric distances, and comparing the volume occupied by the M/L pair to the measured density of VLM dwarfs in the Solar Neighborhood.  We also make preliminary estimates of the individual masses and orbital characteristics of the system.  We discuss our results in $\\S$~4, placing the 2MASS 1707-0558 system in context with other VLM binaries. This work is summarized in $\\S$~5. ",
        "conclusions": "\\subsection{2MASS 1707-0558AB System Characteristics}  We calculate the projected physical separation of 2MASS 1707-0558AB to be 15$\\pm$3 AU, based on the astrometric measurements and derived spectrophotometric distance.  Temperatures for the two components can be derived using the T$_{eff}$/spectral type relation for low mass stars from \\citet{gol04}, yielding T$_{eff}$ = 2400$\\pm$175 K for 2MASS 1707-0558A and 1950$\\pm$190 K for 2MASS 1707-0558B.  We calculate individual masses using the \\citet{burr97} models for ages of 0.5 and 5 Gyr, sampling the typical ages for field dwarfs (see below).  The derived masses are 0.072-0.083 M$_{\\sun}$ and 0.064-0.077 M$_{\\sun}$ for the A and B components, respectively, with a total system mass of 0.136-0.160 M$_{\\sun}$ and a mass ratio $q \\equiv$ M$_2$/M$_1$ = 0.88-0.92.  This mass ratio is again consistent with the properties of VLM binaries, which are predominantly near-equal mass systems.  Combining the estimates of the physical separation and mass indicates an orbital period of roughly 150-300 yr, an unfortunately unreasonable timescale for dynamical mass measurements.  A summary of the 2MASS 1707-0558AB system characteristics can be found in Table 4.  \\subsection{Broader Implications for M and L Dwarfs}  This binary system is similar to other field binary systems, with a small separation and nearly equal mass components.  93\\% of known VLM binaries have separations less than 20~AU, like 2MASS 1707-0558, and 77\\% have $q \\geq 0.8$ \\citep{bur06b}.  Therefore, this system further strengthens the case that VLM binaries have specific traits that may be related to their formation mechanism.  The M9/L3 binary 2MASS 1707-0558 is also a useful probe of the M/L transition, since the components have presumably coevolved with a fixed age and metallicity.  Resolved binaries with spectroscopy are ideal systems to study the complex atmospheric chemistry of low mass stars, especially the thickening of dust, effects of condensation on the spectral energy distribution, and the relative heights of the cloud base.  M dwarfs are characterized by their TiO and VO bands \\citep{kir99}, while the transition into the L dwarf regime is marked by the depletion of these molecules through condensation, leading to the expression of hydrides such as FeH and CrH \\citep{lod02}.  The atmospheric chemistry has been modeled for an assortment of field stars, yielding various masses, ages, and intrinsic metallicities.  A study of resolved binary systems straddling the M/L boundary, such as 2MASS 1707-0558, will permit a fair comparison of low mass atmospheres.  In addition, comparisons of the individual rotational velocities through atmospheric line broadening may constrain particular formation and evolution models.  A direct measurement would determine if the components have equivalent or varied rotation rates, clarifying a dependency on either the mass or the origin of the object. Chromospheric activity is not well understood in low mass stars, but it is likely to be correlated to the spin rate of the star, possibly as a result of the magnetic fields induced by a turbulent dynamo mechanism \\citep[see also \\citealt{cha06}]{dur93}.  Field studies suggest that H$\\alpha$ activity increases from early-type to late-type M dwarfs, and then diminishes through the early-type L dwarfs \\citep{giz00,moh03,wes04}.  An unresolved optical spectrum of 2MASS 1707-0558 shows modest H$\\alpha$ emission \\citep{giz02}, but the origins of this detection are unclear.  Is this activity present in both components, and at what level does each component contribute to this detection?  Furthermore, a measurement of the soft X-ray contributions of each component (resolvable with {\\em Chandra}) will help further constrain the magnetic field mechanisms and their effects on coronal heating.  The substellar nature of brown dwarfs can be directly tested through measurements of the 6708 {\\AA} \\ion{Li}{1} line, but unfortunately the original optical spectrum was unable to resolve this line.  The existence of primordial lithium is a conclusive determinant of substellar nature, as it is quickly destroyed when core temperatures reaches $\\sim$ 2$\\times$10$^{6}$ K, a prerequisite reached by VLM stars with masses greater than $\\sim$0.06M$_{\\sun}$ \\citep{reb92,mag93}.  It should be noted that an object can lie below the hydrogen-burning minimum mass (i.e., substellar) and fail to exhibit \\ion{Li}{1} absorption.  Depending on the age of the system, either one or both of the components will show \\ion{Li}{1} in absorption since they follow different evolutionary tracks.  Hence, in the situation of a binary such as 2MASS 1707-0558 that may straddle the substellar boundary, the existence of the \\ion{Li}{1} line can be also used to trace the age of the system (cf., \\citealt{liu05}). Three possible scenarios exist: 2MASS 1707-0558A and B could both be brown dwarfs and contain \\ion{Li}{1} undepleted from their atmospheres ($t$ $\\le$ 0.5 Gyr), 2MASS 1707-0558A could be a star or brown dwarf with mass greater than $\\sim$0.06M$_{\\sun}$ and 2MASS 1707-0558B a Li-bearing brown dwarf with a mass less than $\\sim$0.06M$_{\\sun}$ (0.5 Gyr $<$ $t$ $<$ 1 Gyr), or both components could be low mass stars or brown dwarfs with masses greater than $\\sim$0.06M$_{\\sun}$ ($t$ $>$ 1 Gyr).  2MASS 1707-0558 is therefore a rare VLM field system that can be assigned an accurate age through the \\ion{Li}{1} diagnostic.  An alternative age diagnostic, although less reliable, is a consideration of space kinematics and chromospheric activity.  Stars are generally born with the space motion of their natal cloud, and over time accumulate higher space velocities through interactions with the components of the Galactic disk \\citep{wei77}.  \\citet{giz00} have found that a sample of field M and L dwarfs separate into two groups according to kinematics and chromospheric activity.  An old population of M dwarfs have high tangential velocities (v$_{tan}$ $>$ 20 km s$^{-1}$) and strong H$\\alpha$ emission, while a younger population exhibits low kinematics and low activity.  Assuming our spectrophotometric distance and the proper motion as measured by the SuperCosmos Sky Survey, 2MASS 1707-0558 has a V$_{tan}$ of 7 km s$^{-1}$.  In comparison to the other late M dwarfs identified in \\citet{giz00}, 2MASS 1707-0558 has one of the lowest tangential velocities and H$\\alpha$ measurements, suggesting that 2MASS 1707-0558 is associated with the young population.  It is therefore likely to be closer to $\\sim$1~Gyr in age, but without additional constraints, we adopt a conservative age range of 0.5-5~Gyr.  Figure 5 compares the properties of the 2MASS 1707-0558 components to evolutionary models from \\citet{burr97}, clearly showing the secondary's possible location below the Li-burning mass limit."
    },
    "0608/astro-ph0608469_arXiv.txt": {
        "abstract": "A new class of accelerating cosmological models driven by a one-parameter version of the general Chaplygin-type equation of state is proposed. The simplified version is naturally obtained from causality considerations with basis on the adiabatic sound speed $v_S$ plus the observed accelerating stage of the universe. We show that very stringent constraints on the unique free parameter $\\alpha$ describing the simplified Chaplygin model can be obtained from a joint analysis involving the latest SNe type Ia data and the recent Sloan Digital Sky Survey measurement of baryon acoustic oscillations (BAO). In our analysis we have considered separately the SNe type Ia gold sample measured by Riess et al. (2004) and the Supernova Legacy Survey (SNLS) from Astier et al. (2006). At 95.4\\% (c.l.), we find for BAO + \\emph{gold} sample, $0.91 \\leq \\alpha \\leq 1.0$ and $\\Omega_{\\rm{M}}= 0.28^{+0.043}_{-0.048}$ while BAO + SNLS analysis provides $0.94 \\leq \\alpha \\leq 1.0$ and $\\Omega_{\\rm M}=0.27^{+0.048}_{-0.045}$. ",
        "introduction": "The impressive convergence of recent observational facts along with some apparently successful theoretical predictions seem to indicate that the simple approach provided by the standard cold dark matter (CDM) model is insufficient to describe the present stage of our universe. From these results, the most plausible picture for our world seems to be a nearly flat scenario dominated basically by CDM and an exotic component endowed with large negative pressure, usually named dark energy. Despite the good observational indications for the existence of these two components, their physical properties constitute a completely open question at present, which gives rise to the so-called dark matter and dark energy problems (see \\cite{rev123} for a recent review on this topic).  Among the many candidates for the dark energy component, a very interesting one was suggested by Kamenshchik {\\it et al.} \\cite{kamen} and developed by Bili\\'c {\\it et al.} \\cite{bilic} and Bento {\\it et al.} \\cite{bento}. Such an exotic fluid, named generalized Chaplygin gas (C-gas), can be macroscopically characterized  by the equation of state (EoS) \\begin{equation}\\label{eq1} p_C = -A/\\rho_C^{\\alpha}, \\end{equation} where $\\alpha$ = 1 and $A$ is a positive constant related to the present-day Chaplygin adiabatic sound speed, $v^2_s = \\alpha A/\\rho_{C_{o}}^{1 + \\alpha}$ ($\\rho_{C_{o}}$ stands for the current C-gas density). In actual fact, the above equation for $\\alpha \\neq 1$ constitutes a generalization of the original C-gas EoS proposed by Bento {\\it et al.} in Ref. \\cite{bento}.   In the last few years, the possibility of describing the unknown dark energy component using the C-gas-type EoS above has provoked a considerable debate in the literature. Theoretical connections between the C-gas and string theory, supersymmetric generalizations \\cite{hope,jackiw}, self-interacting \\cite{zhu}, and even a tachyonic fluid representation \\cite{Bena02} has also been investigated.  Another interesting feature of the above EoS comes from the fact that the C-gas becomes pressureless at high redshifts, which suggests a possible unification scheme for the cosmological ``dark sector\", an interesting idea which has  been considered in different contexts \\cite{quartessence}.  Observational aspects of the above C-gas scenarios have also been largely investigated in the literature. Cosmological tests involving type Ia supernovae (SNe Ia) data \\cite{fabris,alcaniz}, the shape of the matter power spectrum \\cite{avelino}, statistical properties of gravitational lenses \\cite{dev}, the age of the Universe \\cite{jailson},  cosmic microwave background (CMB) measurements \\cite{bento1,bert1,fin,finelli}, galaxy clusters X-ray \\cite{CAL04}, and gamma-ray bursts data \\cite{bertgr}  have been discussed. In general, to perform such analyses, besides the present value of the C-gas density parameter ($\\Omega_C$), the above barotropic EoS implies that one needs to constrain two additional free parameters, namely, $A$ and $\\alpha$. Therefore, in the context of the Friedman-Robertson-Walker (FRW) cosmologies with CDM plus a C-gas, there are at least 4 parameters to be constrained by the data. Actually, this number can be reduced to 3 if one assumes a flat geometry, i.e., $\\Omega_{\\rm{M}}  = 1 - \\Omega_C$ or if a unified dark matter/energy picture involving only the C-gas and baryons is assumed from the very beginning (in this case, the baryonic density ($\\Omega_b$) may be fixed a priori by using, for instance, nucleosynthesis \\cite{Steigmann} or the recent Cosmic Microwave Background (CMB) observations \\cite{wmap}). However, even in this latter cases, there are so many parameters to be constrained by the data, that a high degree of degeneracy on the parametric space becomes inevitable.   Many generalizations of the  original C-gas \\cite{ZhWuZh06,ChLa05,MaHa05,Wang05,GuZh05}, or even of its extended version \\cite{SeSc05} have appeared in literature. In these cases, the number of free parameters is usually increased, and, as consequence, the models become mathematically richer although much less predictive from a physical viewpoint. In this work by following the opposite direction, we propose a simplified version for the generalized C-gas-type EoS which diminishes one of its free parameters. By an additional physical condition, the allowed range of the remaining parameter is also restricted a priori, which makes not only the relevant parametric space bi-dimensional but also (and more important) the model more easily discarded or confirmed by the present set of observations since the range of its free parameter is physically limited from causality considerations. We test the viability of this simplified C-gas approach by discussing the constraints imposed from current SNe Ia observations and Large Scale Structure (LSS) data. ",
        "conclusions": "As widely known, there are many theoretical approaches for describing the exotic dark energy component accelerating the Universe. However, until the present, the available battery of cosmological tests was not capable to decide which is the best theoretical representation. We have argued here that one of such candidates, the so-called Chaplygin type gas (whose equation of state depends on two parameters $A_s$ and $\\alpha$), may have a very simplified description. We postulate that $A_s$ is a function of $\\alpha$ and for simplicity we have taken $A_s = \\alpha$. Thus, similarly to the concordance model ($\\Lambda$CDM), the resulting flat cosmology is completely described only by a pair of parameters ($\\alpha$, $\\Omega_{\\rm M}$).  This SC-gas cosmology mimics the dynamics of the X-matter models  with an extra bonus, namely: the fluid stability and other thermodynamic features are guaranteed from the very beginning.  By considering this particular parameterization we have investigated constraints on the $\\alpha$ parameter from the most recent SNe Ia (\\emph{gold} and SNLS samples) and LSS data. We have found that the limits arising from this particular combination of the data are much more restrictive on this simplified approach than on the generalized C-gas version. In particular, for the the BAO+SNLS combination we found $0.94 \\leq \\alpha \\leq 1.0$ and $\\Omega_{\\rm M}=0.27^{+0.048}_{-0.045}$ (at 95\\% c.l.), which is in agreement with recent estimates of the clustered matter. Naturally, it should be interesting to investigate whether current CMB data and other independent observations can or cannot discard the simplified scenario proposed here."
    },
    "0608/astro-ph0608696_arXiv.txt": {
        "abstract": "{} {We report the X-ray properties of young ($\\sim$3\\,Myr) \\bona~brown dwarfs of the Taurus Molecular Cloud (TMC).} {The \\xest~(XEST) is a large program designed to systematically investigate the X-ray properties of young stellar/substellar objects in the TMC. In particular, the area surveyed by 15 \\xmm~pointings (of which three are archival observations), supplemented with one archival \\cxo~observation, allows us to study 17 brown dwarfs with M spectral types. } {Half of this sample (9 out of 17 brown dwarfs) is detected; 7 brown dwarfs are detected here for the first time in X-rays. We observed a flare from one brown dwarf. We confirm several previous findings on brown dwarf X-ray activity: a log-log relation between X-ray and bolometric luminosity for stars (with $L_*$$\\le$10\\,L$_\\odot$) and brown dwarfs detected in X-rays, which is consistent with a mean X-ray fractional luminosity $<\\!\\log(L_{\\rm X}/L_*)\\!>\\,=-3.5 \\pm 0.4$; for the XEST brown dwarfs, the median of $\\log(L_{\\rm X}/L_*)$ (including upper limits) is $-4.0$; a shallow log-log relation between X-ray fractional luminosity and mass; a log-log relation between X-ray fractional luminosity and effective temperature; a log-log relation between X-ray surface flux and effective temperature. We find no significant log-log correlation between the X-ray fractional luminosity and $EW({\\rm H}\\alpha)$. Accreting and nonaccreting brown dwarfs have a similar X-ray fractional luminosity. The median X-ray fractional luminosity of nonaccreting brown dwarfs is about 4 times lower than the mean saturation value for rapidly rotating low-mass field stars. Our TMC brown dwarfs have higher X-ray fractional luminosity than brown dwarfs in the \\coup. } {The X-ray fractional luminosity declines from low-mass stars to M-type brown dwarfs, and as a sample, the brown dwarfs are less efficient X-ray emitters than low-mass stars. We thus conclude that while the brown dwarf atmospheres observed here are mostly warm enough to sustain coronal activity, a trend is seen that may indicate its gradual decline due to the drop in photospheric ionization degree.} ",
        "introduction": "\\begin{figure*}[!t] \\centering \\includegraphics[width=1.8\\columnwidth]{5559fig1.ps} \\caption{The {\\it XMM-Newton Extended Survey of the Taurus Molecular Cloud}~(XEST). Contours show the $^{12}$CO emission \\citep{dame87} of the Taurus Molecular Cloud (TMC), overlaid on a visual extinction map \\citep[linear colour scale, with black colour corresponding to $A_{\\rm V}\\sim6$\\,mag; based on DSS~I, {6\\arcmin}-resolution map,][]{dobashi05}. The 27 \\xmm~fields of view of the XEST are plotted with continuous/dashed circles (of which seven are archival observations and one is a separate program on T\\,Tau; dashed circles). Note the outlying \\xmm~fields of view around SU\\,Aur (NE corner) and L1551 (S). The two squares near labels 3 and 8 show the archival \\cxo~fields of view used in this work (see Table~\\ref{table:log}). The continuous line shows the survey of brown dwarfs (BDs) performed with the Canada-France-Hawaii telescope \\citep{guieu06}. This region hosts 42 young BDs of the TMC (\\citealt{briceno98,briceno02}; \\citealt{luhman00,luhman04,luhman06}; \\citealt{martin01}; \\citealt{guieu06}). Black dots indicate the 25 BDs of this region which were not surveyed in X-rays with the XEST. White dots show the 8 BDs not detected in X-rays. Crosses show the 9 BDs detected in X-rays (Table~\\ref{table:detections}) with labeled numbers referring to BD names. Only 2 BDs of TMC were previously detected in X-rays by \\rosat: MHO\\,4 \\citep{carkner96,neuhaeuser99} and CFHT-BD-Tau\\,4 \\citep{briceno99,mokler02}. } \\label{fig:map} \\end{figure*}  Pre-main-sequence low-mass stars, i.e.\\ T~Tauri stars, show a high level of X-ray emission, which is generally attributed to an active corona, an enhanced version of the magnetic corona on the contemporary Sun \\citep[see for a review][]{feigelson99}. The dynamo mechanism producing the magnetic field in these fully convective stars is still discussed \\citep{preibisch05b}. X-ray observations of brown dwarfs (BDs) allow investigation of the magnetic activity in the substellar regime, where masses are lower than about 0.075\\,M$_\\odot$.  The first X-ray detection of a BD was reported by \\citet{neuhaeuser98}, who identified in the Cha\\,I dark cloud an X-ray source detected by \\rosat~-- in early 1991, before the near-infrared observations of the first BDs were reported \\citep{nakajima95,rebolo95} -- as the counterpart of a young BD. \\rosat~detected only a few BDs or very low-mass stars, all of them young and located in star-forming regions \\citep{neuhaeuser99,comeron00,mokler02}.  The new generation of X-ray satellites, \\xmm~and \\cxo, allows now detection of more BDs thanks to their increase of sensitivity. The number of young or intermediate-age BDs detected in X-rays \\citep[e.g.,][]{imanishi01,preibisch01,preibisch02,tsuboi03,briggs04,stelzer04b,ozawa05,stelzer06} exceeds largely the one of (older) field BDs, which have been detected so far mainly during X-ray flares \\citep{rutledge00,stelzer04}. Recently, the X-ray properties of 33 young BDs \\citep{slesnick04} in the Orion Nebula Cluster (ONC) located at $\\sim$450\\,pc were studied by \\citet{preibisch05} in the {\\it Chandra Orion Ultradeep Project} (COUP).  \\begin{table*}[!ht] \\caption{The sample of the 17 young BDs surveyed in the XEST. Col.~(2) gives the reference of the discovery paper. Col.~(3) give the 2MASS counterparts. Col.~(7) gives the references for the spectral type (Col.~4), the optical extinction \\citep[when not available derived from $A_{\\rm J}$ using][ Col.~5]{rieke85}, and the reference of the equivalent width of H$\\alpha$ ($EW({\\rm H}\\alpha)$, given in Col.~6) if it is different. Negligible optical extinction in Col.~(5) are indicated by 0. The effective temperature in Col.~(8) has been computed from the spectral type using the temperature scale $T_{\\rm eff}=3841.94-141.17 \\times {SpTyp}$, which is valid for young stars with M spectral type \\citep{guieu06}. The visual extinction in Col.~(5) is taken when available from the literature; for 2MASS\\,J0421 and 2MASS\\,J0422, we averaged the visual extinction computed from $J-H$, $H-K_{\\rm S}$, and the spectral type, using a dwarf sequence \\citep[compiled from the literature; e.g.,][]{leggett98}. The bolometric luminosity of the substellar photosphere in Col.~(9) has been computed from $I$, $J$-band magnitudes and $A_{\\rm V}$ \\citep[see][]{guieu06}; for 2MASS\\,J0414 and MHO\\,4 (without $I$-band magnitude available) the reference for the luminosity is \\citet{luhman04} and \\citet{briceno02}, respectively. Col.~(10) indicates accreting sources based mainly on $EW({\\rm H}\\alpha)$ (see \\S\\ref{accretion}). The last column indicates detection in X-rays (this work; see Table~\\ref{table:detections}). References: B98=\\citet{briceno98}; B02=\\citet{briceno02}; G06=\\citet{guieu06}; L04=\\citet{luhman04}; L06=\\citet{luhman06}; M01=\\citet{martin01}; M05=\\citet{muzerolle05b}.} \\label{table:bd} \\centering \\begin{tabular}{lcclcrccccc} \\hline\\hline \\multicolumn{1}{c}{Name} & Ref.\\ & 2MASS & \\multicolumn{1}{c}{SpTyp} & $A_{\\rm V}$ & \\multicolumn{1}{c}{$EW({\\rm H}\\alpha)$} & Ref.\\ & $T_{\\rm eff}$ & $L_*$ &  Acc. & X\\\\ &                      &   &                         & mag &  \\multicolumn{1}{c}{$\\AA$}  & & K & L$_\\odot$   \\\\ \\multicolumn{1}{c}{(1)} & (2) & (3) & \\multicolumn{1}{c}{(4)} & (5) & \\multicolumn{1}{c}{(6)} & (7) & (8) & (9) & (10) & (11)\\\\ \\hline \\object{2MASS\\,J0414}   & L04 & J04141188+2811535 &  M6.25 & 1.1 & 250.0    &  M05 & 2960 & 0.015 & y & n\\\\ \\object{KPNO-Tau\\,2}    & B02 & J04185115+2814332 &  M6.75 & 0.4 & 8.4\t   &  G06 & 2889 & 0.007 & n & n\\\\ \\object{2MASS\\,J0421}   & L06 & J04215450+2652315 &  M8.5  & 3.0 & \\dotfill &  L06 & 2642 & 0.003 &\\dotfill & n\\\\ \\object{2MASS\\,J0422}   & L06 & J04221332+1934392 &  M8    & 1.0 & \\dotfill &  L06 & 2713 & 0.017 &\\dotfill & y\\\\ \\object{KPNO-Tau\\,4}    & B02 & J04272799+2612052 &  M9.5  & 2.5 & 158.1    &  G06 & 2501 & 0.004 & y & n\\\\ \\object{KPNO-Tau\\,5}    & B02 & J04294568+2630468 &  M7.5  & 0\t& 30.0     &  B02 & 2783 & 0.023 & n & y\\\\ \\object{KPNO-Tau\\,6}    & B02 & J04300724+2608207 &  M9    & 0.9 & 207.9    &  G06 & 2571 & 0.003 & y & n\\\\ \\object{KPNO-Tau\\,7}    & B02 & J04305718+2556394 &  M8.25 & 0\t& 300.0    &  B02 & 2677 & 0.004 & y & n\\\\ \\object{MHO\\,4}         & B98 & J04312405+1800215 &  M7    & 0.5 & 42.0     &  B02 & 2854 & 0.048 & n & y\\\\ \\object{CFHT-Tau\\,5}    & G06 & J04325026+2422115 &  M7.5  & 9.2 & 29.8     &  G06 & 2783 & 0.075 & n & y\\\\ \\object{CFHT-BD-Tau\\,1} & M01 & J04341527+2250309 &  M7    & 3.1 & 19.0     &  M01 & 2854 & 0.017 & n & y\\\\ \\object{KPNO-Tau\\,9}    & B02 & J04355143+2249119 &  M8.5  & 0\t& 20.0     &  B02 & 2642 & 0.001 & n & n\\\\ \\object{CFHT-BD-Tau\\,2} & M01 & J04361038+2259560 &  M7.5  & 0\t& 13.0     &  B02 & 2783 & 0.007 & n & n\\\\ \\object{CFHT-BD-Tau\\,3} & M01 & J04363893+2258119 &  M7.75 & 0\t& 55.0     &  B02 & 2747 & 0.007 & n & y\\\\ \\object{CFHT-Tau\\,6}    & G06 & J04390396+2544264 &  M7.25 & 0.4 & 63.7     &  G06 & 2818 & 0.024 & y & y\\\\ \\object{CFHT-BD-Tau\\,4} & M01 & J04394748+2601407 &  M7    & 2.6 & 340.0&  L04,M01 & 2854 & 0.062 & y & y\\\\ \\object{2MASS\\,J0455}   & L04 & J04552333+3027366 &  M6.25 & 0\t& \\dotfill &  L04 & 2960 & 0.015 &\\dotfill & y\\\\ \\hline \\end{tabular} \\end{table*}  We investigate in this paper the X-ray properties of a sample of 17 BDs in the nearby (140\\,pc) Taurus Molecular Cloud (TMC), surveyed in the {\\it XMM-Newton Extended Survey of the TMC} \\citep[XEST;][]{guedel06b}. First, we define in \\S\\ref{TMC_BD} the TMC BD sample, review previous X-ray results on it, and report the BDs detected in the XEST. Analysis of the X-ray variability of the TMC BDs is presented in \\S\\ref{lc_analysis}, and X-ray spectral properties are determined in \\S\\ref{spectral_properties}. In \\S\\ref{halpha}, we compare the X-ray luminosities of the BDs and low-mass stars observed in the XEST. Comparison of the X-ray fractional luminosities of BDs in the XEST and COUP is made in \\S\\ref{XEST_COUP}. We discuss the origin of the BD X-ray emission in the broader context of the X-ray emission of cool stars in \\S\\ref{discussion}. Finally, we summarize our results in \\S\\ref{summary}. ",
        "conclusions": "\\label{discussion}  We discuss now the origin of the BD X-ray emission in the broader context of the X-ray emission of cool stars. Following \\citet{preibisch05}, we compare the X-ray fractional luminosities of TMC BDs with other cool objects.  Fig.~\\ref{fig:eta_spectyp} shows the X-ray fractional luminosities of TMC objects with spectral type equal or later than M5, X-ray detected BDs \\citep[][ and references therein; Stelzer et al.\\ 2006]{preibisch05}, the evolved BD Gl\\,569 Ba,b \\citep{stelzer04}, and the field BD LP\\,944-20 \\citep{rutledge00,martin02}. Cool field stars with spectral type M5 or later are also indicated: M field stars from \\citet{fleming93}; very-low mass field stars VB\\,10 \\citep{fleming03}, LHS\\,2065 \\citep{schmitt02}, HC\\,722 \\citep{slesnick04,slesnick05,preibisch05}. All these cool objects show a similar level of X-ray fractional luminosities ranging from $\\sim$$10^{-3}$ to $\\sim$$10^{-4}$ (note that, compared to Fig.~\\ref{fig:eta_teff}, the spectral type range has been reduced by a factor of two in Fig.~\\ref{fig:eta_spectyp}).  Based on \\citet{baraffe98} evolutionary tracks, objects of spectral type M7 with an age of 1\\,Gyr are not BDs, but low-mass stars twice as massive as a typical TMC BD having an M7 spectral type and an age of 3\\,Myr. Moreover, such very cool stars also have surface gravities about 40 times higher than in a typical TMC BD. This shows, as found by \\citet{preibisch05} from the COUP BDs, that the X-ray activity of BD coronae is not strongly dependent of the BD mass and the BD surface gravity. Therefore, this implies that the relation that we have obtained, for the young objects of TMC with spectral type M, between X-ray fractional luminosity and mass (Fig.~\\ref{fig:eta_mass}) is in fact the consequence of the more fundamental relation between X-ray fractional luminosity and effective temperature (Fig.~\\ref{fig:eta_teff}). This latter relation agrees with the overall result (field dwarfs and young BDs) in Fig.~\\ref{fig:eta_spectyp}: of the 15 sources shown with spectral types M8.5V or later, only 4 have any detected quiescent emission; the rest are either not detected at all or (in 3 cases) detected only during strong flares. By considering a subsample of evolved BDs (from $\\sim$1\\,Myr to $\\sim$1\\,Gyr) with a small range of masses (0.05--0.07\\,M$_\\odot$) but with effective temperatures ranging from 3000\\,K down to 1000\\,K, \\citet{stelzer06} found evidence for a similar decline of the fractional X-ray luminosity with the effective temperature; showing that the atmospheric temperature plays a crucial role in determining the level of X-ray activity.  The growing evidence for a decline of the coronal activity with the effective temperature is analogous to the decline in chromospheric (H$\\alpha$) activity seen in field dwarfs at similar spectral types, which show a slow decline from mid- to late M in $L_{\\rm H\\alpha}/L_{\\rm *}$ and a sharp drop-off around M9V \\citep{mohanty03}. Although a definitive turn-over of the fractional X-ray luminosity toward L-type BDs is not proven for the present sample due to sparse statistics, we note that the fractional luminosity declines from low-mass stars to M-type BDs, and as a sample, the BDs are less efficient X-ray emitters than low-mass stars. We thus conclude that while the BD atmospheres observed here are mostly warm enough to sustain coronal activity, a trend is seen that may indicate its gradual decline due to the drop in photospheric ionization degree \\citep{mohanty03}."
    },
    "0608/astro-ph0608375_arXiv.txt": {
        "abstract": "We study the formation of giant dense cloud complexes and of stars within them by means of SPH numerical simulations of the collision of gas streams (``inflows'') in the warm neutral medium (WNM) at moderately supersonic velocities. The collisions cause compression, cooling and turbulence generation in the gas, forming a cloud that then becomes self-gravitating and begins to collapse globally. Simultaneously, the turbulent, nonlinear density fluctuations induce fast, local collapse events. The simulations show that: a) The clouds are \\emph{not} in a state of equilibrium. Instead, they undergo secular evolution. During the early stages of the evolution, their mass and gravitational energy $\\Eg$ increase steadily, while the turbulent energy $\\Ek$ reaches a plateau. b) When $\\Eg$ becomes comparable to $\\Ek$, global collapse begins, causing a simultaneous increase in $|\\Eg|$ and $\\Ek$ that maintains a near-equipartition condition $|\\Eg| \\sim 2 \\Ek$. c) Longer inflow durations delay the onset of global and local collapse, by maintaining a higher turbulent velocity dispersion in the cloud over longer times. d) The star formation rate is large from the beginning, without any period of slow and accelerating star formation. e) The column densities of the local star-forming clumps are very similar to reported values of the column density required for molecule formation, suggesting that locally molecular gas and star formation occur nearly simultaneously. The MC formation mechanism discussed here naturally explains the apparent ``virialized'' state of MCs and the ubiquitous presence of HI halos around them. Also, within their assumptions, our simulations support the scenario of rapid star formation \\emph{after} MCs are formed, although long ($\\gtrsim$ 15 Myr) accumulation periods do occur during which the clouds build up their gravitational energy, and which are expected to be spent in the atomic phase. ",
        "introduction": "\\label{sec:intro}  The evolution of molecular clouds (MCs) remains an unsolved problem to date. While traditionally MCs have been thought of as virialized structures in the interstellar medium (ISM) \\citep[e.g., ][]{dBD80, McKee_etal93, BW99, Mc99}, with relatively long lifetimes \\citep{BS80} and a significant delay before they begin forming stars \\citep[e.g.][]{PS00, PS02, TM04, MTK06}, recent studies have suggested that MCs begin forming stars shortly after they themselves form, and are non-equilibrium entities \\citep[][ see also the reviews by Mac Low \\& Klessen 2004 and \\BP\\ et al.\\ 2006]{BVS99, BHV99, Elm00, HBB01, VBK03, PaperI, BH06, Ball06}. One approach that can shed light on this problem is to perform numerical simulations of the MC formation process by generic compressions in the warm neutral medium (WNM). Numerous studies of this kind exist, although self-gravity has in general not been included \\citep[][ hereafter Paper I]{HP99, HP00, KI00, KI02, IK04, AH05, Hetal05, Hetal06, PaperI}.  These studies have shown that a substantial fraction of the internal turbulence, with velocity dispersions of up to a few$\\kms$, can be produced {in shock-bounded layers} by a combined dynamical+thermal instability, whose nature is still not well determined \\citep{Hunter_etal86, WF98, WF00, KI02, IK04, Hetal05, Hetal06, FW06}. Paper I also showed that the dense gas, there defined as the gas with densities larger than $100 \\pcc$, is at systematically higher thermal pressures than the mean interstellar value by factors of 1.5--5, implying that MC formation by colliding streams of diffuse gas can account at least partially for the clouds' excess pressure and turbulent nature.  Only a few works have included self-gravity in simulations of colliding flows. The pioneering work of \\citet{Hunter_etal86} considered the collision of supersonic streams within MCs, finding that the fragments produced by the dynamical instability (which they termed ``Rayleigh-Taylor-like'') were able to accrete mass until they became gravitationally unstable and collapse. However, this study was designed to investigate flows \\emph{within} MCs rather than the formation of the clouds themselves, and so it considered much smaller (sub-parsec) scales and different density, temperature and cooling regimes than are relevant for the process of MC formation.  More recent numerical studies of MC formation not including self-gravity \\citep{KI02, Hetal05} have estimated that the individual clumps formed in the compressed layer formed by the colliding streams are themselves gravitationally stable, although the entire compressed-gas complex would be gravitationally unstable, had self-gravity been included.  This raises the issue of whether the clumps can be considered in isolation.  The question of which gas parcels are involved in the collapse of a particular object is a delicate one, and lies, for example, at the heart of the debate of whether competitive accretion is relevant or not in the process of star formation within molecular clouds: A model in which clumps are isolated dense objects immersed in a more tenuous, unbound medium leads to the conclusion that competitive accretion is irrelevant \\citep{KMK05}. Instead, competitive accretion plays a central role \\citep{BBCP97, KBB98, BBCP01, KB00, KB01, K01a, K01b, BB06} in a scenario in which the local density peaks are the ``tips of the iceberg'' of the density distribution, the flows are organized on much larger scales \\citep{BVS99, BHV99, BBB03}, and the evolution and energetics of the density peaks cannot be considered in isolation, but as part of the global flow instead \\citep{BVS99, KHM00, HMK01, SVB02, TP04, TP05, Ball06, Dib_etal06}. In a similar fashion, then, accurate determination of whether gravitational collapse can be triggered by moderate-Mach number stream collisions in the WNM, requires to self-consistently include self-gravity in the numerical simulations.  {Simulations of molecular cloud formation by the passage of a self-gravitating, clumpy medium through Galactic spiral shocks \\citep{BDRP06, DBP06} have shown that realistic velocity dispersions and densities can be generated in the resulting post-shock clouds, which are furthermore driven to produce local collapse events leading to star formation, although a large fraction of the mass in the clouds remains gravitationally unbound. However, because these simulations were isothermal, the gas had to be assumed to already be clumpy and at very low temperatures ($\\sim$ 100 K) \\emph{previous} to the passage through the shock, while recent studies including cooling leading to thermal bistability of the atomic gas \\citep[][ Paper I]{KI02, AH05, Hetal05, Hetal06} suggest that the clumpiness is generated by the shock compression, which nonlinearly triggers the thermal instability, induces supersonic turbulence, and causes the formation of dense clumps in the compressed layers.}  It is {thus}  essential to model the evolution of the cloud within the frame of its more diffuse environment {including both self-gravity and cooling}. Not only this is indispensible in order to model {the generation of turbulence and clumpiness, but also to model the cloud's evolution as it exchanges mass and energy with its surroundings. Note that the generation of turbulence by the compression itself implies that the turbulence within the cloud is driven while the large-scale converging motions persist, and decaying afterwards. This is a mixed regime that also requires self-consistent modeling of the cloud and its surroundings, rather than, for example, random Fourier driving in closed boxes.}  {Finally, such a unified (or ``holistic'') description of molecular cloud formation and evolution, as outlined by \\citet{VBK03}, including the star formation epoch, allows also the investigation of a number of key problems, such as the evolutionary and star-formation timescales, the evolution of the cloud's gravitational and kinetic energies, and whether it settles into virial equilibrium or not, and the effect of self-gravity on the cloud's physical conditions.}  In this paper we investigate numerically the evolution of interstellar gas as it collects from the WNM, {shocks, suffers a phase transition to a cold and dense state}, and finally begins to form stars, by means of numerical simulations of colliding gas streams in the WNM, in the presence of self-gravity and cooling leading to thermal instability. Since this is a problem that naturally collects large amounts of gas in small volumes, first by compression-triggered thermal instability and then by gravitational instability, it is convenient to use some Lagrangian numerical scheme, and we opt for a smoothed particle hydrodynamics code, sacrificing the possibility of including magnetic fields. We also omit for now a self-consistent description of the stellar feedback on the gas {and of the chemistry}. Thus, our simulations will not adequately describe the latest stages of MC evolution {(in which stellar energy feedback may significantly influence the dynamics of the cloud) nor are they capable of predicting the transition from atomic to molecular gas accurately. We defer these tasks to subsequent papers}.  The plan of the paper is as follows. In \\S \\ref{sec:quali} we first give a qualitative description of the physical system, the main phenomena involved, and our expectations for the evolution. In \\S \\ref{sec:model} we describe the numerical code, the physical physical setting, and the choice of parameters. Then, in \\ref{sec:results} we describe the results concerning the evolution of the cloud's mass and its kinetic and gravitational energies, and the onset of star formation. In \\S \\ref{sec:discussion} we discuss the implications and limitations of the recent work. Finally, in \\S \\ref{sec:concl} we present a summary and draw some conclusions. ",
        "conclusions": "\\label{sec:concl}  In this paper we have presented a suite of numerical simulations designed to investigate jointly the formation of molecular clouds (MCs) and of stars within them. The simulations use an SPH scheme including self-gravity, sink particles and cooling {leading to thermal bistability}. Magnetic fields are neglected. The simulations describe the collision of oppositely-directed gas streams (``inflows'') in the warm neutral medium (WNM) at moderately supersonic velocities (each with a velocity of $9.2 \\kms$, or a Mach number of 1.22 in the unperturbed WNM). Three simulations were considered, varying the length of the inflow and the amplitude of the initial velocity fluctuations in the gas. The collisions trigger a transition to the cold phase in the gas, and simultaneously generate turbulence in the resulting ``cloud'', defined as the gas at densities $n > 50 \\pcc$. The inflows secularly weaken in time, and so does the turbulence level in the cloud, implying that {the turbulence in the clouds gradually transits from being continuously driven to being in a decaying regime}. The cooling and the mass gain of the dense gas eventually cause the cloud to contain a large number of Jeans masses at the mean conditions. Moreover, because the cloud is supersonically turbulent, locally the density can be much larger, with a correspondingly lower Jeans mass and shorter free-fall times than {those} of the whole cloud.  Thus, by the time the inflows have almost subsided, the cloud engages in global gravitational collapse, but shortly thereafter it begins to produce numerous local collapse events that occur on much smaller timescales because of the larger densities, so that by the time the global collapse of the cloud is completed, it has converted most of its mass to collapsed objects (sink particles). The sinks in general represent star clusters, due to the limited mass resolution and lack of modeling of opacity-limited fragmentation \\citep[e.g.][]{BBB03}. Our simulations can only be considered reliable up to the time when the mass in collapsed objects (sinks) implies a high enough number of massive stars that they would disrupt the cloud. Nevertheless, the evolution up to that point shows a number of relevant results:  \\medskip \\noindent 1. The clouds are never in virial equilibrium during this period. Instead, they continually evolve, increasing their mass and self-gravitating energy, until the latter becomes comparable or larger than the thermal and kinetic energies. {Some} $\\sim 5$ Myr {after global contraction began}, star formation begins.  \\medskip \\noindent 2. In spite of not being in equilibrium, the near-equipartition condition $|\\Eg| \\sim 2 \\Ek$ is approximately satisfied as soon as $|\\Eg|$ becomes comparable to $\\Ek$, because of the gravitational contraction of the cloud, with both quantities increasing simultaneously. This occurs long before the onset of star formation. This fact can explain the observed state of apparent virialization of GMCs.  \\medskip \\noindent 3. The near equipartition is a signature of global gravitational collapse, not equilibrium, and suggests a return to Goldreich \\& Kwan's (1974) scenario of global gravitational contraction in MCs.  {However, the criticism by \\citet{ZP74}, namely that an excessively large star formation should result through this process, is avoided in part because the nonlinear turbulent density fluctuations collapse earlier than the whole cloud, involving only a fraction of the total mass, and in part because as soon as the stars form they probably contribute to dispersing the cloud, or at least halting its global collapse. Further reduction of the SFE may occur in the presence of supercritical magnetic fields.}  \\medskip \\noindent 4. Local collapse events begin to occur after global collapse has begun, but they occur rapidly, requiring only $\\sim 3$ Myr to convert $\\sim 15$\\% of the cloud's mass into stars.  \\medskip \\noindent 5. The star formation rate, measured by the slope of the mass in stars \\emph{versus} time, is large from the beginning. Within the framework, assumptions and limitations of our simulations, no {long} period of slow, accelerating star formation is observed.   \\medskip \\noindent 6. Longer inflow durations maintain larger turbulent velocity dispersions in the clouds and delay the onset of both global and local collapse. Instead, larger amplitudes of the initial velocity fluctuations have little effect in delaying the collapse. The latter effect can be attributed to the fact that the turbulence is already dissipating in the clouds by the time global collapse begins.  \\medskip \\noindent 7. A long period ($\\sim 14$--17 Myr) of ``dormancy'' does occur {between the time when the cloud begins to form and the time when} star formation begins. Nevertheless, it is likely that most of this time is spent with the gas being in atomic form, since the column density of star forming regions in our simulations are comparable to values required for molecular gas formation, as reported in the literature. Thus, our simulations support the notion that star formation occurs almost simultaneously with the formation of molecular gas, under local ISM conditions.  \\medskip We conclude that our simulations support the scenario of rapid star formation after molecular gas has formed, involving accumulations of gas from distances of a few hundred parsecs \\citep{HBB01, VBK03}, while simultaneously requiring timescales of the order of the crossing time across the largest scales involved \\citep{Elm00}, and in a systematically out-of-equilibrium fashion \\citep{BVS99, BHV99, KHM00}. However, the evidence will have to be made more compelling as additional physical processes, such as longer-duration inflows, stellar feedback, magnetic fields and chemistry are included."
    },
    "0608/astro-ph0608005_arXiv.txt": {
        "abstract": "The Pierre Auger Observatory detects ultra-high energy cosmic rays by implementing two complementary air-shower techniques. The combination of a large ground array and fluorescence detectors, known as the \\textit{hybrid} concept, means that a rich variety of measurements can be made on a single shower, providing much improved information over what is possible with either detector alone. In this paper the hybrid reconstruction approach and its performance are described. ",
        "introduction": "The Pierre Auger Observatory was designed to observe, in coincidence, the shower particles at ground and the associated fluorescence light generated in the atmosphere. This is achieved with a large array of water Cherenkov detectors coupled with air-fluorescence detectors that overlook the surface array. It is not simply a dual experiment. Apart from important cross-checks and measurement redundancy, the two techniques see air showers in complementary ways. The ground array measures the lateral structure of the shower at ground level, with some ability to separate the electromagnetic and muon components. On the other hand, the fluorescence detector records the longitudinal pro{f}ile of the shower during its development through the atmosphere.  A \\textit{hybrid} event is an air shower that is simultaneously detected by the fluorescence detector and the ground array. The Observatory was originally designed and is currently being built with a \\textit{cross--triggering} capability. Data are recovered from both detectors whenever either system is triggered. If an air shower independently triggers both detectors the event is tagged accordingly. An example of this type of events, known as \\textit{golden} hybrids, is shown in Fig.~\\ref{fig:golden}.  \\begin{figure}[!ht] \\centerline{% \\includegraphics[width=0.4\\textwidth]{mostafa_2006_01_fig01.eps}% \\includegraphics[width=0.4\\textwidth]{mostafa_2006_01_fig02.eps}} \\caption{Example of an air-shower that independently triggers the fluorescence and surface detectors. \\textit{Left:} Fluorescence detector view. The pixel color indicates the timing sequence. The full line is the fitted shower-detector plane. The red squares represent the surface detectors. \\textit{Right:} Surface detector view. The size of the hollow circles is proportional to the logarithm of the signal in each tank. The dotted line is the intersection between the shower-detector plane and the horizontal. The full red line is the projection of the shower axis, and the star marks the location of the shower core.} \\label{fig:golden} % \\end{figure}  There are also cases where the fluorescence detector, having a lower energy threshold, promotes a sub--threshold array trigger. Surface stations are matched by timing and location. This is an important capability because these sub--threshold hybrid events would not have triggered the array otherwise. One example is shown in Fig.~\\ref{fig:subt}.  \\begin{figure}[!hb] \\centerline{% \\includegraphics[width=0.4\\textwidth]{mostafa_2006_01_fig03.eps}% \\includegraphics[width=0.4\\textwidth]{mostafa_2006_01_fig04.eps}} \\caption{Example of a \\textit{sub-threshold} hybrid event. The trigger from the fluorescence detector is matched in time and location by only one of the water tanks. In the array view (\\textit{right}) the fluorescence detector is shown as a black semi-circle.} \\label{fig:subt} % \\end{figure}  The Observatory started operation in \\textit{hybrid} production mode in January, 2004. Surface stations have a $100$\\% duty cycle, while fluorescence eyes can only operate on clear moonless nights. Both surface and fluorescence detectors have been running simultaneously $14\\%$ of the time. The number of hybrid events represents $10\\%$ the statistics of the surface array data. ",
        "conclusions": "We have tested the performance of Pierre Auger Observatory in its hybrid con{f}iguration. Operation started in January, 2004 and over $40000$ hybrid events have been successfully reconstructed up to now. It is important to note that the Observatory is under construction and that results are preliminary. It is already clear that the combination of fluorescence and ground array measurements provides reconstruction of the geometry of the shower with much greater accuracy than is achieved with either detector system on its own. Unprecedented core location and direction precision leads to excellent energy and shower development measurements. Several Auger measurements % pro{f}it from the hybrid capability of the Observatory, including the angular resolution~\\cite{Bonifazi:2005ns}, energy spectrum~\\cite{Sommers:2005vs}, and photon limits~\\cite{Risse:2005hi}."
    },
    "0608/astro-ph0608233_arXiv.txt": {
        "abstract": "Two-dimensional magnetohydrodynamic simulations are performed using the ZEUS-2D code to investigate the dynamics of a collapsar that generates a GRB jet, taking account of realistic equation of state, neutrino cooling and heating processes, magnetic fields, and gravitational force from the central black hole and self gravity. It is found that neutrino heating processes are not so efficient to launch a jet in this study. It is also found that a jet is launched mainly by $B_\\phi$ fields that are amplified by the winding-up effect. However, since the ratio of total energy relative to the rest mass energy in the jet is not so high as several hundred, we conclude that the jets seen in this study are not be a GRB jet. This result suggests that general relativistic effects, which are not included in this study, will be important to generate a GRB jet. Also, the accretion disk with magnetic fields may still play an important role to launch a GRB jet, although a simulation for much longer physical time ($\\sim 10-100$ s) is required to confirm this effect. It is shown that considerable amount of $\\rm ^{56}Ni$ is synthesized in the accretion disk. Thus there will be a possibility for the accretion disk to supply sufficient amount of $\\rm ^{56}Ni$ required to explain the luminosity of a hypernova. Also, it is shown that neutron-rich matter due to electron captures with high entropy per baryon is ejected along the polar axis. Moreover, it is found that the electron fraction becomes larger than 0.5 around the polar axis near the black hole by $\\nu_e$ capture at the region. Thus there will be a possibility that $r$-process and $r/p-$process nucleosynthesis occur at these regions. Finally, much neutrons will be ejected from the jet, which suggests that signals from the neutron decays may be observed as the delayed bump of the light curve of the afterglow or gamma-rays. ",
        "introduction": "There has been growing evidence linking long gamma-ray bursts (GRBs; in this study, we consider only long GRBs, so we refer to long GRBs as GRBs hereafter for simplicity) to the death of massive stars. The host galaxies of GRBs are star-forming galaxies and the positions of GRBs appear to trace the blue light of young stars~\\citep{vreeswijk01,bloom02,gorosabel03}. Also, 'bumps' observed in some afterglows can be naturally explained as contribution of bright supernovae~\\citep{bloom99,reichart99,galama00,garnavich03}. Moreover, direct evidence of some GRBs accompanied by supernovae have been reported such as the association of GRB 980425 with SN 1998bw~\\citep{galama98,iwamoto98}, that of GRB 030329 with SN 2003dh~\\citep{hjorth03,price03,stanek03}, and that of GRB 060218 and SN 2006aj~\\citep{mirabal06,mazzali06}.  It should be noted that these supernovae (except for SN 2006aj) are categorized as a new type of supernovae with large kinetic energy ($\\sim 10^{52}$ ergs), nickel mass ($\\sim 0.5M_{\\odot}$), and luminosity~\\citep{iwamoto98,woosley99}, so these supernovae are sometimes called hypernovae. The total explosion energy of the order of $10^{52}$ erg is too important to be emphasized, because it is generally considered that a normal core-collapse supernova cannot cause such an energetic explosion. Thus another scenario has to be considered to explain the system of a GRB associated with a hypernova. One of the most promising scenarios is the collapsar scenario~\\cite{woosley93}. In the collapsar scenario, a black hole is formed as a result of gravitational collapse. Also, rotation of the progenitor plays an essential role. Due to the rotation, an accretion disk is formed around the equatorial plane. On the other hand, the matter around the rotation axis freely falls into the black hole. MacFadyen and Woosley (1999) pointed out that the jet-induced explosion along the rotation axis may occur due to the heating through neutrino anti-neutrino pair annihilation that are emitted from the accretion disk (see also Fryer and M$\\rm \\acute{e}$sz$\\rm \\acute{a}$ros 2000).  It is true that the collapsar scenario is the breakthrough on the problem of the central engine of GRBs. However, there are many effects that have to be involved in order to establish the scenario firmly. First of all, neutrino heating effects have to be investigated carefully by including microphysics of neutrino processes in the numerical simulations. It is true that MacFadyen and Woosley (1999) have done the numerical simulations of the collapsar, in which a jet is launched along the rotation axis, but detailed microphysics of neutrino heating is not included in their simulations. Secondly, it was pointed out that effects of magnetic fields and rotation may play an important role to launch the GRB jets~\\citep{proga03,mizuno04a,mizuno04b,proga05,shibata05,sekiguchi05,fujimoto06}, although neutrino heating effects are not included in their works. Recently, Rockefeller et al (2006) presented 3-dimensional simulations of collapsars with smoothed particle hydrodynamics code. In their study, 3-flavor flux-limited diffusion package is used to take into account neutrino cooling and absorption of electron-type neutrinos, although neutrino anti-neutrino pair annihilation is not included. They have shown that alpha-viscosity drives energetic explosion through 3-dimensional instabilities and angular momentum transfer, although the jet is not launched and magnetic fields (source of the viscosity) are not included in their study. Thus it is not clear which effects are most important to launch a GRB jet, that is, what process is essential as the central engine of GRBs.  Due to the motivation mentioned above, we have performed two-dimensional magnetohydrodynamic simulations of collapsars with magnetic fields, rotation, and neutrino cooling/heating processes. In our simulations, the realistic equation of state (EOS) of Blinnikov et al. (1996) and effects of photo-disintegration of nuclei are also included. We investigated influence of magnetic fields on the dynamics of collapsars by changing initial amplitude of the magnetic fields. In section~\\ref{method}, models and numerical methods are explained. Results are shown in section~\\ref{results}. Discussions are described in~\\ref{discussions}. Summary and conclusion are presented in section~\\ref{conclusion}. ",
        "conclusions": "We have performed two-dimensional magnetohydrodynamic simulations by the ZEUS-2D code to investigate the dynamics of a collapsar that generates a GRB jet, taking account of realistic equation of state (contribution of electrons, positrons, radiation, and ideal gas of nuclei), neutrino cooling and heating processes, magnetic fields, and gravitational force from the central black hole and self gravity.  We have found that neutrino heating processes (neutrino and anti-neutrino pair annihilation, and $\\nu_e$ and $\\bar{\\nu}_e$ captures on free nucleons) are not so efficient to launch a jet. We have found that a jet is launched by magnetic fields (in particular, $B_\\phi$ fields that are amplified by the winding-up effect). However, the ratio of total energy relative to the rest mass energy in the jet at the final stage of simulations suggests that the bulk Lorentz factor of the jet will not reach to as high as several hundred, so we conclude the jet seen in this study will not be a GRB jet. We also found that the mass accretion rate seems to be enhanced in the models with magnetic fields. This might be because angular momentum is efficiently transfered by the viscosity due to the magnetic fields and multi-dimensional flow.  Since GRB jets are not obtained in this study, we consider that general relativistic effects, by which the efficiency of energy deposition through weak interactions will be enhanced and rotation energy of the black hole will be transfered to matter through the magnetic fields, will be important to generate a GRB jet. Thus we are planning to develop a general relativistic magnetohydrodynamics code in the very near future. Also, the accretion disk with magnetic fields may still play an important role to launch a GRB jet, and it may be seen if we can perform numerical simulations for much longer physical time (say, $\\sim 10-100$ s). To realize such a simulation, the special relativistic code will be helpful because the Alfv\\'{e}n velocity is limited to the speed of light.  We have shown that considerable amount of $\\rm ^{56}Ni$ is synthesized in the accretion disk. Thus if some fraction of the synthesized $\\rm ^{56}Ni$ is ejected without falling into the black hole, there will be a possibility for the accretion disk to supply sufficient amount of $\\rm ^{56}Ni$ required to explain the luminosity of a hypernova. Also, we have shown that neutron-rich matter with high entropy per baryon is ejected along the rotation axis. This is because electron capture dominates positron capture. Moreover, we found that the electron fraction becomes larger than 0.5 around the polar axis near the black hole. This is because $\\nu_e$ capture dominates $\\bar{\\nu}_e$ capture at the region. Thus there will be a possibility that $r$-process and $r/p-$process nucleosynthesis occur at these regions. Finally, much neutrons will be ejected from the jet, which suggests that signals from the neutron decays may be observed as the delayed bump of the light curve of the afterglow or gamma-rays."
    },
    "0608/astro-ph0608227_arXiv.txt": {
        "abstract": "We present measurements of ages, metallicities, and $\\afe$ ratios for 16 globular clusters (GC) in NGC~147, NGC~185, and NGC~205 and of the central regions of the diffuse galaxy light in NGC~185, and NGC~205. Our results are based on spectra obtained with the SCORPIO multi-slit spectrograph at the 6-m telescope of the {\\it Russian Academy of Sciences}. We include in our analysis high-quality HST/WFPC2 photometry of individual stars in the studied GCs to investigate the influence of their horizontal branch (HB) morphology on the spectroscopic analysis. All our sample GCs appear to be old ($T>8$ Gyr) and metal-poor ($\\zh \\la -1.1$), except for the GCs Hubble~V in NGC~205 ($T=1.2\\pm0.6$ Gyr, $\\zh=-0.6\\pm0.2$), Hubble~VI in NGC~205 ($T=4\\pm2$ Gyr, $\\zh =-0.8\\pm0.2$), and FJJVII in NGC~185 ($T=7\\pm3$ Gyr, $\\zh =-0.8\\pm0.2$). The majority of our GC sample has solar $\\afe$ enhancement in contrast to the halo population of GCs in M31 and the Milky Way. The HB morphologies for our sample GCs follow the same behavior with metallicity as younger halo Galactic globular clusters. We show that it is unlikely that they bias our spectroscopic age estimates based on Balmer absorption line indices. Spectroscopic ages and metallicities of the central regions in NGC~205 and NGC~185 coincide with those obtained from color-magnitude diagrams. The central field stellar populations in these galaxies have approximately the same age as their most central GCs (Hubble~V in NGC~205 and FJJIII in NGC~185), but are more metal-rich than the central globular clusters. ",
        "introduction": "Understanding the role of dwarf galaxies and their globular cluster systems as building blocks of massive early-type and spiral galaxies is one of the great challenges of modern astrophysics. Studying the chemical compositions and color-magnitude diagrams (CMDs) of globular clusters (GCs) in the nearby low-mass galaxies is critical to compare properties of these long-living objects situated in galaxies of different type and mass. Only the Local group (LG) dwarf galaxies are close enough so that their globular clusters can be resolved into single red giant and horizontal branch stars by today's observatories.  Three close companions of the Andromeda galaxy, NGC~147, NGC~185 and NGC~205, comprise the brightest end of the LG dwarf early-type galaxy luminosity function. The proximity to M31 in combination with the unusual properties of these three galaxies have attracted great attention of astronomers since the beginning of the 20th century (we refer the reader to a monograph by Van den Bergh 2000 and a review by Mateo 1998 for details). NGC~205 and NGC~185 are very unusual galaxies and do not resemble any galaxy within $\\sim\\!10$ Mpc. Both have a regular spheroidal shape and are composed mainly of old stars. At the same time, both galaxies have a considerable amount of gas, dust, and a significant intermediate-age stellar component.  These unique properties are most likely tightly related their unusual evolutionary histories (e.g. Yong \\& Lo, 1997; Davidge, 1992; Lee et al., 1993; Butler \\& Martinez-Delgado, 2005; Davidge, 2005). Compared to its brighter cousins, NGC~147 is gas and dust free. Asymptotic giant branch (AGB) stars contribute only 2-3\\% of the total light of this galaxy (Davidge, 1994). By means of its stellar content, NGC~147 resembles a typical dwarf spheroidal galaxy in the Local Group. However, it is much more luminous. The distance to NGC~205, NGC~185 and NGC~147, their stellar populations and star formation histories are being studied actively (e.g.~McConnachie 2005, Butler \\& Martinez-Delgado 2005, Dolphin 2005). Table~1 summarises their distances, brightnesses, and reddenings used in this paper for the studied galaxies. We use a standard reddening law with $A_V = 3.315\\cdot E(B-V)$. We note that we find indications for a slightly smaller distance modulus, $(m-M)_0=24.38 \\pm 0.1$ mag, for NGC~205 compared to the values derived by McConnachie (2005), $(m-M)_0= 24.58 \\pm 0.07$ mag, and Dolphin (2005), $(m-M)_0=24.45 \\pm 0.14$ mag, based on the results of our stellar photometry (see Sect.~3). \\begin{table} \\centering \\scriptsize \\caption{Absolute magnitudes, distance moduli and reddenings from Schlegel et al. (1998) used for the galaxies of our study (indices refer to 0: this work; 1: McConnachie et al. (2004).} \\begin{tabular}{@{}lcccc@{}} \\hline NGC   & $M_{V,0}$ & $(M-m)_0$ &    & $A_V$        \\\\ &           &           &         \\\\ \\hline 147   & $-15.1^1$ & 23.95$^1$ &  0.580  \\\\ 185   & $-15.6^1$ & 24.15$^1$ &  0.593  \\\\ 205   & $-16.2^0$ & 24.38$^0$ &  0.206  \\\\ \\hline \\end{tabular} \\end{table}  \\begin{table} \\centering \\scriptsize \\caption{Properties of globular clusters in NGC~147, 185 and 205 summarised from the literature. Columns contain the following data: (1) object name, (2) integrated V magnitude, (3) integrated $U-B$ colour, (4) integrated $B-V$ colour, (5) metallicity (indices refer to 1: Hodge (1974), 2: Hodge (1976), 3: Hodge (1973) 4: Da Costa and Mould (1988), 5: Geisler et al. (1999), 6: Butler and Martinez-Delgado (2005), 7: Sharov and Lyuty (1983), 0: our HST/WFPC2 integrated $VI$ photometry corrected for Galactic extinction using Schlegel et al. (1998) maps. A star like symbol indicates, that $(V-I)_0$ color is given instead of $(B-V)$.} \\begin{tabular}{@{}lcllc@{}} \\hline Object     & V         & $U-B$    &$B-V$   & \\feh\\      \\\\ \\hline {\\bf NGC~147}&         &          &        &                  \\\\ HodgeI     & $17.7^2$  & $0.6^2$  &$0.8^2$ & $-1.9^4$ \\\\ HodgeII    & $16.5^1$  & $0.4:^7$ &$0.6^7$ &  .... \\\\ HodgeIII   & $17.0^2$  & $0.1^2$  &$0.6^2$ & $-2.5^4$ \\\\ &           &          &        &      \\\\ {\\bf NGC~185}&         &          &        &      \\\\ FJJI       & $18.4^1$  & $0.2^1$  &$0.8^1$ & $-1.4^4$ \\\\ FJJII      & $19.7^1$  & ....     &....    & $-1.2^4$ \\\\ FJJIII     & $16.8^1$  & $0.1^1$  &$0.9^1$ & $-1.7^4$, $-1.6^5$ \\\\ FJJIV      & $19.0^4$  & $0.05^1$ &$0.7^1$ & $-2.5^4$, $-1.9^5$ \\\\ FJJV       & $16.7^1$  & ....     &$0.7^1$ & $-1.8^4$, $-1.5^5$ \\\\ FJJVII     & $18.9^{0,*}$ & ....  &$1.0^{0,*}$ &  .... \\\\ &           &          &        &      \\\\ {\\bf NGC~205}&         &          &        &      \\\\ HubbleI    & $16.9^3$  &  ....    &$0.8^3$ & $-1.5^4$ \\\\ HubbleII   & $16.7^3$  & $0.2^3$  &$0.9^3$ & $-1.5^4$ \\\\ HubbleV    & $16.7^3$  & $0.2:^3$ &$0.6^3$ &   .... \\\\ &           &          &        &    .... \\\\ HubbleVI   & $17.9^3$  & $0.8^3$  &$0.7^3$ & $-1.3^4$ \\\\ HubbleVII  & $18.0^3$  & $0.1^3$  &$0.9^3$ & $-1.4^4$ \\\\ HubbleVIII & $16.6^3$  & $0.1^3$  &$0.7^3$ & $-1.9^4$  \\\\ \\hline \\end{tabular} \\end{table} \\begin{table} \\centering \\scriptsize \\caption{Journal of spectroscopic observations} \\begin{tabular}{@{}lclr@{}} \\hline Object      & Date       & Exposure & Seeing         \\\\ \\hline {\\bf NGC~147}      &            &          &                 \\\\ Hodge I, II, III   & 23.09.2003 & 3x900s   & 1.2       \\\\ \\noalign{\\smallskip} {\\bf NGC~185}      &            &          &                 \\\\ FJJIII, IV, V, VI, VII& 22.09.2003 & 4x900s   & 1.2       \\\\ FJJI, FJJII & 26.09.2003 & 3x900s   & 1.2       \\\\ FJJVIII     & 26.09.2003 & 3x600s   & 1.2       \\\\ \\noalign{\\smallskip} {\\bf NGC~205}     &            &          &                 \\\\ Hubble II, V, VI, VII & 23.09.2003 & 4x900s   & 1.2       \\\\ Hubble I    & 25.09.2003 &   900s   & 1.2       \\\\ \\noalign{\\smallskip} \\noalign{\\smallskip} {\\bf Stars}      &            &          &                 \\\\ HD10307  & 25.09.2003 & 2x30s      & 1.2     \\\\ HD26784  & 26.09.2003 & 2x20s      & 1.2     \\\\ HD26757  & 26.09.2003 & 2x20s      & 1.2     \\\\ BD+26377 & 26.09.2003 & 2x30s   &  1.2      \\\\ \\hline \\end{tabular} \\end{table}  Each of our sample galaxies hosts a sizable globular cluster system (GCS). The discovery of globular clusters in NGC~147, NGC~185 and NGC~205 includes the work of Hubble (1932), Baade (1944), Hodge (1974) and Ford, Jacoby, \\& Jenner (1977, hereafter FJJ77) (see Da Costa and Mould (1988) and van den Bergh (2000) for details). All searches of GCs were carried out on photographic plates. The clusters were classified as being associated with the galaxies by their galactocentric radius. However, recent studies of GCSs in dwarf galaxies reveal globular clusters situated far from optical bodies of their host galaxies (see e.g. Hwang et al. 2005). It is likely that detailed wide-field investigations may discover additional GC candidates around NGC~147, NGC~185 and NGC~205 using new observational facilities and improved classification schemes which would allow the measurement of real physical sizes of the GCSs in these dwarf galaxies. Four extended globular clusters were discovered in the outer halo of M31 by Huxor et al., 2005) using INT-WFS images. Recent HST based surveys reveal new globular cluster candidates in M31 (Barmby \\& Huchra 2001, see also the summary of optical and IR photometry in Galleti et al. 2004). For instance, there are $\\sim\\!50$ objects in the catalog of Galleti et al. (2004) within a radius of 20\\farcm\\ around the center of NGC~205. Galleti et al. (2006) emphasize that our knowledge about the M31 GCS is far from complete even in terms of simple membership and undertake a large spectroscopic survey of bona-fide globular clusters and cluster candidates in M31. Obtaining radial velocities for GCs in the M31 halo and in the dwarf galaxies, detailed CMD and abundance studies will certainly help to solve the problem of the origin of M31 and of the satellite galaxies.  The properties of the globular clusters in NGC~147, NGC~185 and NGC~205 from the literature are summarised in Table~2. Da Costa and Mould (1988) first obtained spectra for eight Hubble clusters, the nucleus of NGC~205, FJJ77 clusters I--V in NGC~185 and for the two clusters Hodge~I and III in NGC~147. A comparison of line strengths relative to similar spectra of Galactic globular clusters indicated that with a single exception, of the central GC Hubble~V in NGC~205, all the clusters in NGC~147, NGC~185 and NGC~205 are old and metal-poor. The age estimates for Hubble~V fall in a wide range from $\\sim\\!300$ Myr (based on $BVR$ photometry, Lee, 1996) to intermediate ages (based on optical and far-ultraviolet imaging with HST/WFPC2, Jones et al., 1996). Star cluster formation seems to be a more extended process in NGC~185. A star forming region of $H \\alpha+[NII]$ emission with a diameter of ~50 pc near the center of NGC~185 may be a probable place of a massive young star cluster birth according to a modern evolutionary scheme (Gallagher \\& Grebel, 2001 and references therein). Young \\& Lo (1997) estimated the mass of the brightest central HI clump in NGC~185 $M_{HI}=0.7 \\cdot 10^4 M_{\\sun}$ and the corresponding virial mass a factor of 100 higher. The clump is associated with a massive giant molecular cloud.  \\begin{figure*} \\vspace{-2cm} \\includegraphics[width=0.9\\textwidth]{Fig1.eps} \\vspace{-2cm} \\caption{The upper left panel shows an image of NGC~205 obtained with the SCORPIO multi-slit unit at the 6-m RAS telescope. Settings of slits on the globular clusters and diffuse galactic light regions are indicated. The lower left panel shows a frame with raw non-reduced spectra of objects in NGC~205. Two right panels show Digital Sky Survey 10\\arcmin\\ by 10\\arcmin\\ images of NGC~185 and NGC~147 with globular clusters.} \\label{dss} \\end{figure*}  In this paper, we derive ages, metallicities, and \\afe\\ ratios for globular clusters in NGC~147, NGC~185 and NGC~205 based on measurements of absorption line indices as defined by the Lick standard system (Burstein et al. 1984; Worthey et al. 1994; Worthey \\& Ottaviani 1997; Trager et al. 1998). The main advantage of this system is its ability to disentangle age from metallicity using the information on the strength of prominent diagnostic absorption line features measured in the intermediate-resolution integrated-light spectra.  The paper is organised as follows. In Section~2 we describe our spectroscopic data, the data reduction, as well as our methods and results of measuring Lick line indices in the spectra of globular clusters and the diffuse host galaxy light. In Section~3 we present results from our CMD photometry analysis of globular clusters, where we compare the loci of red giant branch (RGB) and horizontal branch (HB) stars with the Victoria-Regina stellar models of VandenBerg et al. (2006). In Section~4 we discuss the derived evolutionary parameters for globular clusters and compare them with the data available in the literature. ",
        "conclusions": "Although the modest S/N of some of our GC spectra and the relatively moderate photometric depth of our CMDs do not allow us to derive very accurate ages, we conclude that all our sample GCs appear to be old ($T>8$ Gyr) and metal-poor ($\\zh \\la -1.1$), except for the GCs Hubble~V in NGC~205 ($T=1.2\\pm0.6$ Gyr, $\\zh=-0.6\\pm0.2$), Hubble~VI in NGC~205 ($T=4\\pm2$ Gyr, $\\zh =-0.8\\pm0.2$), and FJJVII in NGC~185 ($T=7\\pm3$ Gyr, $\\zh =-0.8\\pm0.2$). We find two intermediate-age GCs (HubbleVI and FJJVII) which are located at projected distances $D_{proj} \\sim0.3$ and 1.0 kpc from the centers of their host galaxies. Deep images taken with HST/ACS are necessary to understand whether these relatively faint star clusters are genuine globular clusters (i.e. simple stellar populations) born during $\\sim\\!4-6$ Gyr old star-formation events in their host galaxies. They could consist also of multiple stellar populations similar to $\\omega$Cen (see e.g. Sollima et al. 2005).  The HB morphologies for our sample GCs follow the same behavior with metallicity as younger halo Galactic globular clusters. It was found for Galactic GC, that age may not be the dominant second parameter determining the shape of globular clusters' HBs (see e.g. Stetson et al., 1996). GCs having blue HBs are in general older than 10 Gyr (Lee et al. 1994). So, we suggest here that the HB morphologies for our sample GCs likely do not bias our spectroscopic age estimates based on Balmer absorption line indices.  We find that most of the GCs in the studied galaxies are weakly or not $\\alpha$-enhanced, in contrast to the population of GCs in nearby early-type galaxies (see Puzia et al. 2006). The chemical composition of globular clusters may turn out to be a powerful tool to discriminate between clusters which formed {\\it in situ} in massive galaxies and those that were formed in smaller sub-units and later accreted in more massive halos.  Spectroscopic ages and metallicities of the central regions in NGC~205 and NGC~185 coincide with those obtained from color-magnitude diagrams. The central field stellar populations in these galaxies have approximately the same age as their most central GCs Hubble~V in NGC~205 and FJJIII in NGC~185, respectively, but are more metal-rich than the central globular clusters."
    },
    "0608/astro-ph0608361_arXiv.txt": {
        "abstract": "A description of the light curves of microlensed Type Ia supernovae (SNe~Ia) as extended and expanding sources is presented.  We give examples of what microlensing by stellar-mass Chang-Refsdal lenses can do to a small percentage of supernova light curves.  We find that in addition to overall brightening, significant changes in light--curve shapes can also occur.  Peaks can be distorted, plateaus can appear, and even secondary peaks can be formed.  The effects of both the relative motion of the lens and the supernova and the expansion of the supernova are given and compared.  The effects of relative motion are more pronounced when a distant supernova ($z_s \\sim 1$) impacts well within the Einstein ring of a nearby microlens ($z_d \\sim 0.05$) and are less important for more distant deflectors.  We also find that the increase in shear that comes with increased deflector distance tends to reduce the time variability of microlensing.  We briefly discuss the probability of observing these effects. ",
        "introduction": "Out of a number of distance indicators, Type~Ia supernovae (SNe~Ia) have emerged as the most promising standard candles. Due to their significant intrinsic brightness and relative ubiquity they can be observed in the local and distant universe. Several teams including the High-z Supernova Search \\citep{Schmidt98} and the Supernova Cosmology Project \\citep{Perlmutter99} have been searching for SNe~Ia at higher redshifts since the early 1990's. Light emitted from these `standard candles' is subject to lensing by intervening objects while traversing the large distances involved \\citep{KVB95}; the further the light source, the higher its chance of being significantly lensed.  In fact, for cosmologically distant sources, the probability is high that a distant point source will be `imaged' \\citep{Press73,BK76,Wyithe02}, particularly by stellar objects (microlensing).  While the systematic errors introduced by K-correction, selection effects, and possible evolution can be removed, lensing might ultimately limit the accuracy of luminosity distance measurements \\citep{Perlmutter03}. Only a large sample of SNe~Ia at each redshift can be used to characterize the lensing distribution and to correct for the effect of weak lensing. Proposed searches, e.g., \\cite{Tyson05, Corasaniti06} are expected to yield hundreds of thousands of SNe~Ia out to $z \\sim 1$, some of which should exhibit lensing effects such as described here.  Properties of microlensed supernovae have been studied previously. \\citet{Schneider87} presented the time-dependent amplification of supernovae caused by their expansion and showed that the related polarization is not likely to exceed 1\\% and hence not possible to detect among cosmologically distant supernovae at this time. \\citet{Linder88} studied amplification of supernovae and developed approximate formulae for the amplification probability distribution. \\citet{Rauch91} studied microlensing of SNe Ia by compact objects and calculated the resulting amplification probability distributions using Monte Carlo simulations.  Assuming point-deflectors with shear \\citet{Kolatt98} estimated the microlensing rate by MACHOs in nearby clusters. The supernovae were treated as point sources and consequently the microlensed light curves were expected to contain amplification peaks similar to caustic crossing events seen in the OGLE data \\citep{Udalski03}.  In this paper, we demonstrate in detail how microlensing by a single stellar deflector in an external shear field can affect light curves of cosmologically distant SNe Ia. We use the lens model of \\citet{Chang79} to calculate amplifications \\citep{Schneider92}.  We ignore any lensing amplification caused by the SN's hosting galaxy and concentrate on the time-dependent effects caused by a single moving stellar deflector.  We model the SNe Ia as expanding light sources with limb-darkening.  In the next section, we present our model for SNe~Ia as microlensed sources.  Section 3 is devoted to a brief discussion of microlensing of finite-size sources, and in $\\S$ 4 we present the results of our calculations.  Throughout, we assume a flat Friedmann-Lama\\^{\\i}tre-Robertson-Walker cosmological model with $\\Omega _{m}=$ 0.3, $\\Omega _{\\Lambda }=$ 0.7, and h$_{100}=$ 0.67 to calculate the distances to the source and deflector. ",
        "conclusions": "We have shown that microlensing can significantly affect light curves of some cosmologically distant SNe~Ia. We restricted our calculation to sources at $z_{s}=1$ in the currently accepted $\\Omega _{m}=0.3$, $\\Omega _{\\Lambda }=0.7$ flat cosmological model. We found that microlensing can not only increase the magnitude of the light curve but also can cause a change in its shape. Relative transverse motion of the SN and lens, when added to the expanding photosphere, can result in features such as an an enhanced peak brightness with a distorted shape as in Figure 3, a post peak plateau as in Figure 4, or even the presence of a wide second peak as in Figure 5.  In the absence of relative lens-source motion \\citet{Schneider87} found distortions to the light curve's peak caused by the supernova's photosphere expanding into a deflector's critical point as well as into its caustics. However, from row 3 of Figures 3-5 it is easy to see how ignoring the relative motion at small deflector redshifts can result in underestimating the variety and intensity of lensing effects on the SNe light curves.  We find that stationary deflectors do not produce features such as bumps or plateaus as often as the corresponding moving deflectors. This is due to the fact that the amplification curves for stationary lenses tend to be flattened and show less change in amplification as the photosphere expands through and beyond the deflector. The effect of amplification by stationary deflectors appears to produce more of an overall upward shift in the supernova's light curve. A look at row 3 of Figures 3 through 5 shows that the magnitude difference of the moving versus stationary cases is measurable for the deflectors at $z_{d}$=0.05 and 0.10 but not as far as 0.35.  The magnitude differences fall below 0.1 at $z \\sim 0.25$.  Also, it should be pointed out that a supernova lensed by a stationary deflector is brighter than one lensed by a moving deflector beyond the time when the projected distance of source and deflector is greater than the initial value $l_{o}$; see, for instance, the right column of row 3 of Figure 3 where the two $z_d=0.05$ curves go negative at $T=169$ days.    In $\\S$ 3 we pointed out, that for microlensing by compact masses distributed through the cosmos, the optical depth is 0.01 $\\left(z_{s}\\sim 1\\right)$ and might reach 0.1 $\\left(z_{s}\\sim 2\\right)$, implying that the overall chance of a distant SN~Ia being microlensed is not negligible. Any multi-band supernova survey aimed at finding supernovae at redshifts around $z=1$ (and above), could discover and identify \\emph{one} microlensed SN Ia event out of roughly \\emph{a hundred} events. However, the low impact parameters required to produce the special features depicted in $\\S$ 4 demand observation of $\\sim 10^{4}$ supernovae at $z_{s}\\geqslant 1$.  To see unusual features such as double peaks, the lensed supernova must be followed for an extended period of $\\sim$ 2 years.  In the Appendix we have made an optical-depth type estimate to include double peak events. For microlensing by stars in the bulge of a galaxy at $z_d=0.05$ we find a max probability of $\\sim 1.7\\times 10^{-3}$. This is an ideal deflector distance for observing double peaks due to transverse motion.  As expected this number is only slightly smaller than the 27\\% optical depth estimate made in $\\S 3.3$ for bulge lensing when corrected for an impact parameter of $u=0.1$. These estimates can double when the observational bias in favor of amplified events are taken advantage of (see, for instance, \\citet{Gunnar03}). It is interesting to note that according to OGLE III \\citep{Udalski03} a histogram of $u$ values for Bulge lensing peaks at $u\\sim 0.1$, probably due to amplification biasing.  Notice that we have not taken into account macrolensing  convergence effects caused by the deflector's host galaxy. The effects of convergence are relatively easy to include, \\eg in our case the amplifications for the Schwarzschild lens given in the second rows of Figure 3-5 would increase by $\\sim$30\\% if the lens galaxy had an unexpectedly high transparent surface mass density of $\\Sigma=$ 1000 M$_{\\sun}$pc$^{-2}$ at the image."
    },
    "0608/astro-ph0608682_arXiv.txt": {
        "abstract": "We solve the one-dimensional neutrino-heated non-relativistic magnetohydrodynamic (MHD) wind problem for conditions that range from slowly rotating (spin period $P\\gtrsim10$\\,ms) protoneutron stars (PNSs) with surface field strengths typical of radio pulsars ($B\\lesssim10^{13}$\\,G), to \"proto-magnetars\" with $B\\approx10^{14}-10^{15}$\\,G in their hypothesized rapidly rotating initial states ($P\\approx1$\\,ms).  We use the relativistic axisymmetric simulations of Bucciantini et al.~(2006) to map our split-monopole results onto a more physical dipole geometry and to estimate the spindown of PNSs when their winds are relativistic.  We then quantify the effects of rotation and magnetic fields on the mass loss, energy loss, and thermodynamic structure of PNS winds. The latter is particularly important in assessing PNS winds as the astrophysical site for the $r$-process.  We describe the evolution of PNS winds through the Kelvin-Helmholtz cooling epoch, emphasizing the transition between (1) thermal neutrino-driven, (2) non-relativistic magnetically-dominated, and (3) relativistic magnetically-dominated outflows.  In the last of these stages, the spindown is enhanced relative to the canonical force-free rate because of additional open magnetic flux caused by neutrino-driven mass loss. We find that proto-magnetars with $P\\approx 1$ ms and $B \\gtrsim 10^{15}$ G drive relativistic winds with luminosities, energies, and Lorentz factors (magnetization $\\sigma \\sim 0.1-1000$) consistent with those required to produce long duration gamma-ray bursts and hyper-energetic supernovae (SNe). A significant fraction of the rotational energy may be extracted in only a few seconds, sufficiently rapidly to alter the asymptotic energy of the SN remnant, its morphology, and, potentially, its nucleosynthetic yield.  We find that winds from PNSs with somewhat more modest rotation periods ($\\approx 2 - 10$ ms) and with magnetar-strength fields produce conditions significantly more favorable for the $r$-process than winds from slowly rotating, non-magnetized PNSs.  Lastly, we argue that energy and momentum deposition by convectively-excited waves may be important in PNS winds.  We show that this further increases the likelihood of successful $r$-process, relatively independent of the PNS rotation rate and magnetic field strength. ",
        "introduction": "\\label{section:intro}  On timescales $\\lesssim$ 1 s following the core collapse of a massive star, neutrino emission from the resulting hot, deleptonizing protoneutron star (PNS) may play an essential role in launching the supernova (SN) shock (Herant et al.~1994; Burrows, Hayes, $\\&$ Fryxell 1995; Janka $\\&$ M{\\\"u}ller 1995) or in generating large-scale anisotropies through hydrodynamical instabilities (Blondin et al.~2003; Scheck et al.~2006; Burrows et al.~2006a,b).  Independent of how the explosion is initiated at early times, a small fraction of the PNS's cooling neutrino emission continues heating the surface layers of the PNS, driving a persistent thermal wind into the cavity evacuated by the rapidly-expanding SN shock (Duncan, Shapiro, $\\&$ Wasserman 1986; Woosley et al.~1994); this post-explosion neutrino-driven mass loss persists for the duration of the PNS's Kelvin-Helmholtz cooling epoch, which lasts a time $\\tau_{{\\rm KH}} \\sim 10-100$ s (Burrows $\\&$ Lattimer 1986; Pons et al.~1999).  Once the SN shock has been launched and the PNS cooling epoch has begun, neutrino-driven winds from non-rotating non-magnetic PNSs are unlikely to be energetically important on the scale of the kinetic energy of the accompanying SN ($E_{\\rm SN}\\approx10^{51}$\\,ergs).  However, in the presence of a sufficiently strong global magnetic field the dynamics of a PNS's neutrino-heated outflow are significantly altered (e.g., Thompson 2003a,b).  This point is germane because $\\sim 10\\%$ (Kouveliotou et al.~1994; van Paradijs et al.~1995; Lyne et al.~1998) of young Galactic neutron stars possess significantly stronger surface magnetic fields ($\\sim 10^{14}-10^{15}$ G) than those usually inferred from pulsar spin-down estimates (``magnetars''; for a recent review see Woods \\& Thompson 2004).  While the precise origin of these large field strengths is uncertain, it has been argued that their amplification occurs via a dynamo during $\\tau_{{\\rm KH}}$ (Duncan $\\&$ Thompson 1992; Thompson \\& Duncan 1993).  The efficiency of the dynamo is determined in part by the core's initial rotation rate $\\Omega_{0} = 2\\pi/P_{0}$, and analytic arguments suggest that the formation of global magnetar-strength fields might require $P_{0} \\sim$ 1 ms rotation at birth.  Such rapid rotation would also alter the dynamics of the PNS wind and provide a reservoir of rotational energy significant on the scale of the accompanying SN explosion: \\be E_{{\\rm rot}} \\simeq 2\\times 10^{52}\\,M_{1.4}\\,R_{10}^{2}\\,P_{{\\rm ms}}^{-2}\\,{\\rm ergs}, \\label{erot} \\ee where $M_{1.4}$ is the PNS mass in units of 1.4 M$_{\\sun}$, $R_{10}$ is the radius of the PNS in units of 10 km, and $P_{{\\rm ms}}$ is the initial PNS rotation period in ms.  Previous authors have suggested that if magnetars are indeed born rapidly rotating, their rotational energy could be efficiently extracted through a magnetized wind (Usov 1992; Thompson 1994; Wheeler et al.~2000; Thompson, Chang, \\& Quataert 2004, henceforth TCQ).  Indeed, a proto-magnetar wind's energetics, timescale, and potential for highly relativistic outflow resemble those of the central engine required to power long-duration gamma-ray bursts (LGRBs). TCQ argue that an accurate description of magnetized PNS spin-down must include the effects of neutrino-driven mass loss.  The significant mass loss accompanying the Kelvin-Helmholtz epoch may open the otherwise closed magnetosphere into a ``split monopole''-like structure, enhancing the early spin-down rate of the PNS.  In addition, the mass loading of a PNS wind, and hence its potential asymptotic Lorentz factor $\\Gamma$, is largely controlled by the neutrino luminosity during the Kelvin-Helmholtz epoch.  This ``baryon-loading'' issue is particularly important in the LGRB context, where $\\Gamma \\sim 10-1000$ are typically inferred.  Because magnetar births are relatively frequent, most cannot produce classical LGRBs, which only occur in $\\sim 0.1-1\\%$ of massive stellar deaths (e.g., Paczynski 2001, Podsiadlowski et al. 2004, Piran 2005); however, a more common observational signature of magnetar birth may include less energetic or mildly relativistic events, which could be observable as X-ray transients or unusual SNe.  Under some circumstances, the asymmetric energy injection from proto-magnetar winds could produce global anisotropies in the SN remnant, as has been detected through polarization measurements of some SNe (e.g., Wang et al.~2001, 2003).  In addition, if significant rotational energy can be extracted sufficiently rapidly following the launch of the SN shock, the nucleosynthetic yield of the SN could be altered (TCQ), which could explain nickel-rich, hyper-energetic SNe such as SN 1998bw (Galama et al.~1998) or SN 2003dh (Hjorth et al.~2003; Stanek et al.~2003).  Lastly, LGRBs occurring without associated SNe (e.g., GRBs 060505, 060614; Fynbo et al.~2006; Gal-Yam et al.~2006; Della Valle et al.~2006) may be accommodated if a proto-magnetar results from the accretion-induced collapse (AIC) of a white dwarf (see \\S\\ref{section:AIC}).  Quite apart from the possible impact of PNS winds on the surrounding SN shock, PNS winds themselves have often been considered a promising site for the production of $r$-process nuclides.  However, the conditions necessary for a successful third peak ($A \\approx 195$) $r$-process have not been realized in detailed studies of non-rotating, non-magnetized PNS winds (e.g., Qian \\& Woosley 1996; Cardall \\& Fuller 1997; Otsuki et al.~2000; Thompson et al.~2001).  Given the, as yet, unidentified site of Galactic $r$-process enrichment and the relatively large birthrate of magnetars, it is essential to consider what effects a strong magnetic field and rapid rotation might have on nucleosynthesis in PNS winds.  Conversely, if the nucleosynthetic yield from rotating, magnetized PNS winds could be well-determined theoretically, the birth period and magnetic field distribution of neutron stars could perhaps be constrained from observations of heavy elemental abundances.  \\subsection{Stages of PNS Evolution} \\label{section:stages}  A sufficiently-magnetized PNS wind goes through at least three distinct stages of evolution following the launch of the SN shock ($t=0$):  (1) During the earliest phase the PNS surface temperature is so high ($\\gtrsim 5$ MeV) and the radius of the PNS is so large ($R_{\\nu} \\sim 20-50$ km) that, even for a magnetar-strength field, the post-explosion outflow is likely to be purely thermally-driven.  (2)  By $t \\sim$ 1 s the SN shock has propagated to well outside the sonic point of the PNS wind.  The PNS will contract and cool to a point at which, if the surface field is sufficiently strong, the outflow becomes magnetically dominated.  The strong magnetic field enhances angular momentum and rotational energy loss by forcing outgoing fluid elements to effectively corotate with the PNS surface out to the Alfv\\'{e}n radius ($R_{A}$) at several stellar radii, in analogy with classic work on non-relativistic stellar winds and the solar wind (e.g., Schatzman 1962, Weber \\& Davis 1967, and Mestel 1968).  For sufficiently rapidly rotating PNSs, the spin-down timescale ($\\tau_{{\\rm J}} \\equiv \\Omega/\\dot{\\Omega}$) can be comparable to $\\tau_{{\\rm KH}}$, implying that much of the rotational energy of the PNS can be extracted in a non-relativistic, but magnetically-dominated wind.  The neutrino luminosity at these relatively early times is still large (e.g., $L_{\\nu} \\sim 10^{52}$ ergs s$^{-1}$; see Pons et al.~1999, Fig.~14), and, for sufficiently rapid rotation, mass loss is significantly enhanced by centrifugal flinging (TCQ).  The outflow during this phase will be collimated about the PNS rotation axis by magnetic stresses (e.g., Bucciantini et al.~2006; hereafter B06).  (3) As the PNS continues to cool, the luminosity and mass-loading decrease to the point at which corotation is sustained out to nearly the light cylinder ($R_{{\\rm L}} \\equiv 2\\pi c/P \\simeq 48\\,P_{{\\rm ms}}^{-1}$ km).  The spin-down rate then becomes approximately independent of the mass loss rate $\\dot{M}$ and the flow becomes relativistic, obtaining high magnetization $\\sigma \\equiv \\Phi_{B}^{2}\\Omega^{2}/\\dot{M}c^{3}$, where $\\Phi_{B}$ is the total open magnetic flux per $4\\pi$ steradian (Michel 1969).  Although an ultra-relativistic, pulsar-like wind is inevitable soon after $\\tau_{{\\rm KH}}$ because $\\dot{M}$ abates as the neutrino luminosity vanishes, relativistic outflow can begin prior to the end of the Kelvin-Helmholtz phase.  Because this outflow is accompanied by significant mass loss, it is generally only mildly relativistic.  Although the energy extracted via relativistic winds is primarily concentrated at low latitudes (B06), the confining pressure of the overlying, exploding, stellar progenitor (e.g., Wheeler et.~al 2000; Uzdensky $\\&$ MacFadyen 2006) or the walls of the collimated cavity carved by the preceding non-relativistic phase may channel the relativistic outflow into a bipolar, jet-like structure.  The primary focus of this paper is to delineate the magnetic field strengths and rotation rates required to significantly alter the characteristics of early PNS evolution by solving the one-dimensional (1D) non-relativistic MHD, neutrino-heated wind problem for conditions that range from normal pulsars to proto-magnetars.  In particular, we quantitatively explore the transition that occurs between stages (1) and (2) above; with these results we analyze some of the immediate consequences of neutron star birth.  We defer a detailed study of the transition between stages (2) and (3) to future work, but we do explicitly address the parameter space of the $\\sigma=1$ boundary.  Relativistic, mass-loaded MHD winds have been studied recently by B06 in two dimensions for both a monopolar and aligned-dipolar field structure, assuming an adiabatic equation of state.  While the work of B06 is critical to understanding the multi-dimensional character of PNS winds (for instance, the degree of collimation and the fraction of open magnetic flux) it does not address the neutrino microphysics necessary for a direct application to PNS environments; this work and that of B06 are thus complementary in studying PNS spin-down in the presence of rapid rotation and a large magnetic field. \\subsection{This Paper} In $\\S$\\ref{section:model} we enumerate the equations of MHD ($\\S$\\ref{section:equations}), discuss the relevant microphysics ($\\S$\\ref{section:microphysics}), and elaborate on our numerical methods ($\\S$\\ref{section:numerical}).  Section \\ref{section:results} presents the results of our calculations and examines the regimes of magnetized PNS wind evolution.  Section \\ref{section:discussion} discusses the implications of this work, examining the time-evolution of a cooling, magnetized PNS ($\\S$\\ref{section:evolution}), weighing the implications for LGRBs and hyper-energetic SNe ($\\S$\\ref{section:LGRB}), and considering the viability for third-peak $r$-process nucleosynthesis in magnetized, rotating PNS winds ($\\S$\\ref{section:rprocess}).  In $\\S$\\ref{section:waveheating} we consider the effects that wave heating have on the $r$-process in PNS winds and in $\\S\\ref{section:AIC}$ we briefly discuss other contexts in which our calculations may be applicable, including the accretion-induced collapse of white dwarfs, ``collapsars,'' and merging neutron star binaries.  Finally, section \\ref{section:conclusions} summarizes the conclusions of our work. ",
        "conclusions": "\\label{section:conclusions}  We have solved the one-dimensional non-relativistic neutrino-heated MHD wind problem in order to study the effects that magnetic fields and rotation have on PNS wind evolution following the launch of the SN shock but prior to the end of the Kelvin-Helmholtz cooling epoch. Figure $\\ref{plot:regimes}$ summarizes the physical regimes of PNS winds in the presence of rotation and magnetic fields.  We map our monopole solutions onto the axisymmetric, relativistic dipole calculations of B06, thus taking into account the effects that neutrino-driven mass loss have on the fraction of open magnetic flux and the PNS spin-down rate.  Our primary conclusions are as follows:  \\begin{itemize} \\item{We identify three types of PNSs based on the dominant character of their energy loss during the Kelvin-Helmholtz epoch (see Figure \\ref{plot:etot}): \\begin{itemize} \\item[(1)]{ For slow rotation and low magnetic field strengths $(B_{14}^{{\\rm dip}} \\lesssim 10^{-2}P_{{\\rm ms}}^{2}$, where $B^{{\\rm dip}} = 10^{14} B_{14}^{{\\rm dip}}$ G is the surface dipole field strength of the PNS and $P = 1 P_{{\\rm ms}}$ ms is its birth period), a neutrino-heated, thermally-driven wind dominates energy and mass loss from the PNS; most radio pulsars were probably this type at birth.} \\item[(2)]{ For larger rotation rates and field strengths $(10^{-2}P_{{\\rm ms}}^{2} \\lesssim B_{14}^{{\\rm dip}} \\lesssim 2P_{{\\rm ms}}^{2}\\exp[2P_{{\\rm ms}}^{-2}])$, a non-relativistic, magnetically-driven wind dominates during the Kelvin-Helmholtz epoch.  Most observed Galactic magnetars ($\\sim 10\\%$ the birthrate of radio pulsars) were probably of this type at birth.  For $B \\sim 10^{15}$ G and $P \\lesssim$ 2 ms, greater than $10^{51}$ ergs can be lost to a non-relativistic, magnetically-driven outflow during $\\tau_{KH}$, and, over a broad range of initial spin period, the energy extracted is many times larger than from a slowly rotating PNS (see Fig.~\\ref{plot:etot}).  The outflow from this type of PNS is likely to be collimated by magnetic stresses and the asymmetric injection of energy may be sufficient to generate an anisotropy in the morphology of the SNe remnant.} \\item[(3)]{For rapid rotation and field strengths somewhat larger than those observed from Galactic magnetars $(B_{14}^{{\\rm dip}} \\gtrsim 2P_{{\\rm ms}}^{2}\\exp[2P_{{\\rm ms}}^{-2}])$, a relativistic, magnetically-driven wind dominates energy loss during the Kelvin-Helmholtz epoch.  Although the birthrate of PNSs of this type is probably small (if they are produced at all), they may be capable of producing hyper-energetic SNe and long-duration gamma-ray bursts.} \\end{itemize} }  \\item{For $P \\approx$ 3(1) ms and $B_{\\nu}^{{\\rm dip}} \\gtrsim 10^{15}$ G the total mass loss during the Kelvin-Helmholtz epoch is enhanced by a factor of $\\approx$ 2($10^{2}$) relative to a non-rotating PNS (Fig. \\ref{plot:mtot}).}  \\item{For initial PNS spin periods of $P \\approx $ 1 ms and magnetic field strengths of $B_{\\nu}^{{\\rm dip}} \\approx 3\\times 10^{15}-10^{16}$ G, we find that $\\gtrsim 10^{52}$ ergs of rotational energy is extracted on a timescale of $10-40$ seconds and that the magnetization $\\sigma$ of the outflow is $\\sim 0.1-1000$.  The energy, luminosity, timescale, and mass-loading of the late-time outflow are all consistent with those required to explain long duration gamma-ray bursts (assuming efficient dissipation of magnetic energy into kinetic energy at large radii; e.g., Drenkhahn \\& Spruit 2002).  In addition, outflows from such PNSs have the property that the PNS rotational energy is extracted with a roughly uniform distribution in log($\\sigma$) over the timescale $\\tau_{{\\rm KH}}$.  For these modest $\\sigma$ winds, especially at early times, energy loss from the PNS is enhanced relative to pure force-free spindown because of additional open magnetic flux (see the discussion in $\\S$\\ref{section:evolution}).  Thus, a significant portion of the PNS rotational energy can be extracted in just a few seconds following the launch of the SN shock, perhaps sufficiently rapidly to increase the nucleosynthetic yield of the SN (e.g., $^{56}$Ni).}  \\item{Winds from PNSs with $B_{\\nu}^{{\\rm dip}}\\gtrsim 10^{15}$ G and $P \\approx 2-10$ ms produce conditions almost an order-of-magnitude more favorable for third-peak $r$-process nucleosynthesis in the space of $(S^{a})^{3}/\\tau_{{\\rm dyn}}$ than do winds from more slowly-rotating, less-magnetized PNSs (see Fig. \\ref{plot:svsomega}).  For these rotation rates, the asymptotic entropy is similar to that of a non-rotating, non-magnetized wind, while the dynamical time is significantly reduced by magnetic acceleration (see Fig. \\ref{plot:tdynplot}).  The very different thermodynamic properties of rapidly rotating proto-magnetar winds (relative to non-rotating PNS winds) may contribute to the inferred diversity of $r$-process sites (e.g., Qian et al.~1998).}  \\item{Heating by outgoing hydrodynamic or MHD waves may be important in PNS winds.  We find that if $\\gtrsim 10^{-4}-10^{-3}$ of the PNS's neutrino luminosity emerges in wave power at early times in the Kelvin-Helmholtz epoch, then r-process is successful in PNS winds.  This conclusion is relatively independent of magnetic field strength and rotation period for $P \\gtrsim$ 2 ms.} \\end{itemize}"
    },
    "0608/astro-ph0608157_arXiv.txt": {
        "abstract": "We  use a  series  of  cosmological N-body  simulations  for a  flat $\\Lambda$CDM cosmology  to investigate the  structural properties of dark matter  haloes, at redshift zero,  in the mass  range $ 3\\times 10^{9}  \\hMsun  \\lta M_{\\rm  vir}  \\lta  3\\times 10^{13}  $$\\hMsun$. These properties include the  concentration parameter, $c$, the spin parameter,  $\\lambda$, and the  mean axis  ratio, $\\qbar$.   For the concentration-mass  relation   we  find  $c\\propto\\Mvir^{-0.11}$  in agreement  with the  model  proposed by  Bullock  \\etal (2001),  but inconsistent with the alternative model of Eke, Navarro \\& Steinmetz (2001).   The  normalization  of  the  concentration-mass  relation, however, is 15 percent lower than suggested by Bullock \\etal (2001). The results  for $\\lambda$  and $\\qbar$ are  in good  agreement with previous studies, when extrapolated  to the lower halo masses probed here, while $c$ and $\\lambda$ are anti-correlated, in that high-spin haloes  have, on average,  lower concentrations.   In an  attempt to remove unrelaxed  haloes from the  sample, we compute for  each halo the offset  parameter, $\\xoff$, defined as the  distance between the most bound particle  and the center of mass, in  units of the virial radius.   Removing  haloes with  large  $\\xoff$  increases the  mean concentration by  $\\sim 10$ percent, lowers the  mean spin parameter by  $\\sim 15$  percent, and  removes  the most  prolate haloes.   In addition, it  largely removes  the anti-correlation between  $c$ and $\\lambda$, though  not entirely.   We also investigate  the relation between halo  properties and their  large-scale environment density. For low  mass haloes we find  that more concentrated  haloes live in denser environments than their less concentrated counterparts of the same  mass, consistent  with recent  correlation  function analyses. Note, however, that the trend  is weak compared to the scatter.  For the halo  spin parameters we  find no environment  dependence, while there is a weak indication  that the most spherical haloes reside in slightly  denser environments.   Finally, using  a simple  model for disk  galaxy formation  we show  that haloes  that host  low surface brightness galaxies are expected to be hosted by a biased sub-set of haloes.   Not only  do these  haloes have  spin parameters  that are larger than  average, they  also have concentration  parameters that are $\\sim 15$  percent lower than the average at  a given halo mass. We discuss  the implications of  all these findings for  the claimed disagreement between halo  concentrations inferred from LSB rotation curves, and those expected for a $\\Lambda$CDM cosmology. ",
        "introduction": "\\label{sec:intro}  The theory of  cold dark matter (CDM) provides  a successful framework for understanding  structure formation  in the universe.   Within this paradigm dark matter collapses first  into small haloes which merge to form progressively  larger haloes over time.  Galaxies  are thought to form out of gas which cools and collapses to the centers of these dark matter  haloes  (White  \\&  Rees  1978).  In the  standard picture of  disk galaxy formation the  structural and dynamical  properties of  disk galaxies  are expected  to  be strongly related to the properties of the  dark matter haloes in which they are embedded.   In  particular   the  characteristic  sizes  and  rotation velocities of  disk galaxies  are determined (to  first order)  by the spin  parameter, concentration parameter,  and size  of the  host dark matter   halo  (e.g.   Mo,   Mao  \\&   White  1998,   hereafter  MMW). Consequently, the detailed rotation curve shapes of disk galaxies can, in principle, be used to  constrain the structural properties of their dark matter haloes. This is especially true for low surface brightness (LSB) galaxies, which are believed to be dark matter dominated even at small radii.  A steadily increasing data base of observed LSB rotation curves has resulted in a heated debate as to whether the slopes of the inner  density profiles  of dark  matter dominated  disk  galaxies are consistent  with the cuspy  profiles found  in N-body  simulations, or similarly,  whether  the  inferred   concentrations  are  as  high  as predicted (see Swaters \\etal 2003 and references therein).  Unfortunately, determining cusp slopes and/or concentration parameters from mass modeling rotation curves is non-unique, even for dark matter dominated  galaxies   (e.g.,  Dutton  \\etal   2005).   In  particular, determining $c$ requires knowledge of the virial radius, which is hard to constrain  using data that only  covers the inner  $\\simeq 10\\%$ of the halo. As an alternative measure  of the central density of a halo, Alam, Bullock  \\& Weinberg (2002) introduced  a dimensionless quantity that  does  not require  knowledge  of  the  halo virial  radius,  and demonstrated  convincingly that  the observed  rotation curves  of LSB galaxies imply halo concentrations  that are systematically lower than predicted  for a  flat $\\Lambda$CDM  cosmology with  a  matter density $\\OmegaM=0.3$ and  a scale-invariant Harrison-Zeldovich power-spectrum with normalization $\\sigma_8=1.0$.  Further  observational  support  for  a  lower  normalization  of  the $c-\\Mvir$  relation  comes  from   the  zero  point  of  the  rotation velocity-luminosity relation, also  known as the Tully-Fisher relation (Tully  \\&  Fisher  1977)  of  disk galaxies  (van  den  Bosch  2000). Detailed disk formation models have clearly demonstrated that the high concentrations of  CDM haloes cause an overprediction  of the rotation velocities  at a  fixed disk  luminosity,  at least  for a  `standard' $\\LCDM$    cosmology    with    $\\OmegaM=0.3$,   $\\OmegaL=0.7$,    and $\\sigma_8=0.9$ (e.g., Dutton \\etal 2007; Gnedin \\etal 2006).  Although these  discrepancies may indicate  a genuine problem  for the CDM paradigm, there are a number of alternative explanations: First of all, as shown by various authors (e.g.  Swaters \\etal 2003; Rhee \\etal 2004;  Spekkens, Giovanelli, \\&  Haynes  2005)  the observed  rotation curves could be hampered by a variety of observational biases, such as beam smearing, slit offsets and inclination effects, all of which tend to underestimate the circular velocity in the central regions.  Secondly,  the dark matter  distribution could  have been  modified by astrophysical  processes   such  as  bars   (e.g.   Holley-Bockelmann, Weinberg, \\& Katz  2005) or dynamical friction (e.g.   Mo \\& Mao 2004; Tonini \\etal 2006). These processes tend to lower the concentration of the  dark  matter halo,  bringing  it  in  better agreement  with  the observations.   On  the   other  hand,  adiabatic  contraction  (e.g., Blumenthal \\etal 1986) thought to  be associated with the formation of disk galaxies, actually tends  to increase the halo concentration, and it remains to be seen whether the above mentioned processes are strong enough to undo  this contraction and still cause  a relative expansion of the inner halo (see Dutton \\etal 2007 for a detailed discussion).  A third option is that the data-model comparison has been made for the wrong  cosmology.   In  particular,   a  reduction  in  the  power  of cosmological density fluctuations on small scales causes a significant reduction of the predicted  halo concentrations (e.g.  Eke, Navarro \\& Steinmetz  2001; Zentner  \\& Bullock  2002; Alam  \\etal 2002;  van den Bosch, Mo, \\& Yang 2003).  Most data-model comparisons have been based on   a   flat    $\\Lambda$CDM   cosmology   with   $\\OmegaM=0.3$   and $\\sigma_8=0.9$.  However,  recently the  third year data  release from the WMAP mission has advocated  a model with $\\OmegaM \\simeq 0.25$ and $\\sigma_8 \\simeq  0.75$ (Spergel  \\etal 2006).  This  relatively small change in  cosmological parameters  causes a significant  reduction of the  predicted halo  concentration parameters,  bringing them  in much better agreement with the data (e.g., Yang, Mo, \\& van den Bosch 2003; van den Bosch, Mo \\& Yang 2003).  Another potentially important cause for the discrepancy are systematic errors  in the actual  model predictions.  Both Bullock  \\etal (2001a; hereafter  B01) and Eke,  Navarro \\&  Steinmetz (2001;  hereafter ENS) presented analytical  models that allow  one to compute the  mean halo concentration   for  given   halo  mass,   redshift  and   cosmology. Unfortunately, at  redshift zero the  predictions of these  models are divergent below  $\\sim 10^{11} \\Msun$,  with the ENS  model predicting halo  concentrations  that  are   significantly  lower.   This  is  of particular importance  for LSB (and  dwarf) galaxies with $V  \\lta 100 \\kms$, which  are thought  to typically reside  in haloes  with masses below this  value. Both  B01 and ENS  calibrated their  models against numerical  simulations. Those  of B01  probed the  mass  range between $3\\times  10^{11} h^{-1}  \\Msun$ and  $6\\times 10^{13}  h^{-1} \\Msun$, while those of ENS probed an even narrower range from $3\\times 10^{11} h^{-1}  \\Msun$ to  $3\\times10^{12} h^{-1}  \\Msun$ (albeit  with higher resolution).  What is needed  to discriminate between these models are simulations that resolve a large  population of low mass haloes, which is one of the main objectives of this paper.  Another important issue  that we wish to address in  this paper is the possibility  that  the  LSB  disk  galaxies that  have  been  used  to constrain  halo concentrations reside  in a  biased sub-set  of haloes (see discussion  in Wechsler \\etal 2006).   Numerical simulations have shown that there is a  significant scatter in both halo concentration, $c$, and  halo spin parameter, $\\lambda$,  at a given  halo mass (e.g. Bullock \\etal 2001a,b).  Thus if disk galaxies form in a biased subset of haloes, this  could lead to an apparent  discrepancy between theory and observation. In  fact, there are a number  of potential causes for such  a   bias.   First  of   all,  disk  galaxies  are   expected  to preferentially  form in haloes  that have  not experienced  any recent major  merger.  There is  evidence that  such a  subset of  haloes has higher mean $c$,  lower mean $\\lambda$, and lower  scatter in both $c$ and  $\\lambda$  (Wechsler  \\etal  2002;  D'Onghia  \\&  Burkert  2004). Clearly, this would worsen the disagreement between model and data. On the  other  hand,  it  has  also  been  suggested  that  LSB  galaxies preferentially  reside in haloes  with relatively  low concentrations. First  of  all,   since  disks  are  thought  to   be  in  centrifugal equilibrium,  less concentrated haloes  will harbor  less concentrated (i.e., lower  surface brightness) disk galaxies (e.g.,  MMW, B01).  In addition, using  numerical simulations Bailin \\etal  (2005) found that haloes   with  higher   spin  parameters   have,  on   average,  lower concentration parameters.  Since LSB  galaxies are thought to be those with  high spin  parameters such  a correlation  would imply  that LSB galaxies  reside in  haloes  with relatively  low concentrations.   If confirmed this could offer an alternative explanation as to why (some) LSB galaxies have lower concentrations than predicted.  Note, however, that previous  studies (B01, Navarro,  Frenk \\& White  1997, hereafter NFW),   have  found   no  correlation   between  spin   parameter  and concentration.  Another potential bias for disk  galaxy formation could arise if there is  a correlation  between  environment (defined  as  the large  scale matter  density) and $c$  or $\\lambda$.   In particular,  Harker \\etal (2006) found  evidence  that low  mass  haloes  in dense  environments assemble  earlier  than  haloes   of  the  same  mass  in  under-dense environments. Note  however, that for the  lowest density environments this trend reverses so that formation redshifts actually increase with decreasing  density.   Since  haloes  that  assemble  later  are  less concentrated  (Wechsler \\etal  2002), one  thus may  expect  a similar correlation between halo  concentration and environmental density.  If dwarf and LSB galaxies preferentially form in under-dense regions this could also help explain the lower than expected halo concentrations of these galaxies.  In this  paper we study galaxy size  dark matter haloes from  a set of cosmological N-body simulations with  the following goals: (i) to test the predictions  of B01 and  ENS regarding the halo  concentrations of low mass haloes  (down to $\\sim 3 \\times 10^9  h^{-1} \\Msun$), (ii) to determine the scatter in concentration, spin parameter, and halo shape at a  given mass,  (iii) to determine  whether there is  a correlation between the  spin and concentration parameters, and  (iv) to determine whether $c$,  $\\lambda$ and  halo shape depend  on the density  of the environment in which  the halo is located.  Our  paper is organized as follows.   Section~2  describes our  set  of  N-body simulations.   In Section~3 we  discuss how halo concentration, halo  spin parameter and halo shape  depend on halo mass.  Sections~4  investigates whether $c$ and  $\\lambda$  are  correlated,   while  Section~5  focuses  on  the environment dependence of halo  properties. In Section~6 we use simple models  for disk  formation  to investigate  whether  one expects  LSB galaxies to  reside in haloes  with a biased concentration  parameter. Finally, we summarize our results in Section~7. ",
        "conclusions": "\\label{sec:conclusion}  In this paper we have used a set of cosmological N-body simulations to study   the  scaling   relations,  at   redshift  zero,   between  the concentration   parameter,  $c$,   spin   parameter,  $\\lamp$,   shape parameter,  $\\qbar$, and  mass, $\\Mvir$,  of  a large  sample of  dark matter haloes.  Due  to the combined set of  simulations, we were able to  extend previous  studies to  a  mass range  at least  an order  of magnitude smaller, covering the full range of masses in which galaxies are expected to  form: $3\\times 10^{9} \\hMsun \\lta  \\Mvir \\lta 3\\times 10^{13} \\hMsun$.  For  this mass  range  we find  $c\\propto\\Mvir^{-0.11}$,  which is  in agreement with the model of Bullock \\etal (2001a), but in disagreement with the  model of Eke, Navarro  \\& Steinmetz (2001)  which predicts a significantly shallower  slope.  The single free parameter  of the ENS model only  controls the normalization  of the $c-\\Mvir$  relation, so that their model  cannot be tuned to fit the data.   The ENS model has also  been shown  to be  unable to  match the  slope of  the $c-\\Mvir$ relation for low mass haloes  at $z=3$ (Col{\\'{\\i}}n \\etal 2004).  For the Bullock \\etal  (2001a) model our data is well  fitted with a model with  $F=0.01$ and  $K=3.4  \\pm  0.1$ (where  the  error reflects  our estimate  of cosmic  variance).  Note  that this  normalization  is 15 percent lower  than the $K=4.0$  originally proposed by  Bullock \\etal (2001a), but it is in good agreement with Zhao \\etal (2003) and Kuhlen \\etal  (2005), who  found a  best-fit normalization  of  $K=3.5$. This discrepancy is at least partially due to an inconsistency with the use of transfer functions  in the work of Bullock et al.   If they use the same  transfer functions  to set  up the  initial conditions  of their simulations  and  to  compute   the  model  predictions,  they  obtain $K=3.75$.  This is however  significantly higher than can be accounted for by our estimate of  cosmic variance.  We find an intrinsic scatter in $c$  and $\\lamp$ at  fixed $\\Mvir$ of  $\\sigma_{\\ln c}=0.33\\pm0.03$ and $\\sigma_{\\ln  \\lam}=0.55\\pm0.01$, and  a median spin  parameter of $\\lamp=0.034\\pm0.001$,  all  in  good  agreement  with  Bullock  \\etal (2001a,b).  In an attempt  to distinguish between relaxed and  unrelaxed haloes we introduce a new and simple  parameter: $\\xoff$ which is defined as the distance between  the most bound particle  and the center  of mass, in units  of   the  virial  radius.   The  distribution   of  $\\xoff$  is approximately log-normal  with a median  $\\overline{x}_{\\rm off}\\simeq 0.04$.   The full  set  of haloes  shows  strong correlations  between $\\xoff$  and   the  residuals  of  the   $c-\\Mvir$  and  $\\lamp-\\Mvir$ relations, such  that haloes  with larger $\\xoff$  have a  larger than average $\\lamp$  and a lower  than average $c$.  Removing  haloes with large  $\\xoff$ therefore  results in  a higher  mean  concentration, a lower mean spin  parameter, and in less scatter  in both the $c-\\Mvir$ and $\\lamp-\\Mvir$  relations.  The median spin  parameter of `relaxed' (GOOD) haloes  is $\\lamp=0.030\\pm0.003$  with an intrinsic  scatter of $\\sigma_{\\ln\\lamp}=0.52\\pm0.01$.  The  average $c(\\Mvir)$ of `relaxed' haloes is  again in  good agreement  with the model  of B01,  but with $F=0.01$  and  $K=3.7$,  and  with  a  reduced  intrinsic  scatter  of $\\sigma_{\\ln c}=0.26\\pm0.03$.  A better  fit to the $c-\\Mvir$  relation for high  mass haloes ($\\Mvir \\gta 10^{13} \\hMsun$) can be  obtained with $F=0.001$ and $K=2.6$ (for all haloes), and  $K=2.9$ (for GOOD haloes). This is, however, at the expense of under predicting the concentrations for the low mass haloes ($\\Mvir \\lta 10^{11} \\hMsun$).  We also find  a strong correlation between the mean  axis ratio of the haloes,  $\\qbar$, and $\\xoff$,  such that  more prolate  haloes (i.e., those  with lower  $\\qbar$)  have higher  $\\xoff$,  on average.   This suggests  that the  majority  of  haloes with  small  axis ratios  are unrelaxed, and thus  that fragile LSB galaxies are  unlikely to reside in haloes  that are  strongly aspherical.  This  makes it  less likely that the discrepancy between observed and predicted rotation curves is due to  the fact that disks  are strongly elliptical,  as suggested by Hayashi \\etal (2004).  We have also investigated  the environment dependence of the residuals in $c$  and $\\lamp$.  This  is interesting because disk  galaxies, and LSBs   in  particular,   are  preferentially   found  in   regions  of intermediate to low density (Mo \\etal 1994; Rosenbaum \\& Bomans 2004). Defining  `environment' by  the matter  overdensity within  spheres of radii $1,2,4$ and $8h^{-1}$Mpc, we  find at the low mass end ($M<M^*$) that more  concentrated haloes live in denser  environments than their less concentrated  counterparts of the  same mass. This  is consistent with the studies of Wechsler \\etal  (2006) and Gao \\& White (2007) who found a similar trend using correlation functions. However, as we have shown, this trend is weak  compared to the scatter, which explains why Lemson \\&  Kauffmann (1999) did not notice  any environment dependence in their  simulation.  Contrary to  the halo concentrations,  the halo spin parameters  reveal no environment dependence at  fixed mass. This is at  odds with  the results of  Gao \\&  White (2007) who  found that haloes with  a large spin  parameter are more strongly  clustered than low spin haloes  of the same mass.  Finally we find  a weak trend that the most spherical haloes reside in slightly denser environments.  Finally,  using   a  simple  model  for  disk   galaxy  formation,  we investigated the properties of the  (expected) host haloes of LSB disk galaxies  (i.e., those  with  a central  surface  brightness $24  \\gta \\mu_{0,\\rm B}  \\gta 23 \\,{\\rm magn.} {\\rm  arcsec}^{-2}$). In addition to  having higher  than  average spin  parameters,  in agreement  with numerous  other studies  (e.g.,  Dalcanton \\etal  1997; Jimenez  \\etal 1998),  we  also  find that  the  host  haloes  of LSB  galaxies  have concentrations  that  are  biased   low  by  about  15  percent.  This correlation between halo concentration and disk surface brightness (or density) owes largely to a (weak) anti-correlation between $\\lamp$ and $c$, and to  the fact that centrifugal equilibrium  commands that less concentrated haloes  host less  concentrated disks.  The  amplitude of this correlation is significantly smaller than what has been advocated by Bailin \\etal  (2005), but this owes to the  fact that these authors did not  remove unrelaxed haloes,  which are unlikely to  host fragile LSB galaxies, from their analysis.  All these  results have important implications  for the interpretation of the halo concentrations inferred from LSB rotation curves. Numerous studies in the past have argued that these are too low compared to the predictions  for  a $\\Lambda$CDM  cosmology  (e.g.,  Alam \\etal  2002; McGaugh \\etal  2003; de Blok  \\etal 2003). However, there  are several reasons why we now believe  that the model predictions where too high. First of all, virtually all  previous predictions were made for a flat $\\Lambda$CDM  cosmology  with  $\\Omega_m=0.3$ and  $\\sigma_8=0.9$  (or $\\sigma_8=1.0$).  However, if one  adopts the cosmology favored by the three year data release of  the WMAP mission (Spergel \\etal 2006), one predicts  concentrations that  are  about 25  percent lower  (Macci\\`o \\etal,  in  preparation).   Compared  to  a  $\\Lambda$CDM  model  with $\\Omega_m=0.3$  and $\\sigma_8=1.0$ the  concentrations are  35 percent lower.   Secondly,   the  B01   model  (with  $F=0.01$   and  $K=4.0$) overpredicts the halo concentrations by $\\sim 15$ percent.  If we take into account that LSBs only  reside in relaxed haloes, this is lowered to a $\\sim  8$ percent effect.  And finally, one  needs to correct for the fact that LSB galaxies reside in a biased sub-set of haloes, which is  another 15  percent effect.   Combining all  these effects,  it is clear that the halo concentrations  predicted were almost a factor two too large.  This brings models and data into much better agreement."
    },
    "0608/astro-ph0608631_arXiv.txt": {
        "abstract": "The high-quality cosmological data, which became available in the last decade, have thrusted upon us a rather preposterous composition for the universe which poses one of the greatest challenges theoretical physics has ever faced: the so-called dark energy. By focusing our attention on specific examples of dark energy scenarios, we discuss three different candidates for this dark component, namely, a decaying vacuum energy or time-varying cosmological constant [$\\Lambda(t)$], a rolling homogeneous quintessence field ($\\Phi$), and modifications in gravity due to extra spatial dimensions. As discussed, all these candidates [along with the vacuum energy or cosmological constant ($\\Lambda$)] seem somewhat to be able to explain the current observational results, which hampers any definitive conclusion on the actual nature of the dark energy. % ",
        "introduction": "According to Einstein's general theory of relativity, the dynamic properties of a given space-time are fully determined by its total energy content. In the cosmological context, for instance, this amounts to saying that in order to understand the space-time structure of the Universe one needs to identify the relevant sources of energy as well as their contributions to the total energy momentum tensor $T_{\\mu \\nu}$. Matter fields (e.g., baryonic matter and radiation), are obvious sources of energy. Nevertheless, according to current observations, two other components, the so-called dark matter and dark energy (or \\emph{dark pressure}), whose origin and nature are completely unknown thus far, are governing the late time dynamic properties of our Universe. In the standard lore, the later actor plays a special role because it is supposed to be driving the current cosmic acceleration.  Dark energy has been inferred from a number of independent astronomical observations, which includes distance measurements of type Ia supernovae (SNe Ia) \\cite{perl, rnew, Legacy2005}, estimates of the age of the Universe \\cite{age}, measurements of the Cosmic Microwave Background (CMB) anisotropies \\cite{boom,wmap}, and clustering estimates \\cite{calb}. While the combination of the two latter results suggests the existence of a smooth component of energy that contributes with $\\simeq 3/4$ of the critical density, SNe Ia observations require such a component to have a negative pressure,  which generates repulsive gravity and accelerates the cosmic expansion. In addition to these recent observational results, the astonishing success of the inflationary Cold Dark Matter (CDM) paradigm in explaining high precision measurements of CMB anisotropy, galaxy clustering, the Ly$\\alpha$ forest, gravitational lensing and other astrophysical phenomena can also be thought of as an indirect but important evidence for the existence of a dominant repulsive dark energy component  In this short contribution we review some aspects of the so-called dark energy problem. This paper does not aim to be (and is far from being) an exhaustive or complete account of the problem (for that end we refer the reader to Ref. \\cite{revde}). Here, we focus our attention on three  possible mechanisms capable of explaining the current cosmic acceleration, namely, a decaying $\\Lambda$ term, the potential energy density associated with a dynamical scalar field and modifications in gravity due to extra dimensions effects. As discussed below, all these possibilities (together with the standard cosmological constant $\\Lambda$) seem somewhat to explain the current observational results, which makes the nature of dark energy a completely open question nowadays in Cosmology. ",
        "conclusions": "The results of observational cosmology in the past decade have opened up an unprecedented opportunity to establish a more solid connection between fundamental physics and cosmology. Surely, the most remarkable finding among these results comes from SNe Ia observations which suggest that the cosmic expansion is undergoing a late time acceleration. These SNe Ia results have been checked and confirmed by other cosmological observables, including the angular distribution of the 3-K CMB temperature and the clustering and dynamical estimates on the matter density parameter, which certainly makes a serious case for dark energy. Theoretical efforts to explain the nature of this dark energy component are vast and include, besides the possibilities discussed here, models with mass varying neutrinos \\cite{neutrinos}, holographic dark energy \\cite{elcio}, two-field quintom models \\cite{quintom}, models driven by dissipative process \\cite{mc}, unification dark matter/dark energy \\cite{cg}, among many others. %  Here, we have explored three specific scenarios of dark energy and shown that, although completely different from the physical viewpoint, all of them have attractive features capable of explaining the current cosmic acceleration and other recent astronomical observations. We emphasize that a possible way to distinguish among some of these dark energy candidates is through the observational limits on the cosmic EoS ($w$) and its time derivative ($w'$), i.e., if upcoming observations confirrm $w = 1$, then one will need to review the $\\Lambda$ problem, in order to fathom out why the formally infinite quantity $\\rho_{\\Lambda}$ is in fact so very small. On the other hand, if either $w \\neq \u22121$ or it is a time-dependent quantity($w' \\neq 0$) , then one will be able to rule out the cosmological constant as the actual nature of the dark energy and search for more realistic quintessence,  BW and other dark energy models. We expect that future observations along with theoretical developments will be able to decide this matter and shed some light on the origin and nature of the dark energy.  \\vspace{0.2cm}"
    },
    "0608/astro-ph0608407_arXiv.txt": {
        "abstract": "We present new weak lensing observations of 1E0657$-$558 ($z=0.296$), a unique cluster merger, that enable a direct detection of dark matter, independent of assumptions regarding the nature of the gravitational force law.  Due to the collision of two clusters, the dissipationless stellar component and the fluid-like X-ray emitting plasma are spatially segregated. By using both wide-field ground based images and HST/ACS images of the cluster cores, we create gravitational lensing maps which show that the gravitational potential does not trace the plasma distribution, the dominant baryonic mass component, but rather approximately traces the distribution of galaxies. An $8\\sigma$ significance spatial offset of the center of the total mass from the center of the baryonic mass peaks cannot be explained with an alteration of the gravitational force law, and thus proves that the majority of the matter in the system is unseen. ",
        "introduction": "We have known since 1937 that the gravitational potentials of  galaxy clusters are too deep to be caused by the detected baryonic mass and a Newtonian $r^{-2}$ gravitational force law \\citep{ZW37.1}.  Proposed solutions either invoke dominant quantities of non-luminous ``dark matter''  \\citep{OO32.1} or alterations to either the gravitational force law \\citep{BE04.1,BR06.1} or the particles' dynamical response to it \\citep{MI83.1}. Previous works aimed at distinguishing between the dark matter and alternative gravity hypotheses in galaxies \\citep{BU02.1,HO04.1} or galaxy clusters \\citep{GA02.1,PO05.1} have used objects in which the visible baryonic and hypothesized dark matter are spatially coincident, as in most of the Universe.  These works favor the dark matter hypothesis, but their conclusions were necessarily based on non-trivial assumptions such as symmetry, the location of the center of mass of the system, and/or hydrostatic equilibrium, which left room for counterarguments. The actual existence of dark matter can only be confirmed either by a laboratory detection or, in an astronomical context, by the discovery of a system in which the observed baryons and the inferred dark matter are spatially segregated.  An ongoing galaxy cluster merger is such a system.  \\begin{figure*} \\plottwo{f1a.new.jpeg2ps.ps}{f1b.new.jpeg2ps.ps} \\caption{ Shown above in the top panel is a color image from the Magellan images of the merging cluster 1E0657$-$558, with the white bar indicating 200 kpc at the distance of the cluster.  In the bottom panel is a 500 ks Chandra image of the cluster.   Shown in green contours in both panels are the weak lensing $\\kappa$ reconstruction with the outer contour level at $\\kappa  =$ 0.16 and increasing in steps of 0.07.  The white contours show the errors on the positions of the $\\kappa$ peaks and correspond to $68.3\\%$, $95.5\\%$, and $99.7\\%$ confidence levels.  The blue $+$s show the location of the centers used to measure the masses of the plasma clouds in Table \\ref{table2}.} \\label{fig1} \\end{figure*}  Given sufficient time, galaxies (whose stellar component makes up $\\sim 1-2 \\%$ of the mass \\citep{KO03.1} under the assumption of Newtonian gravity), plasma ($\\sim 5-15$\\% of the mass \\citep{AL02.1,VI06.1}), and any dark matter in a typical cluster acquire similar, centrally-symmetric spatial distributions tracing the common gravitational potential.  However, during a merger of two clusters, galaxies behave as collisionless particles, while the fluid-like X-ray emitting intracluster plasma experiences ram pressure.  Therefore, in the course of a cluster collision, galaxies spatially decouple from the plasma.  We clearly see this effect in the unique cluster 1E0657$-$558 \\citep{TU98.1}.  The cluster has two primary galaxy concentrations separated by 0.72 Mpc on the sky, a less massive ($T\\sim 6$ keV) western subcluster and a more massive ($T\\sim 14$ keV) eastern main cluster \\citep{MA02.1}.   Both concentrations have associated X-ray emitting plasma offset from the galaxies toward the center of the system. The X-ray image also shows a prominent bow shock on the western side of the western plasma cloud, indicating that the subcluster is currently moving away from the main cluster at $\\sim$4700 km/s.  As the line-of-sight velocity difference between the components is only $\\sim 600 $ km/s \\citep{BA02.1}, the merger must be occurring nearly in the plane of the sky and the cores passed through each other $\\sim 100$ Myr ago.  Two galaxy concentrations that correspond to the main cluster and the smaller subcluster have moved ahead of their respective plasma clouds that have been slowed by ram pressure.  This phenomenon provides an excellent setup for our simple test.  In the absence of dark matter, the gravitational potential will trace the dominant visible matter component, which is the X-ray plasma.  If, on the other hand, the mass is indeed dominated by collisionless dark matter, the potential will trace the distribution of that component, which is expected to be spatially coincident with the collisionless galaxies.  Thus, by deriving a map of the gravitational potential, one can discriminate between these possibilities.  We published an initial attempt at this using an archival VLT image \\citep{CL04.1}; here we add three additional optical image sets which allows us to increase the significance of the weak lensing results by more than a factor of 3.  In this paper, we measure distances at the redshift of the cluster, $z=0.296$, by assuming an $\\Omega_\\mathrm{m} =0.3, \\lambda =0.7, H_0=70\\mathrm{km/s/Mpc}$ cosmology which results in $4.413$ kpc/$\\arcsec$ plate-scale.  None of the results of this paper are dependent on this assumption; changing the assumed cosmology will result in a change of the distances and absolute masses measured, but the relative masses of the various structures in each measurement remain unchanged. ",
        "conclusions": "A key limitation of the gravitational lensing methodology is that it only produces a two-dimensional map of $\\kappa$, and hence raises the possibility that structures seen in the map are caused by physically unrelated masses along the line-of-sight.  Because the background galaxies reside at a mean $z\\sim1$, structures capable of providing a significance amount of $\\kappa$ must lie at $z\\la 0.8$.  By comparing the measured shear for galaxies divided into crude redshift bins using photometric redshifts \\citep{WI03.1}, we further limit the redshift of the lensing objects to $0.18<z<0.39$. This range is consistent with the cluster redshift, but corresponds to a large volume in which a structure unassociated with the cluster could exist and be projected onto the lensing map.  However, the number density of structures with these lensing strengths in blank field surveys is $\\sim 10^{-3}$ arcmin$^{-2}$ for the subcluster \\citep{WI06.1}, and an order of magnitude less for the main cluster, resulting in a $\\sim 10^{-7}$ probability of having two structures within a square arcminute of the observed cluster cores.  Further, all such lenses observed in these cosmic shear surveys are clusters with enough plasma and galaxies to be easily observable. There is no evidence, however, in our deep imaging for additional cluster sized concentrations of galaxies or of plasma hotter than $T\\sim 0.5$ keV (the lower bound of the Chandra energy band) near the observed lensing structures.  Another alternate explanation of the lensing signal is related to the fact that clusters form at the intersections of matter filaments \\citep{BO96.1}. In principle, one could imagine two line-of-sight filaments of intergalactic gas (too cool to be visible with Chandra and too diffuse to have cooled into stars) extending from the cluster at the locations of the weak lensing peaks. To explain the measured surface mass density, such filaments would have to be several Megaparsecs long, very narrow, and oriented exactly along the line of sight.  The probability of such an orientation for two such filaments in the field is $\\sim 10^{-6}$. Further, because the two cluster components are moving at a relative transverse velocity of 4700 km/s compared to the typical peculiar velocities in the CMB frame of a few hundred km/s, the filaments could coincide so exactly with each of the BCGs only by chance.  This is an additional factor of $\\sim 10^{-5}$ reduction in probability. While such projections become more important in non-standard gravity models because in such models the thin lens approximation breaks down \\citep{MO01.1} and structures with a given surface density produce a greater amount of lensing the more they are extended along the line-of-sight, two such projections would still have a $\\ll 10^{-8}$ probability. Finally, we mention that two other merging clusters, MS1054$-$03 \\citep{JE05.1} and A520 (in preparation), exhibit similar offsets between the peaks of the lensing and baryonic mass, although based on lensing reconstructions with lower spatial resolution and less clear-cut cluster geometry.  A final possibility is that some alternative gravity models may be able to suppress the lensing potential of the central peak in a multiple-peak system, as in \\citet{AN06.1}.  That work used a model of a gas disk located between two symmetric mass concentrations representing the galaxy subclusters.  In their $\\kappa$ map, derived in the TeVeS framework \\citep{BE04.1}, the relative signal from this disk may be suppressed, but would still be easily visible with the noise levels of our reconstruction.  Our $\\kappa$ map, however, has no evidence of any mass concentration between the two galaxy subclusters other than the small perturbations consistent with the gas mass contribution in Newtonian gravity.  Furthermore, such a suppression has also only been shown to work for symmetric systems which have the central peak directly between the two outer peaks. In 1E0657$-$558, however, the X-ray plasma, which would provide the central peak, lies north of the line connecting the two $\\kappa$ peaks.  Further, the absolute $\\kappa$ levels of the peaks observed in 1E0657$-$558 are in good agreement with those in systems with similar velocity dispersions and X-ray temperatures \\citep[e.g.][]{CL02.1} which have the gas and the galaxies coincident.  The $\\kappa$-to-light ratios are also consistent with those in normal clusters with coincident gas and galaxies. Therefore one would need to not only suppress the inner peak in the $\\kappa$ map relative to the two outer peaks in this system, but also enhance the strength of the outer peaks to make up for the missing plasma mass.  Any non-standard gravitational force that scales with baryonic mass will fail to reproduce these observations.  The lensing peaks require unseen matter concentrations that are more massive than and offset from the plasma. While the existence of dark matter removes the primary motivation for alternative gravity models, it does not preclude non-standard gravity. The scaling relation between $\\kappa$ and surface mass density, however, has important consequences for models that mix dark matter with non-Newtonian gravity:  to achieve the $\\sim 7:1$ ratio in $\\kappa$ between the dark matter + galaxy component and the plasma component (Table~\\ref{table2}), the true ratio of mass would be even higher (as high as 49:1 for a constant acceleration model; although MOND \\citep{MI83.1} would not reach this ratio as the dark matter density would become high enough to shift the acceleration into the quasi-Newtonian regime), making the need for dark matter even more acute.  Such high concentrations of dark matter, however, are extremely unlikely based on the measured X-ray plasma temperatures \\citep{MA02.1} and cluster galaxy velocity dispersions \\citep{BA02.1}.  The spatial separation of the dominant baryonic component in a galaxy cluster from the hypothesized dark matter produced during a cluster merger has enabled us to directly compare the dark matter hypothesis to one with only visible matter but a modified law of gravity.  The observed displacement between the bulk of the baryons and the gravitational potential proves the presence of dark matter for the most general assumptions regarding the behavior of gravity."
    },
    "0608/astro-ph0608294_arXiv.txt": {
        "abstract": "Despite the common picture of an early-type dwarf (dE) as a quiescent galaxy with no star formation and little gas, we identify 23 dEs that have blue central colors caused by recent or ongoing star formation in our sample of 476 Virgo cluster dEs. In addition, 14 objects that were mostly classified as (candidate) BCDs have similar properties. Among the certain cluster members, the dEs with blue centers reach a fraction of more than 15\\% of the dE population at brighter ($\\mb\\le16$) magnitudes. A spectral analysis of the centers of 16 galaxies reveals in all cases an underlying old population that dominates the mass, with $M_{\\rm old}\\ge90\\%$ for all but one object. Therefore the majority of these galaxies will appear like ordinary dEs within $\\sim$ one Gigayear or less after the last episode of star formation. Their overall gas content is less than that of dwarf irregular galaxies, but higher than that of ordinary dEs. Their flattening distribution suggests the shape of a thick disk, similar to what has been found for dEs with disk features in Paper I of this series. Their projected spatial distribution shows no central clustering, and their distribution with projected local density follows that of irregular galaxies, indicative of an unrelaxed population. This is corroborated by their velocity distribution, which displays two side peaks characteristic of recent infall. We discuss possible formation mechanisms (ram-pressure stripping, tidally induced star formation, harassment) that might be able to explain both the disk shape and the central star formation of the dEs with blue centers. ",
        "introduction": "\\label{sec:intro}   Early-type dwarf galaxies are the most numerous type of galaxy in clusters, suggesting that the vigorous forces acting within a cluster environment might actually be creating them. At least theoretically, this could come about via a morphological transformation of infalling galaxies, like e.g.\\ ram-pressure stripping of irregular galaxies \\citep{gun72}, or so-called harassment of late-type spirals \\citep{moo96}. While these scenarios would in principle also be able to explain the famous morphology-density relation \\citep[e.g.][]{dre80}, unambiguous observational proofs are difficult to obtain. Early-type dwarfs are characterized by their smooth, regular appearance, while any possible progenitor galaxy probably displays a different overall structure, before a morphological transformation occurs. Moreover, early-type dwarfs themselves are not a homogeneous class of objects. In addition to the classical dwarf ellipticals, \\citet{san84} introduced the class of dwarf S0 (dS0) galaxies, which were conjectured to have disk components, based on indications like high flattening or a bulge-disk-like profile \\citep{bin91}. Following up on the discovery of spiral structure in an early-type dwarf \\citep{jer00a} and on similar other discoveries, we identified 41 Virgo cluster early-type dwarfs with disk features in \\citet[][hereafter Paper I]{p1}, and argued that they constitute a population of disk-shaped objects. Furthermore, nucleated and non-nucleated dwarf ellipticals show significantly different clustering properties \\citep{bin87} and flattening distributions \\citep{ryd94,bin95}.  Aside from this structural heterogeneity, color differences have been reported as well \\citep{rak04,lis05}, hinting at a range of stellar populations in early-type dwarfs. While the classical dwarf ellipticals are typically considered not to have recent or ongoing star formation and no significant gas or dust content \\citep[e.g.][]{gre01,conIV}, the well known Local Group galaxy NGC 205 constitutes a prominent example for an early-type dwarf that does exhibit central star formation, gas, and dust \\citep{hod73}. More examples are NGC 185 \\citep{hod63} and the apparently isolated galaxy IC 225 \\citep{gu06}. Furthermore, \\citet{conIV} report a 15\\% HI detection rate for early-type dwarfs in the Virgo cluster. All these objects are morphologically early-type dwarfs, but form a heterogeneous family of objects, which we assign the common abbreviation ``dE'', thereby including galaxies classified as dwarf elliptical or as dS0.   Based on the general idea of gas removal due to effects like ram-pressure stripping, numerous observational studies have focused on comparing the properties of gas-rich dwarf irregular (dIrr) galaxies and gas-poor dEs and dwarf spheroidals (dSphs), attempting to confirm or reject a possible evolutionary relation between them. \\citet{thu85} found in his study of optical--near-infrared colors the metallicity ranges of dIrrs and Virgo dEs to be ``mutually exclusive''. In their spectroscopic analysis of oxygen abundances of planetary nebulae in Local Group dEs and dSphs, \\citet{ric98} also found a significant offset to the respective abundances of dIrrs. Similarly, \\citet{gre03} recently showed for the Local Group that this metallicity difference exists even when considering only the respective old stellar populations of dIrrs and dSphs, indicative of more intense star formation and enrichment in dSphs. \\citet{thu85} and \\citet{bin85} also noted that many Virgo dEs are nucleated, while none of the dIrrs are. \\citet{bot86} concluded that Virgo dIrrs would have to fade by $\\sim$1.5 mag in B to fall into the optical--near-infrared color range of dEs. However, the bulk of faded dIrrs would then have an effective surface brightness $\\mu_{\\rm B,e} > 25 $mag/arcsec$^2$, substantially lower than the observed values of dEs.  Despite these counter-arguments, potential dIrr/dE ``transition types'' have frequently been discussed \\citep[e.g.][]{fer94,joh97,kne99}. Recently, \\citet{vZe04a} pointed out the ``remarkable commonality'' between dEs and dIrrs with respect to their surface brightness distributions and metallicity-luminosity relations. These authors found significant rotation in several bright dEs, upon which they base their discussion about the possibility that these very dEs may have formed via ram-pressure stripping of dIrrs. While \\citet{vZe04a} admitted that dEs (in Virgo) might actually form through various processes, they argued that every dE must have been gas-rich and star-forming in the past, and would thus inevitably have been classified as dIrr. However, these properties also apply to blue compact dwarf (BCD) galaxies, which have often been discussed in the literature as potential dE progenitors.  An important difference between dIrrs and BCDs was highlighted by \\citet{bot86}, who characterized dIrrs as an ``odd combination of rather blue colors, yet quite low surface brightness,'' indicative of a low surface mass density. In contrast, BCDs have a much higher surface brightness, suggesting a higher surface mass density, since they are not bluer than the dIrrs. Whether or not BCDs could indeed be the gas-rich, star-forming progenitors of dEs has been discussed controversially \\citep[e.g.][]{bot86,dri96,guz96,pap96}. A glance at the color images\\footnote{Using the Sloan Digital Sky Survey Image List Tool, http://cas.sdss.org/astro/en/tools/chart/list.asp, Authors: J.\\ Gray, A.\\ Szalay, M.\\ Nieto-Santisteban, and T.\\ Budavari} of Virgo cluster galaxies that were classified as (candidate) BCDs reveals what might be one reason for this controversy: these objects do not constitute a homogeneous class, but they vary strongly in size, (ir)regularity, and the fraction of area dominated by blue light with respect to the total area of the galaxy. It might thus be more promising to simply look for plausible dE progenitors among the BCDs, instead of attempting to draw conclusions about this class as a whole.  Interestingly, \\citet{bot86} mentioned that the BCDs in their sample look similar to NGC 205. Could dEs with central star formation thus bridge the evolutionary gap from quiescent dEs to potentially star-bursting progenitors? To our knowledge, \\citet{vig84} were the first to identify a central star formation region in a dwarf elliptical outside the Local Group. Recently, \\citet{gu06} presented a similar dE with ongoing star formation in its center. In Paper I of this series \\citep{p1}, we identified nine Virgo early-type dwarfs with central irregularities likely to be caused by dust or gas, and found that all of them have a blue center. In the present paper (Paper II), we follow up on these objects, and systematically search the Virgo cluster for such dEs with blue centers. This is made possible by the publicly available data of the Sloan Digital Sky Survey (SDSS) Data Release 4 \\citep[DR4,][]{sdssdr4} which covers almost the whole Virgo cluster with multiband optical imaging and partly with spectroscopy. Our data and sample are described in Sects.~\\ref{sec:data} and ~\\ref{sec:sample}, respectively. Results from image analysis are presented in Sect.~\\ref{sec:images}, followed by the spectral analysis in Sect.~\\ref{sec:spectra}. Section~\\ref{sec:gas} focuses on the gas content of our galaxies. The systematic properties of dEs with blue centers are given in Sect.~\\ref{sec:stuff}. Their evolutionary role is discussed in Sect.~\\ref{sec:discuss}, followed by a summary and outlook in Sect.~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:discuss}   The SDSS enables for the first time a systematic study of several hundred dEs in the Virgo cluster with optical imaging and partly with spectroscopy. These data provide the basis for this series of papers, which aims at disentangling the various subpopulations of early-type dwarfs and uncovering their evolutionary histories. In Paper I, we found that dEs with disk features (dE(di)s) constitute a disk-shaped, unrelaxed dE population that is clearly different from classical dwarf ellipticals. In the current paper, we focus on another conspicuous feature that is common to several dEs: a blue center caused by recent or ongoing central star formation. We have shown that these dE(bc)s constitute a significant fraction (more than 15\\%) of bright dEs in the Virgo cluster. This number declines to almost zero beyond $\\mb > 16$, which is most likely a real decline and is not mainly due to signal-to-noise effects.  Two dE(bc)s of the main sample are detected in HI, with gas-to-baryonic mass fractions of 1\\% and $3-6$\\%, respectively. These values and the upper limits for several other dE(bc)s suggest an average gas content larger than that of ordinary dEs, but lower than that of dwarf irregulars. Note that we do not know where in the galaxies the gas is located: if it were concentrated in the central region, it would probably be able to fuel star formation for a longer time than if the gas were distributed homogeneously.  As soon as star formation has ceased, each dE(bc)'s color -- which provided the basis of its selection -- will become indistinguishable from that of an ordinary dE within $\\sim$ a Gigayear or less (see Sect.~\\ref{sec:sub_color}). However, the statistical properties of the dE(bc) \\emph{population} are unlike those of ordinary dEs: the projected spatial distribution and the flattening distribution of the dE(bc)s are similar to those of the dE(di)s and are different from those of dEs that have no blue centers or disk features. Both the dE(bc) and dE(di) population show no central clustering, which, along with the side-peaks of their velocity distributions, hints towards fairly recent infall. How long ago this infall could have taken place depends on relaxation timescales. \\citet{conI} derived a two-body relaxation time for the Virgo dEs of much more than a Hubble time. Even violent relaxation, which might only apply for the case of merging groups or subclusters, would take at least a few crossing times $t_{\\rm cr}$, with $t_{\\rm cr}\\approx 1.7$ Gyr for the Virgo cluster \\citep{bos06}. Thus, a dE population built out of infalling galaxies remains in an unrelaxed, non-virialized state for many Gyr. We conclude that the dE(bc)s most likely formed through infall of progenitor galaxies.  The shape of the dE(bc)s, as deduced from their flattening distribution, is hardly consistent with their being spheroidal objects, and instead implies that the dE(bc)s are rather thick disks, i.e.\\ oblate-shaped objects with intrinsic axial ratios around $\\sim$0.4. This is only somewhat thicker than the dE(di)s, for which we derived an axial ratio of $\\sim$0.35 in Paper I. How could these disk-shaped dEs be produced?  \\subsection{Formation scenarios} \\label{sec:sub_discuss_formation}   In the galaxy harassment scenario \\citep{moo96}, an infalling late-type disk galaxy gets transformed into a dE through high-speed encounters with massive cluster galaxies. This obviously leads to an increase in the axial ratio during the transformation process, since a disk is converted into a spheroid. However, a thick stellar disk may survive and lead to lenticular systems \\citep{moo96,mas05}. These may form a bar and spiral features, and retain them for some time, depending on the tidal heating of the galaxy \\citep{mas05}. Galaxy harassment thus appears to be a plausible scenario to explain the formation of disk-shaped dEs, and of dE(di)s in particular. Moreover, it predicts gas to be funneled to the center and form a density excess there \\citep{moo98}. While Moore et al.\\ compare this to the presence of a nucleus in many dEs, they admit that their simulations ``were not designed to probe the inner 200 pc''. The radii of many of the blue central regions of the dE(bc)s are only slightly larger than this value (see Figs.~\\ref{fig:pics1a} to \\ref{fig:pics1c}); consequently, the central gas density excess could well explain the blue centers. The harassment scenario thus describes a possible evolution of an infalling late-type disk galaxy to a dE(bc) and to a dE(di).  Another mechanism to form disk-shaped dEs could be ram-pressure stripping \\citep[e.g.][]{gun72} of dwarf irregulars. Depending on galaxy mass, the gas might be significantly removed except around the central region \\citep{mor00}, which would seem to be consistent with the central star formation of the dE(bc)s. As shown by \\citet{vZe04a}, several dEs have significant rotation, and could thus be the descendants of dIrrs, which are also known to be mostly rotationally supported, at least at the luminosities considered here. However, apart from the problems with this scenario discussed in Sect.~\\ref{sec:intro}, like the metallicity offset between dEs and dIrrs or the too strong fading of dIrrs, the flattening distribution of dIrrs shown by \\citet{bin95} is not quite like that of our dE(bc)s. Instead, dIrrs have a (primary) axial ratio $\\ge$0.5. On the other hand, significant mass loss due to stripped gas would be expected to affect also the \\emph{stellar} configuration of the galaxies and could thus possibly account for the difference.  As outlined in Sect.~\\ref{sec:intro}, several studies claimed that BCDs might be progenitors of dEs. Their flattening distribution, as analyzed by \\citet{bin95}, is somewhat more like that of our dE(bc)s, though still slightly rounder. These BCDs behave similarly to the galaxies of our additional sample -- which were mostly classified as (candidate) BCDs, but were selected only if their appearance was similar to the dE(bc)s. Overall, they have more extended blue regions, stronger emission lines, and clearly a higher gas content than the dE(bc)s of the main sample. They are also dominated by an underlying old population -- only one out of seven has $M_{\\rm old}<90$\\% -- and they have fairly regular outer shapes. Their spatial distribution also hints at an unrelaxed population, and their velocity distribution is asymmetric.  Tidally induced star formation in dIrrs \\citep{dav88} might be able to link BCDs and dEs, and could at the same time overcome the problems of the ram-pressure stripping scenario. The initially lower metallicity and surface brightness of a dIrr are increased by several bursts of star formation, during which the galaxy appears as blue compact dwarf (BCD). After that, it fades to become a dE. The last star formation burst might occur in the central region of the dwarf \\citep{dav88}, consistent with the appearance of the dE(bc)s.  \\subsection{Presence of nuclei} \\label{sec:sub_discuss_nuclei}  Many late-type spirals -- which act as dE progenitors in the harassment scenario -- host a compact nucleus \\citep[e.g.][]{boe04}. If the nucleus survives the morphological transformation, it would become the nucleus of the resulting dE(bc) or dE(di). A second scenario for nucleus formation in dEs is late star formation out of central gas \\citep[e.g.][]{oh00}, which in principle might be taking place in the centers of the dE(bc)s.  The fraction of nucleated galaxies among the dE(bc)s is 7/16 (44\\%), which is based on the VCC classification and has been verified by us through visual inspection of images and unsharp masks. We find one more dE(bc) to have ``multiple nuclei'' (VCC 0021), and one to have a possible nucleus (VCC 1512). If those two objects were counted as being nucleated, the fraction would increase to 56\\%. Among the dE(di)s, 26 out of 36 galaxies are nucleated (72\\%). Ordinary dEs (i.e.\\ excluding dE(bc)s and dE(di)s) have a nucleated fraction of also 72\\% among the brighter objects ($\\mb\\le 16$), which is the magnitude range of most dE(di)s and of all but one dE(bc). If we assumed a 72\\% nucleated fraction for the dE(bc)s as well, we would expect to find 11.5 nucleated dE(bc)s among our 16 objects, with a standard deviation $\\sigma = 1.8$. Our observed number thus lies within $1.4 \\sigma$ of the expected value if dE(bc)s had the same nucleated fraction as dEs and dE(di)s, and if the two uncertain objects were counted as nucleated. Thus, there is no significant difference between dE(bc)s, dE(di)s, and the other dEs with respect to the presence or absence of a nucleus. Still, the somewhat smaller number of nuclei in dE(bc)s could be a hint that nuclei are indeed forming in their centers. Of the galaxies in the additional sample, none displays a nucleus, but the significant amount of central gas and dust might leave the possibility of a hidden nucleus, or of a nucleus just being formed.   We point out that the ACS Virgo Cluster Survey \\citep{acsvir} finds a much higher frequency of compact stellar nuclei in early-type galaxies than \\citet{vcc} did, primarily due to the much higher resolution of space-based studies as compared to ground-based ones \\citep{cot05}. However, if we took into account these results, a dE classified as nucleated in the VCC would then simply be termed \\emph{a dE with a nucleus bright enough to have been detected by \\citet{vcc}}.   \\subsection{Presence of dE(bc)s in less dense environments} \\label{sec:sub_discuss_ngc205}  Early-type dwarfs with blue centers are not only present in the Virgo cluster. NGC 205, with its central region of young stars and central dust clouds \\citep{hod73}, is a well known local example of what we term a dE(bc); the same applies to NGC 185 \\citep{hod63}. On the one hand, NGC 205 might be considered as special case due to its clear signs of tidal interaction with M 31 \\citep{mat98}. On the other hand, tidal interaction might not be something special but instead something very common and even required for the formation of dEs, and of dE(bc)s in particular. While galaxy harassment is negligible in groups \\citep{may01}, a similar mechanism is provided there by the so-called tidal stirring scenario \\citep{may01a}, in which a dIrr that suffers repeated tidal shocks is transformed into a dE or a dSph. In this model, the galaxies with higher surface brightness can reach a central gas density excess like what was described above for harassment. Tidal stirring might thus provide a consistent explanation for dE(bc)s in groups.  Contrary to the above interpretation of NGC 185 and NGC 205 as possible result of tidal stirring, \\citet{gu06} recently presented an apparently isolated dwarf elliptical with a blue center (IC 225) at a distance of 20.6 Mpc. Its spectrum displays Balmer line emission, and based on its appearance it would doubtlessly enter our dE(bc) sample. The authors base their conclusion about the galaxy's isolation on their failure to find a potential companion within 30 arcmin, using the NED. This angular search radius corresponds to 180 kpc at their given distance. We increased the search radius, and found the small galaxy group USGC U124 \\citep{ram02} to lie at the same distance as IC 225 (within 25 km s$^{-1}$ in radial velocity), and at an angular separation of 145 arcmin, or 870 kpc. The brightest group member (NGC 0936) is an early-type spiral galaxy with $\\Mb\\approx-20.5$. Since 870 kpc would seem a rather large distance for IC 225 to be a bound companion, we test whether a single interaction could have occured in the past. As an example, if the relative velocity was 100 km s$^{-1}$, the encounter would have occured 8.5 Gyr ago, much too long for the current star formation to have been triggered by it.  While the tidal forces of smaller galaxies closer to IC 225 could possibly affect the galaxy's gas distribution, \\citet{bro04} argued that tidal interactions might not be necessary for activating star formation, and that instead, the dynamics of gas masses in a dark matter gravitational potential could be the primary trigger. Nevertheless, the main problem would be to explain how an \\emph{isolated dE} could have formed at all. Therefore, if IC 225 is indeed a truly isolated galaxy, one has to think of other mechanisms for dE formation than the ones we have discussed. It also demonstrates along with NGC 185 and NGC 205 that central star formation does not only occur in cluster dEs, and that, consequently, the mechanisms for dE formation might be similar in different environments. In fact, in the Local Group, where we can study galaxies at the highest angular resolution, many dSphs show star formation histories that extend over many Gyr. In all of these cases, the younger populations are more centrally concentrated than the old ones \\citep{har01}. Apart from interactions, this may also be a consequence of longer-lived gas reservoirs at the centers of these galaxies' potential wells.   Finally, we caution against calling the dE(bc)s ``dE/dIrr transition types'': this would suggest that every dE(bc) has a dIrr progenitor, which might not be the case as discussed above. Similarly, \\citet{gre03} argued for Local Group dwarfs that the so-called dIrr/dSph transition types are plausible progenitors of dSphs, while dIrrs themselves seem less likely.     We have presented a study of Virgo cluster early-type dwarf galaxies (dEs) with central star forming regions, based on photometric and spectroscopic data from the SDSS DR4. These ``dE(bc)s'' are not rare objects, but they reach a fraction of more than 15\\% among the bright ($\\mb < 16$) dEs. Their spatial distribution and their distribution with projected local density suggest that they are an unrelaxed population. Their flattening distribution is consistent with them being disk-shaped objects like the dEs with disk substructure (dE(di)s) identified in Paper I. Even in the very center, where their colors are bluest, 90\\% or more of their mass belongs to an old stellar population. Thus, they will appear like ordinary dEs within about one Gigayear after the end of their star formation. The gas content of the dE(bc)s is lower than in dwarf irregulars, but probably somewhat higher than in classical dEs, implying that at least in some dE(bc)s star formation might still continue for some time. Plausible formation mechanisms that could explain both the disk shape and the central star formation of the dE(bc)s are galaxy harassment \\citep{moo96}, which describes the transformation of infalling late-type disk galaxies into dEs, tidally induced star formation in dIrrs \\citep{dav88}, and possibly also ram-pressure stripping \\citep{gun72} of dIrrs and star formation induced by gas compression due to ram pressure.  We have started our series of papers on Virgo early-type dwarfs by describing two sorts of dEs with ``special features'', i.e., dEs with a blue center (this study) and dEs with disk substructure (Paper I). It is important to stress that we are not looking at one single population of dEs, with those features just being some extra ``flavour'', but that these dE subtypes constitute populations with distinct properties that differ from the rest of dEs. To complicate the issue even more, ordinary dEs (i.e., exluding dE(bc)s and dE(di)s) are probably not a homogeneous population either, since e.g.\\ the clustering properties of nucleated and non-nucleated dEs differ significantly \\citep{bin87}. Whether or not all of these dE subtypes simply reflect different evolutionary stages of one single class of galaxy, or whether they are indeed different classes of early-type dwarfs, will be discussed in detail in a forthcoming paper."
    },
    "0608/astro-ph0608588_arXiv.txt": {
        "abstract": "A $\\sim$50 ks  XMM-Newton observation of  \\src\\ has been carried out in September 2005, after almost three years during which no bursts were detected from this soft gamma-ray repeater. The 0.8-10 keV spectrum was well fit by a power law plus blackbody model with photon index $\\Gamma$=1.9$\\pm$0.1, temperature kT=0.47$\\pm$0.02 keV and N$_H$ = $(2.12\\pm0.08)\\times10^{22}$ cm$^{-2}$, similar to previous observations of this source. The flux was $\\sim5\\times10^{-12}$ \\flux , a factor 2 dimmer than the typical value and the smallest ever seen from \\src . The long term fading of the persistent emission has been interrupted by the recent burst reactivation of the source. A target of opportunity XMM-Newton observation performed in April 2006 showed a flux $\\sim$15\\% higher. This variation was not accompanied by significant changes in the spectrum, pulsed fraction and light curve profile. We searched for emission and absorption lines in the spectra of the two observations, with negative results and setting tight upper limits of 50--200 eV (3$\\sigma$), depending on the assumed line energy and width, on the equivalent width of lines in the 1-9 keV range. ",
        "introduction": "The few  X-/gamma-ray sources known as Soft Gamma-ray Repeaters (SGRs) are generally believed to provide the most convincing evidence for the existence of magnetars, i.e. neutron stars with magnetic fields well above the quantum critical value  $B_{c} = \\frac{m_{e}^{2}c^{3}}{e\\hbar}=4.4\\times10^{13}$~G \\citep{duncan92, paczynski92}.  The first sources of this class were discovered in the seventies as transient emitters of short (tens of ms) bursts of hard X-rays \\citep{mazets79,cline80,laros86}. Only in more recent years it was possible to identify and study in detail their counterparts in the classical 1--10~keV X-ray band. This led to the discovery of periodic pulsations and secular spin-down in SGR 1806--20 (P=7.5 s, \\citealt{kouveliotou98}) and \\src\\ (P=5.2 s, \\citealt{hurley99,kouveliotou99}) thus confirming the neutron star nature of this class of sources. Occasionally, SGRs emit very energetic giant flares, during which up to a few 10$^{46}$ ergs are released in a few tenths of a second. The extreme properties of these events, of which only three, each one for a different source, have been observed to date, are the main motivation for the magnetar interpretation \\citep{thompson95,thompson96}.  Among the four confirmed SGRs, SGR 1806--20 and \\src\\, offer the best prospects for detailed spectral and timing studies in the soft X-ray band (E$<$10 keV). Both these sources, despite alternating periods of bursting activity with intervals of quiescence lasting months or years, remained at flux levels of the order of 10$^{-11}$ \\flux. The other Galactic soft repeater, SGR 1627--41, had a similar X-ray flux when it was discovered as a bursting source in 1998 \\citep{woods99c}, but since then its luminosity decreased by a factor $\\sim$25 and it is now a rather faint source \\citep{mereghetti06}. SGR 0526--66 has a relatively high luminosity ($\\sim2\\times10^{35}$ \\lum), but being at the Large Magellanic Cloud distance, its  observed flux is only $\\sim$10$^{-12}$ \\flux\\ \\citep{kulkarni03}. Furthermore, its study with low spatial resolution instruments is complicated by the presence of diffuse emission from the surrounding supernova remnant N49.  We started in 2003 a long--term monitoring program to study the time evolution of the spectral properties of SGR 1806--20 using the XMM-Newton X-ray satellite. Thanks to its imaging capability and large effective area we could obtain spectra of a much better quality than those available from previous satellites. The emission of the 27 December 2004 giant flare was also a fortunate occurrence, since we could observe how the source properties evolved in the two years leading to the flare and how they changed after this dramatic event \\citep{mte05,tiengo05}.  A similar observational campaign could not be carried out for \\src. In fact this source is located in a sky region that, until recently, was not accessible to XMM-Newton observations due to technical constraints in the satellite pointing. Thus the first observation of \\src\\ with this satellite could   be obtained only in September 2005, during a long period  of inactivity (the last bursts before the XMM-Newton observation were reported in November 2002, \\citealt{hurley02}).  \\src\\ became active again in March 2006: bursts were observed by Swift \\citep{palmer06} and Konus-Wind \\citep{golenetskii06}. We therefore requested a target of opportunity XMM-Newton observation, that was carried out on 1 April 2006. ",
        "conclusions": "Thanks to the high sensitivity of the EPIC instrument on XMM-Newton we have obtained the first high quality spectra of the persistent X--ray emission from \\src , setting tight limits on the presence of emission and absorption lines.  In September 2005 the source was found at a luminosity level of 1.3$\\times10^{35}$ \\lum\\ (for d=15 kpc), a factor two smaller than the typical value observed in the past, and in line with the trend of  luminosity decrease already observed in the latest BeppoSAX observations performed in April 2002. The target of opportunity XMM-Newton observation of April 2006 showed that the decreasing luminosity trend in \\src\\ has been interrupted by the recent onset of bursts emission. However, the moderate flux increase was not associated with significant changes in the X--ray spectral and timing properties, probably because the source is, up to now, only moderately active. Future  observations with XMM-Newton will be essential to monitor the spectral and flux variations for this source, possibly in connection with its renewed bursting activity, as it has been successfully done for its twin source SGR 1806--20.     We thank N.Schartel and the staff of the XMM-Newton Science Operation Center for performing the target of opportunity observation. This work has been partially supported by the Italian Space Agency and INAF through contract ASI/INAF I/023/05/0 and by the MIUR under grant PRIN 2004-023189."
    },
    "0608/astro-ph0608541_arXiv.txt": {
        "abstract": "We present optical spectra of 45 intermediate mass Herbig Ae/Be stars. Together with the multi-epoch spectroscopic and photometric data compiled for a large sample of these stars and ages estimated for individual stars by using pre-main sequence evolutionary tracks, we have studied the evolution of emission line activity in them. We find that, on average, the H$\\alpha$ emission line strength decreases with increasing stellar age in HAeBe stars, indicating that the accretion activity gradually declines during the PMS phase. This would hint at a relatively long-lived (a few Myr) process being responsible for the cessation of accretion in Herbig Ae/Be stars. We also find that the accretion activity in these stars drops substantially by $\\sim$ 3 Myr. This is comparable to the timescale in which most intermediate mass stars are thought to lose their inner disks, suggesting that inner disks in intermediate mass stars are dissipated rapidly after the accretion activity has fallen below a certain level.  We, further find a relatively tight correlation between strength of the emission line and near-infrared excess due to inner disks in HAeBe stars, indicating that the disks around Herbig Ae/Be stars cannot be entirely passive.  We suggest that this correlation can be understood within the frame work of the puffed-up inner rim disk models if the radiation from the accretion shock is also responsible for the disk heating. ",
        "introduction": "There has been considerable interest in the evolutionary studies of circumstellar disks around young stars in the past few years. The primary motivation for such studies has been the realisation that young circumstellar disks are the potential sites of planet formation. There is now overwhelming observational evidence for the presence of circumstellar disks around a majority of young pre-main sequence (PMS) stars \\citep[e.g.][]{becksar96, mansar97, wein02, eisner03}. The observed disk masses, sizes and chemical composition of young circumstellar disks are analogous to that of the {\\it protosolar nebula} \\citep[e.g.][]{beck99, hill05}. How these disks evolve into planetary systems is one of the fundamental questions in astronomy.  In the early PMS phase, young stars actively accrete matter from the surrounding optically thick, gas-rich disk \\citep [e.g.][]{hart98} . As the star evolves, the disk accretion rate goes down and eventually the accretion is terminated. By the time the central stars reach the main sequence, young circumstellar disks surrounding them lose most of the material due to planet formation and other disk dispersal processes. Characterising this evolution of circumstellar disks, therefore, is critical to improve our understanding of the planet formation process and disk dispersal mechanisms and their associated timescales.  Temporal evolution studies of circumstellar disks use different observational diagnostics of the disks such as emission lines, continuum excess in the UV, near-infarred (NIR) excesses and millimiter and sub-mm excesses which trace different properties of the disks with varying efficacy. For example, emission lines and continuum excess in the UV trace level of accretion activity in young stars, NIR excess traces the hot inner disk and millimeter and sub-mm excesses trace cold dust in the outer disk.   Most disk evolution studies in the literature have been based on the NIR excess which traces the evolution of inner accretion disks \\citep[e.g.][]{strom89, skrut90, sek93, haisch01, haisch01b, hill03, hill05}. From the JHKL imaging studies of embedded and revealed clusters of different ages, \\citet{haisch01} have found that essentially all the stars in a cluster lose their disks in $\\sim$ 6 Myr. \\citet{hill05}, for a much larger sample of young clusters and associations obtained the median lifetime of the optically thick accretion disks to be 2-3 Myr, but with a large scatter, especially at young ages ($<$ 1 Myr). All these studies are consistent with the fact that the inner disks in most young stars do not last for more than 5 Myr.   How long does emission line activity persist in young stars? Emission lines are more closely related to accretion activity than the JHKL excess \\citep[e.g.][]{muz04}. Excesses at the NIR wavelengths can be produced even by passive disks if there is hot ( T $\\sim$ 1000 K) dust present close to the star. Therefore the evolution of accretion is better diagnosed with emission lines. In this contribution we study the evolution of emission line activity in intermediate mass young stars. With the data compiled for a large sample of Herbig Ae/Be( hereafter HAeBe) stars, which are PMS objects of intermediate mass ($1\\leq M/M_{\\odot} \\leq 8$), we study the evolution of emission line activity in these stars.   One of the prominent observational features of HAeBe stellar group and that which distinguishes these stars from the normal main sequence stars is the presence of emission lines in the spectra. Apart from the H$\\alpha$ line which is almost always in emission, the other emission lines that are often observed in the optical spectra of HAeBe stars are HeI ($\\lambda\\:5876 \\AA$ \\& $\\lambda\\:6678 \\AA$), OI ($\\lambda \\:7774 \\AA$ \\& $\\lambda\\: 8446 \\AA$) and CaII triplet ($\\lambda \\lambda \\:8498, 8542, 8662 \\AA$) \\citep{herbig60, cohenkuhi79, finkmun84, hp92, bohm95} and forbidden emission lines such as [OI] ($\\lambda\\:6300$ and $\\lambda\\:6364$) \\citep{herbig60, ham94, bc94, bh97, corray98, vieira03, acke05}.  In the low mass PMS T Tauri stars, the origin of emission lines is currently understood in the framework of magnetospheric accretion model where emission lines are believed to be formed in the magnetospheric accretion columns \\citep{us85, kon91, hhc94, mch98, muzerolle01}. However, in HAeBe  stars the situation is less clear. Nevertheless, recent high spatial resolution interferometric observations at millimeter (mm), sub-mm and NIR wavelengths have revealed the presence of gas and dust circumstellar disks with inner holes, and possibly with puffed up inner rims, around a large number of HAeBe stars \\citep{ngm00, mansar97, mansar00, petu03, ohashi05, mill01, eisner03, eisner04, monn05}. Observed disk properties are in general similar to those found for the disks surrounding T Tauri stars. Spectropolarimetric observations of linear polarization across emission lines in HAeBe stars have clearly demonstrated that the emission lines originate in a compact region close to the star \\citep{vink02, vink05}. Magnetic fields also have been detected in a few HAeBe stars \\citep{hub04, wade05}. Magnetospheric accretion model has now been succesfully applied to HAeBe stars \\citep{muz04}. This evidence supports the fact that emission lines in HAeBe stars are related to accretion.    We obtained optical spectra of 45 HAeBe stars in order to investigate the evolutionary trends in the emission line activity of these stars. We also compiled available photometric and spectroscopic data for a large number of these stars. Based on the multi-epoch photometric and spectroscopic data thus obtained for a sample of 91 stars, we studied the evolution of emission line activity in HAeBe stars. In Section 2 of the paper we describe our sample. Spectroscopic observations obtained are presented in Section 3. Main results of our study are presented in Section 4 and discussed in detail in Section 5. Finally, in Section 6 we summarize our conclusions. ",
        "conclusions": "In this paper we have studied the temporal evolution of emission line activity in intermediate mass PMS stars by compiling  multi-epoch spectroscopic and photometric observations for a large sample of HAeBe stars. We find that, on average, the H$\\alpha$ emission line strength decreases with increasing stellar age in HAeBe stars, suggesting that the accretion activity gradually declines during the PMS phase. This would hint at a relatively long-lived ( a few Myr) process being responsible for the cessation of accretion in intermediate mass stars thereby implying that the evolution of accretion and disk dissipation is probably not entirely stochastic in nature. We also show that  in most HAeBe stars the H$\\alpha$ emission has weakened considerably in $\\sim$ 3 Myr, indicating that the accretion activity in them has dropped significantly. Most intermediate stars also appear to lose their inner disks on similar timescale, suggesting that inner disks dissipate rapidly after the accretion activity has fallen below certain level.  Futher, we find a relatively tight correlation between emission line strength and NIR excess due to inner disks in HAeBe stars. Such a correlation strongly suggests a physical connection between accretion activity and the strength of emission from the pufffed-up inner rim of the disks. We suggest that this correlation can be explained if the radiation from the accretion shock is also responsible for heating the inner rim of the disk. At higher accretion rates, more will be the contribution from the accretion shock radiation which will increase the effective surface area of the inner rim, thereby increasing the NIR excess emission."
    },
    "0608/astro-ph0608030_arXiv.txt": {
        "abstract": "Hydrogen atmosphere white dwarfs with metal lines, so-called DAZs, show evidence for ongoing accretion of material onto their surfaces.  Some DAZs are known to have unresolved M dwarf companions, which could account for the observed accretion through a stellar wind.  I combine observed Ca abundances of the DAZs with information on the orbital separation of their M dwarf companions to infer the mass loss rate of the M dwarfs.  I find that for three of the six known DAZs with M dwarf companions, a stellar wind can plausibly explain the observed accretion on the white dwarfs assuming Bondi-Hoyle accretion of solar abundance stellar winds on the order of 10$^{-14}$-10$^{-16}\\Msun$\\ yr$^{-1}$.  The rest of the sample have companions with orbits $\\gtorder$ 1~AU, and require companion mass loss rates of $> 10^{-11}\\Msun$\\ yr$^{-1}$.  I conclude that there must be an alternative explanation for accretion of material onto DAZs with widely separated companions. The inferred winds for two of the close binaries are orders of magnitude smaller than typically assumed for the angular momentum loss of red dwarf-white dwarf pairs due to magnetic braking from a stellar wind and may seriously affect predictions for the formation rate of CVs with low mass companions. ",
        "introduction": "Hydrogen atmosphere white dwarfs with spectral absorption lines due to metals are a potential puzzle for white dwarf evolution.  Metals present in the upper atmossheres of white dwarfs with T$_{eff}$ $<$ 20.000 K generally diffuse out of the upper atmosphere on timescales ranging from days to 10$^6$ yr \\citep{paquette86}.  It was therefore expected that white dwarfs with cooling ages $\\gg$10$^{8}$~yr should show no detectable metal line absorption.  However a number of both hydrogen (class DA) and helium (class DB) white dwarfs show the evidence of heavy metals in their atmospheres (classes DAZ and DZ). The initial explanation for metal contamination was the accretion of solar abundance ISM material\\citep{dupuis93b,dupuis92,dupuis93a}.  But the lack of correlation between white dwarfs and dense ISM clouds, a large number of DAZs residing at $<$70 pc within the Local Bubble, the high inferred accretion rates for DAZs, and the lack of hydrogen in DZ atmospheres makes ISM accretion problematic\\citep{aannestad93,zuckerman98,zuckerman03}. Recent surveys of white dwarfs have turned up several new DAZs, and these have been used to put forth a separate ISM accretion argument for within the Local Bubble \\citep{koester05,koester06}.  In this case, nearby white dwarfs accrete continuously from clouds of warm, partially ionized ISM  Other explanations exist, however.  The direct accretion of heavy metals through cometary impacts with white dwarfs was invoked to explain the first DAZ, G74-7 \\citep{alcock86}.  This explanation can deliver a moderate amount of material to the white dwarf surface, but fails to explain high abundance DAZs with short settling times and dusty disks such as G29-38 and GD 362 unless a planetary system perturbs a steady supply of material close to the white dwarf \\citep{zuckerman87,koester97,holberg97,debes02,gianninas04}.  In a survey of DAs for metal lines, six of ten known unresolved DA+M dwarf systems were DAZs compared to 1/9 for double degenerates and 0/7 for resolved binaries\\citep{zuckerman03}.  This raises the possiblity that DAZs may be caused by hitherto unseen stellar or substellar companions that deposit material through a stellar wind onto the surfaces of the white dwarfs.  The evidence for apparently single DAZs with unknown companions is thin.  Only two known substellar companions to white dwarfs exist, in widely separated binaries \\citep{zuckerman92,farihi04}.  This is despite wide searches for unresolved stellar or substellar objects in close orbits \\citep{probst82,zuckerman92,dobbie05,farihi05}.  Furthermore, concerted searches of DAZs themselves show that stellar companions are ruled out at all separations and substellar companions ruled out for all but the closest orbital separations \\citep{debes05a,debes05b,debes05c}.  M Dwarf stellar wind mass loss rates are poorly understood and hard to study--upper limits can only be placed on the nearby M5.5 Proxima Centauri based on X-ray and Ly $\\alpha$ studies with upper limits of 6$\\times 10^{-14}M_{\\odot}$/yr (3 $\\dot{M}_\\odot$) and 4$\\times 10^{-15} M_\\odot$/yr (0.2 $\\dot{M}_{\\odot}$) respectively \\citep{wargelin02,wood01}.  Radio observations with marginal detections and spectroscopic observations of coronal mass ejections from low mass stars have been used to propose very high mass loss rates from dwarf stars, which has a theoretical motivation \\citep{mullan92,foing90,badalyan92}.  These results are controversial, however, and depend on the model assumptions made \\citep{lim96}. The recent results for Proxima Centauri suggest that perhaps M dwarf stars have winds comparable to or smaller than solar type stars, rather than the larger rates.  The strength of young M dwarf winds can affect the lifetime of gaseous material in disks that are important for planet formation and limit the presence of dusty disks \\citep{laughlin05,plavchan05}.  In orbit around white dwarfs, M dwarf winds are important for the origin of cataclysmic variables \\citep{g03}.  Under the assumption that the cause of the metal lines in DAZs is due to M dwarf companions, the stellar wind rates for known DAZ+M dwarf systems can be calculated.  These calculations require the known orbits of the M dwarf companions.  Three of the known DAZ+M systems have well known orbital periods and masses since they are detached transiting systems. The other three have recent HST Advanced Camera for Surveys (ACS) images that allow orbital information to be estimated.  In this paper I infer the stellar wind mass loss rates of the M dwarf companions to six DAZ+M dwarf binaries.   In Section \\ref{s1} I estimate the orbital semi-major axis for each binary.  In Section \\ref{s2} I use the semi-major axes and white dwarf diffusion coefficients to infer the mass accretion rate onto the white dwarfs and thus the mass loss rate of the companion M dwarfs.  Finally, in Section \\ref{s3} I discuss my results. ",
        "conclusions": "\\label{s3} Our results break down along the lines of those DAZs with known close companions (a$<$0.02 AU) and those with greater separations.  At smaller separations, the calculated mass loss rates agree reasonably well with the upper limit to the wind around Proxima Centauri \\citep{wargelin02,wood02}.  The mass loss rates I calculate are about two orders of magnitude smaller than Proxima Centauri's smaller upper limit, with the exception of WD 0419-487 which is comparable, though larger than the other two M dwarfs.  This larger rate could be explained by moderate Roche lobe overflow, evaporation of the companion by the DAZ, or efficient capture of the companion's wind by a magnetic field.  If WD 0419-487 was efficiently capturing all of its companion's wind, $\\dot{M}_{RD}$ would fall nicely in with those observed for the other two close binaries.  On the other hand, the mass loss rates determined for the widely separated companions are three to four orders of magnitude larger than the Solar wind.  This is despite a slightly higher uncertainty of the accretion rate onto the white dwarfs. Most of these uncertainties would conspire to create a higher accretion rate.  For WD 1210+464, a lower accretion rate is possible if the detected equivalent width corresponds to a lower abundance than assumed.  However, even at the smallest lower limit of the \\citet{zuckerman03} survey ([Ca/H]$\\sim$12.8), the inferred mass loss rate of the companion would be equivalent to the Solar Wind and two orders of magnitude higher than the close binaries.  It is possible that these systems are hierarchical triples with companions undetected by radial velocity observations but in orbits similar to the close binaries.  WD 1210+464 and WD1049+103 have F814W photometry consistent with single DAZs, neglecting their resolved companions. This strongly argues that any further unresolved companions would have to be quite dim and of low mass. Conversely, M dwarfs could have the super solar rates predicted by the earlier results, but in light of the estimated winds of Proxima Centauri and the three close M dwarfs this seems unlikely.  Furthermore, given the inferred total ages of the host white dwarfs, the companions would have either completely evaporated or lost a large fraction of their total mass.  The low mass loss rates for the three closest binaries has two possible interpretations.  Either the mechanism for accretion is suppressed relative to Bondi-Hoyle accretion by several orders of magnitude if one expects M dwarf winds to be similar to the Sun, or M dwarf winds are quenched even in situations where they are rotating quickly and should have significant activity due to strong magnetic fields.  Some evidence for the quenching of winds for low mass stars comes from \\citet{mohanty03}, who find that very late spectral type stars have lower indicators of activity due to a corona or chromosphere. This has been noted in studies of the angular momentum evolution of CVs, where fully convective companions were believed to have lost less angular momentum due to an inefficient dynamo process \\citep{durney93}.  If the winds from these low mass companions are orders of magnitdue smaller than the solar wind mass loss rate, this has significant implications for the angular momentum evolution of CV progenitors and CVs themselves.  Typically, it has been assumed that the angular momentum evolution of a CV or CV-progenitor occurred due to magnetic braking through a stellar wind, with angular momentum loss:  \\begin{eqnarray} \\label{eq:ang} \\dot{J} & = & -K_w \\left(\\frac{R_{sec}}{\\Rsun}\\right)^{2-N}\\left(\\frac{M_{sec}}{\\Msun}\\right)^{-N/3} \\left(\\frac{\\dot{M}}{10^{-14} \\Msun s^{-1}}\\right)^{1-2N/3}\\omega^3\\  \\mbox{for}\\ \\omega \\le \\omega_{crit} \\\\ & = & -K_w \\left(\\frac{R_{sec}}{\\Rsun}\\right)^{2-N}\\left(\\frac{M_{sec}}{\\Msun}\\right)^{-N/3} \\left(\\frac{\\dot{M}}{10^{-14} \\Msun s^{-1}}\\right)^{1-2N/3}\\omega \\omega_{crit}^2\\ \\mbox{for}\\ \\omega > \\omega_{crit} \\\\ \\end{eqnarray} where $\\omega_{crit}$ represents a cutoff rotation speed where the magnetic field saturates. This equation is a modified form of that used by \\citet{kawaler88}, following the work of \\citet{mestel87} and \\citet{weber67}.  This prescription was used by others to estimate the timescale for the evolution of single low mass stars, the onset of mass transfer in CV progenitors, and the timescale for the evolution of CVs \\citep{sills00, g03, andronov03}.  In all of these cases, it was implicitly assumed that the mass loss rate was approximately the solar wind mass loss rate.  $K_w$ and $N$ were essentially free parameters but were chosen to be 2.7$\\times$10$^{47}$ g cm s and 1.5 respectively in these works.  $K_w$ essentially collects all the uncertainties in the properties of the wind and magnetic field of a particular star and $N$ corresponds to the particular magnetic field geometry.  It is immediately apparent that for the given $N$, $\\dot{J}\\propto\\dot{M}$ and the consequent angular momentum loss rate for the observed DAZs would be 1-2 orders of magnitude smaller than would be calculated by Equation \\ref{eq:ang}.  The consequent timescale for these binaries to become CVs would be dominated by gravitational radiation, and would not be affected significantly by magnetic braking.  I also investigate a plausible orbital separation cutoff where a typical measured WD accretion rate can be maintained by an M dwarf wind.  Using a similar calculation as I used for my widely separated binaries, I place a cutoff for reasonable mass loss rates due to winds at the stellar wind upper limits for Proxima Centauri (See Table \\ref{resulttab}). Figure \\ref{fig:cutoff} shows the result for varying levels of mass accretion that should be typical for DAZs.  The lowest accretion rates allow plausible companions at wider separations, but the maximum allowed is $\\sim$1~AU.  It may be instructive to separate DAZs into three classes for ease of identifying the underlying mechanism for accretion.  The first class would be DAZs with companions in orbits $\\ltorder$1~AU.  These are most plausibly explained by a wind scenario and can be used to measure M dwarf winds to rates orders of magnitude smaller than other methods \\citep{wood01,wood02}.  The second class would encompass all singular WDs and WDs in binaries $\\gtorder$1~AU, with distances from the Sun $<$70~pc.  These DAZs cannot easily be explained by either ISM accretion or the companion wind explanation.  The third class of DAZs would be at distances $\\gtorder$70~pc and be either singular or have companions at separations $\\gtorder$1~AU.  These objects would either have to show ISM accretion through correlation with a known dense cloud or through spectroscopic evidence of a surrounding medium.  Otherwise they would have the same origin as those in the second class.  ISM accretion remains viable for the two binaries that are at larger distances, though one would expect these two objects to correspond with known clouds or show some other evidence of a fairly dense ISM component. WD~0354+463, however, is well within the Local Bubble. Both wind accretion and ISM accretion would have difficulty explaining this object.  Planetesimal accretion could still be possible. This can either occur through comet impacts and disruptions such as suggested by \\citet{alcock86} and \\citet{debes02}, or through the tidal disruption of asteroids as suggested by \\citet{jura03}. Relic planetesimals from a circumbinary ring around the two stars could provide the necessary materials.  The M dwarf could be slowly perturbing the planetesimals. Since extrasolar planets are known to exist in binary systems, the interaction between a planet and M dwarf could be perturbing planetesimals into WD crossing orbits.  If this is the case, the resevoir of planetesimals might be detectable with {\\em Spitzer} or at longer wavelengths.  Finally, there are roughly 15 other detached systems with M dwarf companions that have orbital information and masses for both the white dwarf and the companion that could be analyzed in the same way \\citep{ritter03}.  Medium to high resolution optical spectra could be used to find out how many of these objects possess metal lines, and whether the mass loss rates inferred for those systems match the binaries in my sample."
    },
    "0608/astro-ph0608206_arXiv.txt": {
        "abstract": " ",
        "introduction": "Cosmology currently poses some of the most stringent bounds on the properties of sterile neutrinos $\\nu_s$ -- Standard Model (SM) gauge group singlet fermions -- namely, their mass and mixing parameters with their active counterparts $\\nu_e$, $\\nu_\\mu$ and $\\nu_\\tau$.  In particular, by comparing the predictions of the yield of primordial light elements with observations, Big Bang Nucleosynthesis (BBN) efficiently constrains the possible manifestations of sterile neutrinos. Having a greater than 3 effective number of neutrinos would increase the rate of expansion of the universe and leads to larger $^4$He and Deuterium abundances. Even if steriles are produced only after the decoupling of the active neutrinos from the cosmological plasma, so that the effective number of neutrinos remains roughly constant at 3, they would lead to a depletion of electron neutrinos and anti-neutrinos, which also boosts the amount of $^4$He and Deuterium produced~\\cite{historical}. In addition, observations of the matter power spectrum of the Large Scale Structures (LSS) in the Universe limits the total mass in neutrinos today, including the sterile ones~\\cite{LesgourguesPastorReview}.  However, naturally light or cleverly lightened sterile fermions arise in many theories beyond the Standard Model. On the concrete experimental side, moreover, the result of the LSND experiment~\\cite{LSND} still stands. The most common interpretation requires the existence of a sterile neutrino with a mass $\\mathcal{O}(1 \\eV)$ mixing with both the electron and muon neutrino with angles $\\mathcal{O}(10^{-3})$. Such parameters are solidly ruled out by the constraints of standard BBN and by LSS~\\cite{CMSV,cosmonu,di bari,LSS}.  As well known, the same region of the parameters space is currently being scrutinized by the MiniBooNE experiment~\\cite{miniboone}, which is expected to deliver results soon. Should MiniBooNE confirm the existence of the LSND-like sterile neutrino, this would surely signal the need of a drastic revision of standard cosmology. If instead MiniBooNE strengthens the experimental exclusion contour beyond the LSND region, the bound from BBN on this particular mixing configuration is still stronger -- a few orders of magnitude in mass and mixing angle -- and these issues remain: in what ways and how much of the parameter space can be reopened?\\footnote{This could also be of importance for the $r-$process nucleosynthesis in SuperNov\\ae, as Ref.~\\cite{rprocess} finds that the region of successful synthesis in the presence of an active/sterile oscillation has some overlap with the LSND area.}  \\medskip  The bounds from cosmology are based on the assumption of ``standard evolution''. They relax considerably if some non standard modification is postulated. One such modification that has attracted much attention is to invoke the presence of a large primordial asymmetry for one or more of the active neutrinos seas. Of course, one might expect the value for a primordial asymmetry to be set by that of the baryonic asymmetry $\\eta = ( n_{\\rm B} - n_{\\rm \\bar{B}} ) / n_{\\gamma} \\simeq 6 \\cdot 10^{-10}$. To respect charge neutrality, the asymmetry in the charged leptons must be of the same order. But there is no such constraint on the asymmetries of the electrically neutral neutrinos, and they are therefore allowed to take values that can be orders of magnitude larger, up to $\\mathcal{O}(10^{-2})$~\\cite{bounds on xi}. Indeed, there are many examples of cosmological scenarios that lead to a large primordial neutrino asymmetry while preserving the observed smallness of the baryonic one, such as GUT scenarios, the Affleck-Dine mechanism, Q-balls, etc.~\\cite{scenariosLargeL}. From the purely phenomenological point of view, this is thus an important possibility and it has relevant consequences on the sterile neutrino bounds.  The basic idea is the following~\\cite{original Foot Volkas}. A non vanishing asymmetry $L_\\nu$ produces an additional term in the matter potential for the active neutrinos that is proportional to the asymmetry $L_\\nu$ itself. This increases the magnitude of the diagonal entries of the real part of the effective Hamiltonian written in the flavor basis relative to its off diagonal ones, thus suppressing the active neutrino oscillations into the sterile states. Since sterile neutrinos can only be created through active-sterile oscillations, this means they will be less efficiently or not produced at all in the Early Universe.  \\medskip  The main goal of this paper is to study quantitatively the extent to which such a mechanism works. Some aspects of this issue have been addressed before (see e.g.~\\cite{asymm, BargerHiding, dolgov villante, fuller}). We aim at an analysis that includes as much content of the full neutrino evolution equations as possible, such as the established active-active neutrino mixing, and that covers a broad range of parameters. We compute the effects in terms of observable cosmological quantities, such as the yield of the primordial elements, and compare them with the most recent results, which have gone through significant changes.  \\medskip  The paper is organized as follows. In Sec.\\ref{formalism} we describe the neutrino mixing formalism, identify the parameter spaces we are interested in, and present the evolution equations that we solve. In Sec.\\ref{results} we present our results and in Sec.\\ref{conclusions} we summarize our conclusions. Appendix \\ref{appendix} lays down some relatively standard details that govern the evolution of the temperature of the primordial plasma in the time of interest. ",
        "conclusions": "\\label{conclusions}  We investigated the extent to which introducing a primordial leptonic asymmetry $(n_{\\nu} \\neq n_{\\bar\\nu})$, much larger than the baryonic one, relaxes the bounds on light sterile neutrinos imposed by Big Bang Nucleosynthesis and Large Scale Structure, by suppressing via matter effects the production of the sterile states in the Early Universe. Our momentum-averaged neutrino evolution equations took into account the background cosmological evolution and included all the relevant features of neutrino mixing, both active-active and active-sterile, in the presence of a primordial leptonic asymmetry.  Working in the framework of constant leptonic asymmetry, we solved numerically the neutrino densities as a function of temperature and computed the relevant cosmological observables such as the yield of $^4$He and Deuterium and the neutrino contribution to the energy density in the Universe $\\Omega_\\nu h^2$, comparing them with the most recent determinations.  We applied this formalism to the study of the generic cases of a sterile neutrino mixing with a pure electron neutrino and with a pure muon one. For different values of the primordial asymmetry $L_\\nu$, taken to be the same for all the active flavors, we found the relaxed constraints on the parameter spaces. The results are presented in Fig.~\\ref{fig:esmus} across a broad range of mixing parameters.  We then studied the specific case of the LSND sterile neutrino, which mixes with a specific combination of electron and muon neutrinos. The constraint from standard cosmology in the absence of asymmetries are confirmed to rule out the entire LSND region, by both BBN $^4$He overproduction and the cosmological neutrino mass bound. We found that a primordial asymmetry of order $\\mathcal{O}(10^{-4})$ allows the LSND regions around $(\\Delta m^2_{\\rm LSND}, \\sin^22\\theta_{\\rm LSND}) \\simeq (1\\ {\\rm eV}^2, {\\rm few}\\ 10^{-3})$ to be reopened.   \\appendix"
    },
    "0608/astro-ph0608495_arXiv.txt": {
        "abstract": "If the dark matter particle is a neutralino then the first structures to form are cuspy cold dark matter (CDM) haloes collapsing after redshifts $z\\approx 100$ in the mass range $10^{-6}-10^{-3}\\, \\Msun$. We carry out a detailed study of the survival of these micro-haloes in the Galaxy as they experience tidal encounters with stars, molecular clouds, and other dark matter substructures.  We test the validity of analytic impulsive heating calculations using high resolution $N$-body simulations. A major limitation of analytic estimates is that mean energy inputs are compared to mean binding energies, instead of the actual mass lost from the system.  This energy criterion leads to an overestimate of the stripped mass and underestimate of the disruption timescale since CDM haloes are strongly bound in their inner parts. We show that a significant fraction of material from CDM micro-haloes can be unbound by encounters with Galactic substructure and stars, however the cuspy central regions remain relatively intact. Furthermore, the micro-haloes near the solar radius are those which collapse significantly earlier than average and will suffer very little mass loss. Thus we expect a fraction of surviving bound micro-haloes, a smooth component with narrow features in phase space, which may be uncovered by direct detection experiments, as well as numerous surviving cuspy cores with proper motions of arc-minutes per year, which can be detected indirectly via their annihilation into gamma-rays. ",
        "introduction": "If dark matter is composed mainly of the lightest supersymmetric partner particle, the neutralino, the first self-gravitating structures in the Universe are Earth-mass haloes forming at high redshifts \\citep*{hofmann_etal01,diemand05}. As many as $10^{15}$ could be within our Galactic halo today. These abundant cold dark matter (CDM) micro-haloes have cuspy density profiles that can withstand the Galactic tidal field at the solar radius. The numbers of such haloes that lie within the vicinity of the solar system depends on how many survive the complex merging history of early hierarchical structure formation. $N$-body simulations of CDM satellites indicate that tightly bound cusps are very stable against tidal stripping (Kazantzidis et al. 2004), and therefore dense micro-haloes accreting late onto more massive structures may survive relatively intact. The exact distribution of dark matter in the solar vicinity is important for direct and indirect dark matter detection experiments.  Substructures that survive the merging process will experience continuous perturbative encounters with stars, molecular clouds, and other dark matter subhaloes. As discussed in \\citet{diemand05}, we expect that these encounters lead to some mass loss but that the cusps of most micro-haloes remain intact. Recent studies by \\citet{zhao_etal05a, zhao_etal05b}, \\citet{green_goodwin06} and \\citet{berezinsky} have raised the question whether these first haloes would be completely disrupted by close encounters with stars. Crossing the Galactic disc would also cause additional tidal heating. \\citet{moore_etal05} argued that the analytical impulse approximation and the semi-analytic models used in these studies may not fully describe the disruption of the high-density inner cores. Particle orbits deep in the cusp may remain adiabatically invariant to the perturbations and preserve the structure of the cusp. Only direct numerical simulations can describe these complex dynamical processes. In this paper, we use several sets of high resolution $N$-body simulations to test the validity of analytical heating models.  An important factor in the survival statistics of micro-haloes is how many survive similar-mass mergers during the build-up of the Galactic halo \\citep{diemand2006}. Even if only a few percent survive the hierarchical growth, many micro-haloes would still lie within one parsec from the Sun. Their dense cuspy cores would be sources of gamma-ray emission due to self-annihilation, which could be uniquely distinguished by their high proper motions on the sky of the order arc-minutes per year (Moore et al. 2005, Koushiappas 2006). ",
        "conclusions": "We have studied the disruption of dark matter micro-haloes by stars and other substructures using both analytical impulse approximation and self-consistent $N$-body simulations. The analytic calculations presented here are quite similar to those of \\citet{green_goodwin06} and we come to similar conclusions. Our calculations differed in that we used more realistic cuspy n-body models and we studied how the internal structure of these systems evolve due to perturbations. See also \\citet{angus} for an independent and complementary study. Earlier studies  e.g. \\citet{aguilar86}, also studied cuspy systems and found similar robustness to tidal heating in the central regions as we find. However our resolution allows us to study the response of CDM haloes deep within their central regions.  \\begin{itemize}  \\item The impulse approximation predicts that those micro-haloes in the solar vicinity which formed after z = 130 (about 80\\% of the local micro-halo population) should lose most of their mass due to close encounters with disc stars.  \\item Numerical simulations of individual encounters demonstrate that the usual condition of disruptive heating used in analytical studies, $\\Delta E = E_{\\rm b}$, does not lead to complete dissolution of haloes with cuspy density profiles. For the inner logarithmic slope $\\gamma = 1.2$, on average only 30\\% of the mass is lost from the system for this energy change. The relation between the fractional mass loss and the energy input in the tidal approximation is given by equation (\\ref{eq:dm}): $\\Delta M/M \\approx (1 + 2 (\\Delta E/E_{\\rm b})^{-1})^{-1}$.  \\item The change of particle energies, after the system settles into a new virial equilibrium following the tidal encounter, is described accurately by the extension of the impulse approximation accounting for virial oscillations. An apparent resistance to tidal heating of the material deep in the cusp is due to the high binding energy inside the cusp.  \\item Repeated tidal encounters lead to diminishing mass loss from the same micro-halo.  After 10 identical encounters at impact parameter $b = 2 R_{\\rm vir}$, 10\\% of the halo still remains self-bound even though $\\Delta E = 5E_{\\rm b}$.  \\end{itemize}  Near the solar radius within the Galaxy most of the mass of the micro-haloes is tidally removed. This material forms cold streams in phase space providing a unique signal for direct detection experiments. The dense cuspy cores of these haloes survive reasonably intact, although the mass loss leads to a reduction in annihilation products of about a factor of only two to three. These cores could be distinguished by their high proper motions on the sky of the order arc-minutes per year."
    },
    "0608/astro-ph0608340_arXiv.txt": {
        "abstract": "KASCADE-Grande is a new extensive air shower experiment co-located to the KASCADE site at Forschungszentrum Karlsruhe. The multi-detector system allows to investigate the energy spectrum, composition, and anisotropies of cosmic rays with unprecedented prevision in the energy range from $10^{14}$-$10^{18}$~eV. The primary goals besides investigating the origin of the knee at $E \\simeq 3 \\cdot 10^{15}$~eV, are to verify the existence of the second knee at $E \\sim 10^{17}$~eV and to measure the composition in the expected transition region of galactic to extragalactic cosmic rays.  The performance of the apparatus and shower reconstruction methods will be discussed on the basis of detailed Monte Carlo simulations and first data. First results based on slightly more than a year of data taking are presented. ",
        "introduction": "Despite intense research over several decades and measuring an impressive set of cosmic ray (CR) data, the origin of non-solar CRs remains unknown.  As an example about the available data, Fig.\\,\\ref{fig:knee} shows the so called {\\em all-particle} spectrum.  It has become standard practice to present the spectral data as the flux $dJ/dE$ times a power of the energy $E$, so that the visibility of spectral features is emphasized. The most prominent feature in the spectrum is the {\\em knee} at about $3\\cdot10^{15}$~eV where the index $\\gamma$ of the power law spectrum $dJ/dE \\propto E^{-\\gamma}$ changes from $\\gamma \\simeq 2.75$ to $\\simeq 3.1$.  The origin of this phenomenon is still unknown, but it is generally attributed to either the limiting energy of galactic CR accelerators and/or to the onset of strong diffusion losses of particles out of our Galaxy.  Also, particle physics explanations have been put forward, either assuming the onset of a new interaction mechanism above the knee energy so that part of the CR-energy remains invisible to experiments or by considering propagation effects between the sources and the solar system.  Generally, supernova remnants (SNR) are conceived to be the major sources for CRs up to the knee with the acceleration taking place by the first order Fermi mechanism.  Based on simple dimensional arguments, the maximum energy would then be given by $E_{\\rm max} \\propto Z \\cdot R \\cdot B$, with $Z$ being the charge of the CR particle and $R$ and $B$ being the size and magnetic field of the source, respectively.  As a consequence, the knee position for different CR primaries would be expected to scale with their atomic number $Z$.  Similarly, diffusion losses out of the Galaxy would result in a $Z$ dependence of the knee position, since magnetic confinement is the underlying process.  For a recent review on astrophysical models, the reader is referred to \\cite{Hillas05} and references therein.  Particle physics models generally predict a scaling of the knee energy with the atomic mass $A$, since the reaction mechanism at high energies is governed by the energy per nucleon rather than by the total energy of the CR particles.  More recently, the cannonball (CB) model was suggested to explain the origin of all non-solar CRs \\cite{Dar06}.  It predicts a scaling of the knee position with $A$ rather than with $Z$, because of deflections of CRs at the relativistically moving CBs.  In any of the astrophysical models, light particles drop out of the acceleration and/or confinement region first -- this is what should cause the knee -- so that the composition is expected to change from light to heavy when crossing the knee energy.   \\begin{figure*}[t] \\centering \\vspace*{0.1cm} \\includegraphics[width=100mm]{allparticle-spectrum.eps} \\caption[xx]{The all-particle CR energy spectrum weighted by $E^{3}$ showing the knee at $3\\cdot10^{15}$~eV, a possible second knee at $\\sim 10^{17}$~eV, and the ankle at about $3\\cdot10^{18}$~eV. The energy range covered by KASCADE-Grande is indicated by the dashed box.} \\label{fig:knee} \\end{figure*}   A na\\\"{i}ve extrapolation of this idea suggests that the putative second knee would mark the position of the Fe-knee.  No consensus on a preferred accelerator site or mechanism exists for the energies between the (second) knee and ankle.  It has long being argued that CRs above the ankle are of extragalactic origin, as they cannot be isotropized by the Galactic magnetic fields, but their arrival directions are surprisingly isotropic.  At the highest energies, above $5 \\cdot 10^{19}$~eV, the GZK\\footnote{Greisen, and Zatsepin \\& Kuzmin} cut-off is expected.  Because of the low flux involved, such energies cannot be measured with KASCADE-Grande.  This is the domain of the Pierre Auger Observatory \\cite{Auger04}. ",
        "conclusions": "KASCADE-Grande has started to take data.  All components are working stable and perform well.  The presently installed FADC system will extend its capabilities and provide detailed information about the longitudinal profile of the shower disk and allow for additional electron - muon separation.  The LDFs of both electrons and muons are found to agree well with EAS simulations.  From their obtained electron and muon numbers a first estimate of the primary energy spectrum could be derived. Within the overlap region with KASCADE, a good agreement is found.  However, the statistics beyond $3 \\cdot 10^{17}$\\,eV is still to poor to allow conclusions about the second knee.  This will change in the near future with more data becoming available. KASCADE-Grande keeps the multi-detector concept for tuning different interaction models at primary energies up to \\mbox{$10^{18}$ eV}.  It also provides a perfect environment for detecting radio emission in extensive air showers, which is the aim of the LOPES project~\\cite{lopes}.  \\medskip \\noindent {\\bf Acknowledgments} \\\\ KHK wishes to thank the organizers for their invitation to the Vulcano workshop conducted in a very pleasant and fruitful atmosphere.  KASCADE-Grande is supported by the Ministry for Research and Education of Germany, the INFN of Italy, the Polish State Committee for Scientific Research (KBN grant for 2004-06) and the Romanian National Academy for Science, Research and Technology."
    },
    "0608/hep-ph0608055_arXiv.txt": {
        "abstract": "We investigate dark matter candidates emerging in recently proposed technicolor theories. We determine the relic density of the lightest, neutral, stable technibaryon having imposed weak thermal equilibrium conditions and overall electric neutrality of the Universe. In addition we consider sphaleron processes that violate baryon, lepton and technibaryon number. Our analysis is performed in the case of a first order electroweak phase transition as well as a second order one. We argue that, in both cases, the new technibaryon contributes to the dark matter in the Universe. Finally we examine the problem of the constraints on these types of dark matter components from earth based experiments. ",
        "introduction": "The origin of the dark matter is one of the most intriguing open problems of modern particle physics and cosmology. A few decades ago, astronomers realized that it was impossible to explain the motion of galaxies, just by accounting the luminous part of the galaxy. The most important observational evidence comes from the rotation curves of spiral galaxies. Astronomers are able to estimate the velocities of clouds of hydrogen atoms, just by looking at the Doppler shifted 21 cm emission. Naively one would expect by using Newton's law, that the velocity of clouds like these should fall as $\\propto r^{-1/2}$, when $r$ becomes larger than the radius of the luminous part of the galaxy ($r$ being the distance to the center of the galaxy). However, astronomers observe an almost constant velocity, independent of $r$. Unless one proposes a radical change of the laws of gravity or of motion, it is logical to assume that matter that does not interact ``too much'' and therefore appears dark to us, has to fill the space of the galaxy.  Once we assume that a type of dark matter (DM) is responsible for the discrepancies of the motion of the galaxies, then there are two distinct possibilities for the origin of the DM. The first type of candidates are of baryonic origin and they are called MACHO (Massive Compact Halo Objects). The MACHO can be brown dwarf stars, giant planets and massive black holes. Massive black holes with masses near $100M_{\\bigodot}$ can be remnants at the center of supernovae explosions. The brown dwarf stars are stellar objects with masses less than $0.1M_{\\bigodot}$. Since a proton star needs at least a mass of $0.1M_{\\bigodot}$ to ignite nuclear fusion, brown dwarfs never get to begin the nuclear fusion of hydrogen. Giant planets can be also a component of MACHOs with masses of the order of $0.001M_{\\bigodot}$. However, observations so far showed with a high level of confidence that MACHOs cannot account for more than $20\\%$ of the DM \\cite{Alcock:2000ph}.  The second possibility is to have matter of nonbaryonic origin. In contrast with MACHOs that are compact objects of baryonic matter, in this case we have particles that are neutral and interact only through gravitational and weak forces. The name WIMP (Weakly Interacting Massive Particles)  is frequently used for these particles. The Standard Model does not have the particles with the desired properties. Ordinary neutrinos are the only electrically neutral objects that interact weakly within the SM. They are, however, too light and would compose part of a {\\it hot}-type dark matter, which is usually not considered in a viable cosmological model. This is so since hot dark matter would smear out large scale structure of galaxies. Supersymmetry, for example, provides some natural WIMP candidates for cold dark matter such as neutralinos.  In the search of WIMP candidates, particles related to technicolor theories were also investigated as possible sources of cold DM. This idea was pioneered by Nussinov in \\cite{Nussinov:1985xr,Nussinov:1992he} and further investigated in \\cite{Barr:1990ca}. He imagined that DM could be accounted by a technibaryon asymmetry which can be ultimately related to the ordinary baryon asymmetry. Although it was clear that technicolor can give ``reliable'' DM candidates, little interest was shown in the past because of the severe problems that most technicolor theories suffer from such as large flavor changing neutral current processes and/or problems with the Electroweak Precision Measurements. Progress in this direction has been made recently. We have constructed explicit technicolor extensions of the SM passing the precision tests \\cite{Sannino:2004qp,Hong:2004td, Dietrich:2005jn,Gudnason:2006ug}. Because of their walking nature they have very much reduced flavor changing neutral current processes. Therefore it does make sense to revisit the possibility of technibaryons as components of DM. We have started this analysis in \\cite{Gudnason:2006ug} for the models with technifermions in the adjoint representation of the technicolor gauge group. In this case the technibaryon is one of the would-be Goldstone bosons of the underlying technicolor theory. In Ref.~\\cite{Appelquist:2003hn} has been also suggested that technicolor theories may lead to dark matter candidates of similar nature.  The goal of the present work is threefold: First, we will provide the basic set up while showing the detailed computations needed to determine the present relic density of our technibaryon DM candidate. Second, we consider two different orders of the electroweak phase transition and compare the results. Third, using our low energy effective theory developed in \\cite{Gudnason:2006ug} we compute the relevant cross sections needed to examine the problem of the constraints on these types of DM components from earth based experiments.   In the next section we review the basic properties of the new technicolor theory and present the lightest neutral technibaryon (LTB) of the theory. In section III we present our main calculation of the technibaryon contribution to DM. In section IV we investigate the experimental constraints and comment on the detection of our technibaryon WIMP. Finally we briefly conclude in section V. ",
        "conclusions": "We have investigated in much detail dark matter candidates emerging in our recently proposed technicolor theories. We have determined the contribution of the lightest, neutral, stable technibaryon to the dark matter density having imposed weak thermal equilibrium conditions and overall electric neutrality of the Universe. Sphaleron processes have been taken into account. We performed the analysis in the case of a first order electroweak phase transition as well as a second order one. In both cases, the new technibaryon contributes to the dark matter in the Universe. We have also examined the problem of the constraints from earth based experiments. {We find that quite a substantial amount of DM can be explained within our model}. The new generation of DM-detection experiments can put very strict limits or even rule out the present type of DM component. {We should stress that our framework can be applied to any model featuring a new baryonic type particle at the electroweak scale whose new baryon-type charge is violated only by weak interactions.}"
    },
    "0608/astro-ph0608710.txt": {
        "abstract": "%\\vspace {5pt} % The inner couple hundred pcs of our Galaxy is characterized by significant amount of synchrotron-emitting gas, which appears to co-exist with a large reservoir of molecular gas.  Many of the best studied sources in this region exhibit a mixture of 6.4 keV Fe K$\\alpha$ emission, molecular line emission and nonthermal radio continuum radiation.  The spatial correlation between\u00a0 fluorescent~Fe K-$\\alpha$ line \u00a0emission at 6.4 keV and molecular line emission~ \u00a0from Galactic center molecular clouds has been explained as reflected X-rays from a past outburst of Sgr~A*. Here we present multi-wavelength study of a representative Galactic center cloud Sgr C using {\\it Chandra}, {\\it VLA} and {\\it FCRAO}.  We note a correlation between the nonthermal radio filaments in Sgr C and the X-ray features, suggesting that the two are related.  This correlation, when combined with the distribution of molecular gas suggests against the irradiation of Sgr~C by Sgr A*.  Instead, we account for this distribution in terms of the impact of the relativistic particles from local (nonthermal filaments) and extended  sources with diffuse neutral gas producing both a nonthermal bremsstrahlung X-ray continuum emission, as well as diffuse 6.4 keV line emission.  The production rate of Fe K$\\alpha$ photons associated with the injection of electrons into a cloud as a function of column density is calculated.  The required energy density of low-energy cosmic rays associated with the synchrotron emitting radio filaments or extended features is estimated to be in the range between 20 and $\\sim10^3$ eV cm$^{-3}$ for Sgr C, Sgr B1, Sgr B2, and ``the 45 and -30 \\kms'' clouds.  We also generalize this idea to explain the cosmic-ray heating of molecular gas, the  interstellar cosmic ray ionization,  the  pervasive production of diffuse K${\\alpha}$ line and TeV emission from the Galactic center molecular clouds.  In particular, we suggest that Inverse Compton scattering of the sub-millimeter radiation from dust by relativistic electrons may contribute substantially to the large-scale diffuse TeV emission observed towards the central regions of the Galaxy. %We argue that %that the distribution of the 6.4 keV line emission with respect to %$^{13}$CO emission  is %inconsistent with the irradiation model of Sgr~C. ",
        "introduction": "\\label{introduction}  It is well recognized that the Galactic center region hosts several sources of energetic activity in the form of nonthermal linear filaments,  supernova remnants and colliding winds of massive stars. This region is also recognized to be the site of massive molecular clouds with pockets of on-going and past massive star formation. The most prominent of these clouds are associated with Sgr B2, Sgr C, as well as the clouds associated with radio continuum arc and Sgr A. Sgr~C or G359.43--0.09 is  one of the best examples  of a massive star-forming region which hosts prominent thermal and nonthermal radio continuum, millimeter, sub-millimeter and X-ray sources in the Galactic center region. This massive star forming region is identified by its H\\,{\\sc ii} region, a dense molecular cloud and nonthermal radio filaments. This source is also known as an X-ray source based on ASCA observations (Murakami et al. 2001a) showing the evidence of iron K-shell fluorescence (K${\\alpha}$) line emission. To probe the nature of high-energy activity in the Galactic center region, we have selected to study Sgr~C in detail, as well as star forming regions near the continuum arc ($l \\sim 0.2^{\\circ}$) and the Sgr B complex.  Sgr~C consists of a thermal radio continuum, a far-IR and a sub-millimeter source, as well as a synchrotron emitting source associated with magnetized filaments (Odenwald \\& Fazio 1984; Caswell \\& Haynes 1987; Liszt 1992; Tsuboi et al. 1991; Lis \\& Carlstrom 1994). The  Sgr~C H\\,{\\sc ii} complex is thought to be powered by a single O4  ZAMS star. The  molecular mass of Sgr C is estimated to be $\\sim5-6\\times10^5$ \\msol with column density of $N_{\\rm H_2}\\sim10^{23}$ cm$^{-2}$, based on sub-millimeter and $^{13}$CO observations (Liszt \\& Spiker 1995; Lis \\& Carlstrom 1994). Recent radio continuum survey of the Galactic center region has identified additional nonthermal filaments distributed in the immediate vicinity of the Sgr C H\\,{\\sc ii} region (Nord et al. 2004; Yusef-Zadeh, Hewitt \\& Cotton 2004). The main nonthermal radio filament (NRF)  appears to end abruptly inside the Sgr~C molecular H\\,{\\sc ii} complex M359.5--0.15 with a velocity of --65 \\kms (Liszt \\& Spiker 1995).  Similar radio recombination line velocity has also been detected from the Sgr C H\\,{\\sc ii} region (Anantharamaiah \\& Yusef-Zadeh 1989). HI absorption measurements toward the nonthermal filament and the H\\,{\\sc ii} region constrain the distance to both sources within the inner 200 pc of the Galactic center. However, Roy (2003)  suggests that the Sgr~C filament and H\\,{\\sc ii} region   are two separate objects and may not be interacting with each other (Roy 2003).  The ASCA measurements show the presence of strong fluorescent 6.4 keV line emission from neutral iron gas implying H$_2$ column density similar to that of molecular line observations (Murakami et al. 2001a). The nature of the 6.4 keV line emission has been considered to be due to the irradiation of the Sgr~C molecular cloud originated by past energetic outburst from Sgr~A* (Murakami et al. 2001a). Similar mechanism  has been postulated to explain the origin of the strong 6.4 keV line emission detected from other Galactic center molecular clouds (e.g., Sgr~B2) by either Sgr A* or from the interaction of the Sgr A East and ``the 50 \\kms cloud'' (Koyama et al. 1996;  Murakami et al. 2000, 2001b; Revnivtsev et al. 2004; Fryer et al. 2006). An alternative explanation to the X-ray reflection nebula model has  accounted for the origin of X-ray emission in terms of  the impact of low-energy cosmic ray particles with neutral gas associated with the M0.11--0.08 molecular cloud (Yusef-Zadeh, Law \\& Wardle 2002). These authors speculated that the same process could also explain the origin of the 6.4 keV line emission from other  Galactic center molecular clouds.  The structure of this paper is as follows. In $\\S$2, we re-analyze the archival {\\it Chandra} observations of the Galactic center region followed by the analysis of new millimeter line and radio continuum data using FCRAO and VLA.  In $\\S3$, we report the discovery of X-ray filamentary structures in Sgr C. The morphology of X-ray and nonthermal radio continuum features suggest that many of the K$\\alpha$ 6.4 keV line features lie near the nonthermal radio filaments at 20 and 90~cm wavelengths. We also examine the morphology of \\thco\\, and Fe K$\\alpha$ lines and estimate column densities using millimeter and X-ray techniques. We  present spectral index determination in $\\S3.1.5$   and polarization measurements of the brightest filament in Sgr C at radio wavelengths in the Appendix. The synchrotron parameters of the Sgr C filament is critical in support of our picture that the low-energy component of the synchrotron emitting particles penetrate molecular clouds. We then compare the distribution of 6.4 keV and TeV emission from prominent galactic center molecular clouds. The overall correlations of K$\\alpha$ line emission, molecular line emission from the ``45 and -30 \\kms\\ clouds'', TeV emission and nonthermal radio continuum features are described. In $\\S$4, we argue in support of enhanced cosmic rays in the central region of the Galaxy and compare the irradiation model versus the low-energy cosmic ray model to account for the origin of the 6.4 keV emission. We then describe a quantitative mode of injection of cosmic rays into molecular clouds and the subsequent heating of gas before this model is applied to the Galactic center molecular clouds. Unlike the irradiation model, the picture that is presented here can naturally explain the high temperature of molecular gas in the galactic center region. Lastly, we investigate the origin of diffuse TeV emission from Galactic center clouds in terms of the inverse Compton scattering of dust emission by the high energy component of the cosmic rays distributed throughout this unique region of the Galaxy.   %bombardment of molecular clouds by %nonthermal electrons. %Following a detailed study of Sgr C, we model the nature of 6.4 keV emission %from  Galactic center  molecular clouds that are correlated spatially with %6.4 keV line and nonthermal radio continuum emission. %Given the role of low-energy  cosmic rays in %production of  diffuse 6.4 keV emission, % app to Sgr C  and %low-energy %impact of cosmic rays to  other Galactic center molecular ",
        "conclusions": "The distributions of \\thco\\, line, X-ray, nonthermal radio continuum and 6.4 keV line emission suggest against the irradiation origin of X-ray emission from Sgr C. We believe that the origin of X-ray continuum and line emission can naturally be explained by the interaction of cosmic rays with prominent Galactic center molecular clouds such as Sgr C, the ``45 and -30 \\kms clouds'', Sgr B1 and Sgr B2. We have argued that the sources of nonthermal particles are due to a combination of the Galactic center radio filaments which are localized in the immediate vicinity of molecular clouds, as well as the diffuse nonthermal component distributed throughout the Galactic center region. This hypothesis is consistent with similar value of 6.4 keV equivalent width throughout the Sgr~C, the -30 \\& 45 \\kms clouds near the radio arc and Sgr B1 regions.  In the context of this model, we explain the high value of the equivalent width of Fe K$\\alpha$ line and unusually strong K$\\alpha$ line emission from Sgr B2 in terms of higher metallicity and the contribution of a diffuse nonthermal particles in this source. The interaction of the nonthermal radio filaments and the molecular cloud implies that these features are at the same distance from each other and that the long-standing problem of the origin of high temperature molecular gas can naturally be explained in terms of the LECR heating of the clouds. In addition, we estimate the high cosmic-ray ionization rate in the Galactic center region which could explain the origin of warm molecular gas as traced by H$_3^+$ absorption study. In addition, we provide an alternative model to explain the origin of diffuse TeV emission from molecular clouds in the central region of the Galaxy by using the high energy component of nonthermal electrons to upscatter the submillimeter emission from dust clouds in this region. The remarkable correlation of K$\\alpha$ line emission, TeV emission and molecular gas suggest that similar processes are at work in explaining both diffuse X-ray and $\\gamma$-ray emission from the Galactic center clouds. Lastly, we believe  that the interplay between the cosmic rays and  neutral gas in the central  region of the Galaxy is  significant and provides much insight in our understanding of the physical processes that are involved in both the  nuclei  and the disk of galaxies.   %In the cosmic ray interaction picture, %we speculate  the possibility  that the He-like 6.7 keV line emission line from Fe may arise %via  charge exchange as the fast moving Fe ions   transfer electrons with neutral gas. %The  association of  6.7 keV line emission with the Galactic center molecular clouds %could be used as additional support for the cosmic ray interation model discussed here."
    },
    "0608/astro-ph0608062.txt": {
        "abstract": "We build, as far as theory will permit, self consistent model  \\HII\\ regions around central clusters of aging stars. These produce strong emission line diagnostics applicable to either individual \\HII\\ regions in galaxies, or to the integrated emission line spectra of disk or starburst galaxies. The models assume that the expansion and internal pressure of individual \\HII\\ regions is driven by the net input of mechanical energy from the central cluster, be it through winds or supernova events. This eliminates the ionization parameter as a free variable, replacing it with a parameter which depends on the ratio of the cluster mass to the pressure in the surrounding interstellar medium. These models explain why \\HII\\ regions with low abundances have high excitation, and demonstrate that at least part of the warm ionized medium is the result of overlapping faint, old, large, and low pressure \\HII\\ regions. We present line ratios (at both optical and IR wavelengths) which provide reliable abundance diagnostics for both single \\HII\\ regions or for integrated galaxy spectra, and we find a number that can be used to estimate the mean age of the cluster stars exciting individual \\HII\\ regions. ",
        "introduction": "Introduction} Much of what we have learnt about the chemical evolution of the Universe, or of individual galaxies throughout cosmic time has been gleaned from the study of the spatially unresolved emission line specta of distant galaxies \\citep{Steidel96,Kobulnicky99a,Kobulnicky99b,Kewley05}. This emission line spectrum arises from the ensemble average of \\HII\\ regions within the galaxy, characterized by a wide range of physical parameters.  For example, in disk galaxies, the chemical abundance of heavy elements falls continuously from centre to edge. In like manner, both the pressure and density of the interstellar medium in the disk drops exponentially with radius. For a given input of mechanical energy by the central star cluster, this means that the outermost \\HII\\ regions expand more rapidly, and are both larger and have lower internal pressure at a given age. However, the properties of the central cluster of a given mass change systematically with chemical abundance. Massive stars with higher chemical abundances lose more mass prior to supernova explosion, have stronger stellar winds, evolve more rapidly, and have lower effective temperatures while they are on the Main Sequence.  Finally, the temporal evolution of individual \\HII\\ regions is extremely important in determining the integrated emission line spectrum. \\citet{Leitherer92} demonstrated the importance of aging of the OB stars in lowering the number of ionizing photons as the most massive stars evolve away from the Zero Age Main Sequence (ZAMS) to become supergiants. Later, these stars enjoy a brief resurgence in their ionizing photon and mechanical energy production as they become Wolf-Rayet stars.  With so many temporal and physical variables, it is perhaps not surprising that the nebular modelling community (ourselves included) have in the past been reduced to making severe, even gross, oversimplifications in attempts to model the integrated strong emission line spectra of galaxies. Indeed, frequently only a single line ratio has been used to determine the metallicity of extragalactic \\HII\\ regions or even of whole galaxies. This is the famous  \\R23 ratio; \\ratioR23.  This ratio was first proposed by \\citet{Pagel79}. The logic for its use is impeccable, since it uses the two strongest lines of the strongest coolant of \\HII\\ regions, and it should therefore be sensitive to the total oxygen abundance. By using two stages of ionization, it accounts for the emission in the bulk of the ionized nebular volume. However, the calibration of this ratio in terms of the abundance has proved to be very difficult and many different and often contradictory calibrations of \\R23 are available, including \\citet{Pagel79,Pagel80,Edmunds84,McCall85,Dopita86,Torres-Peimbert89,Skillman89,McGaugh91,Zaritsky94,Pilyugin00,Kewley02}.  The basic cause of this difficulty in calibrating the \\R23 ratio is common to many ratios formed from the flux of optical forbidden line divided by the flux of a recombination line of hydrogen. Initially, as $Z$ is increased, the forbidden line increases in flux, and the line ratio increases. However, as the abundance increases further, the cooling by forbidden lines lowers the electron temperature making it more difficult to collisionally excite the optical forbidden lines. Eventually there comes a point at which an increase in abundance is matched by a decrease in collisional excitation and the line ratio reaches its maximum. Any further increase in $Z$ then leads to a \\emph{decrease} in the relative strength of the forbidden line. \\R23, like other optical diagnostics is therefore a two-valued function of $Z$. Because infrared lines have lower thresholds for collisional excitation, they remain monatonic functions of $Z$ up to much higher metallicities. Recognizing this physics, the S$_{23}$ ratio; \\ratioS23, or even more complex ratios, have been used either in the place of or to supplement the \\R23 ratio \\citep{Dennefeld83,Kennicutt96,Oey02}.  The calibration problems of the \\R23 ratio are made worse by the scatter, at fixed  \\R23 and within each given branch, both between models and observational data, and amongst different models. The observational material is affected by uncertainty in the reddening correction applying to the \\OIILs\\ , since classical extinction laws are usually used, rather than the attenuation laws which more correctly apply to extended objects \\citep{Fischera03,Fischera05}. This point will be discussed more fully below. As far as the theoretical models are concerned, scatter in the relationship is easily generated, since the stellar atmospheres used for single stars can show very wide differences in the number of ionizing photons produced per unit frequency \\citep{Morisset04}.  Moving beyond the use of single line ratios, a more sophisticated analysis is made possible by the understanding that the strong line emission spectrum of an individual \\HII\\ region is controlled by three physical parameters \\citep{Dopita86}. These parameters are, the effective temperature of the exciting stars, the chemical abundance set (or ``metallicity\", $Z$) and the ionization parameter, defined either as the ratio of the mean ionizing photon flux to the mean atom density, $q$ or in its dimensionless form as the ratio of mean photon density to mean atom density ${\\cal{U}}=q/c$. Thus, for a given stellar input spectrum a ${\\cal{U}} : Z$ grid of models should provide the ideal means of interpreting the observational data \\citep{Kewley02}.  This analysis leaves open the issue of the age distribution of the exciting stars. The metallicity and the age of the exciting stars determines the energy distribution of the EUV photons. In the modelling hitherto, this problem is usually treated in two limits;  the instantaneous burst approximation (for which zero age is often, but not universally assumed) or the continuous star formation approximation in which stars are assumed to be born continually within the ionized region (see \\citet{Kewley01a}, and references therein). All this modeling  effectively assumes a single \\HII\\ region of a given metallicity and ionization parameter. Whilst such an approach is acceptable in attempts to understand the emission line spectrum of individual  \\HII\\ regions, it is certainly inadequate in attempts to interpret the integrated spectrum of a whole galaxy, or even of a complex of many \\HII\\ regions.  The advent of the  \\emph{Starburst 99 v.5} \\citep{Leitherer99} stellar spectral synthesis code finally allowed the construction of fully self-consistent models for \\HII\\ regions. For a given stellar initial mass function (IMF), this code delivers not only the pan-spectral distribution of photons, including ionizing photons as a function of age, but also provides tables of the time-dependent mechanical energy input of the stars including OB-star stellar winds, red giant winds, Wolf-Rayet winds and the energy input from supernova explosions.  In the first paper of this series (\\citet{Dopita05}, hereinafter referred to as Paper I), we computed the (one dimensional) evolution of \\HII\\ regions around clusters of a single mass. We demonstrated how the size~:~age relationship is essentially determined by the density (or pressure) in the surrounding interstellar medium (ISM). This in turn determines the dust `temperature' described by the wavelength of the peak of the far-IR dust re-emission feature.  In the second paper of this series, Dopita et al. (2006), in press, hereinafter referred to as Paper II), we investigated the role of metallicity and of the cluster mass function in controlling the excitation of the populations of \\HII\\ regions in galaxies, be they disk or starburst.  We demonstrated that the ionization parameter of \\HII\\ regions, previously treated as a free variable, is determined at any time by the instantaneous ratio of the ionizing photon flux to the mechanical energy flux of the central stars. This is because the HII region is evolving as a mass-loss bubble pressurized by the combined ram pressure of the stellar wind and supernova explosions, so HII region density and HII region radius are closely coupled. Although this conclusion was obtained using a simple 1-D spherical evolution model for the \\HII\\ region, it should remain valid in more complex geometries, since wherever un-ionized or molecular inclusions exist, these will be pressure confined by the hot shocked stellar wind gas, and this gas pressure is coupled with the overall size of the \\HII\\ region, which itself determines the dilution of the radiation field of the central stars.  Because the effective temperature of the exciting stars increases as metallicity is lowered, but the mechanical energy flux in the stellar winds decreases towards to lower metallicity, low abundance HII regions are characterized by higher ionization parameters, a fact that had been frequently noted by observers, but which had not previously been explained in a satisfactory manner.  Using these insights, in this third paper of the series,  we proceed to the next stage of sophistication in modeling -- the construction of time-dependent photoionization models of individual \\HII\\ regions, and the construction of ensemble averages of aging \\HII\\ regions to provide galaxy-wide averaged spectra at a given metallicity.  Such models applied to individual \\HII\\ regions make it possible to estimate the ages  of the exciting stars from the positions of the observations on theoretical \\HII\\ region isochrones. Previously, this could not be done, since the ionization parameter and the effective temperature of the cluster  (determined by the ageing of its stars) can both produce similar effects on the emitted spectrum. This degeneracy is raised thanks to the self-consistent geometry that we have developed, which includes a dynamical evolution of the \\HII\\ region consistent with the properties of the stellar wind generated by the central cluster.  In this paper, we apply these new techniques to produce line ratio diagnostics which enable the independent determination of both the stellar ages and metallicities of the exciting clusters for individual \\HII\\ regions. These should prove useful in the analysis of abundance gradients in resolved disk galaxies. More importantly, we develop line ratio diagnostics for the determination of abundances in ensembles of \\HII\\ regions in starburst and normal galaxies. This should greatly assist in analyses of the strong emission lines of the more distant and unresolved galaxies, observed either at optical wavelengths or in the IR. ",
        "conclusions": "\\label{discussion}  It should now be clear that these models represent a significant advance in our understanding of the strong line spectra of \\HII\\ regions. In particular, the models explicitly account for the dynamical evolution of the \\HII\\ region controlled through the mechanical energy input by both the stellar winds and supernova explosions. Two factors largely control the excitation of the young \\HII\\ regions, both of these dependent on the chemical abundance. These are: \\begin{itemize} \\item{The effective temperature of the exciting stars, which increases with decreasing metallicity, and} \\item{The control of the ionization parameter by the stellar winds in the expanding \\HII\\ regions. This was shown in Paper II to increase with decreasing metallicity} \\end{itemize} Since both of these effects operate to increase the excitation of low abundance \\HII\\ regions, we therefore find an abundance dependent excitation, which much improves the fit of the theory compared to models in which the ionization parameter is treated as a free variable. In our models the ionization parameter is replaced by the $R$ parameter, the ratio of the cluster mass in solar masses to the pressure in the ISM; $P/k$~(cm$^{-3}$K). This parameter is used because all models with a given $R$ have a unique relationship between their instantaneous ionization parameter and time. This allows us to construct evolutionary tracks for ensembles of \\HII\\ regions.  With increasing age, the ionization parameter falls almost monotonically, and the \\HII\\ region fades after reaching its maximum brightness in \\Ha\\ after 1-2~Myr. The presence of Wolf-Rayet stars after about 3.5~Myr causes the effective temperature of the cluster to briefly increase. This increases the excitation for a short time, causing the 3~Myr and the 4~Myr isochrones to cross over each other on diagnostic plots for the cases of high metallicity ($Z/Z_{\\odot} > 0.5$, approximately).  In this paper, we have identified a number of useful diagnostics for deriving the abundance and either the age, in the case of isolated \\HII\\ regions, or the $R$ parameter, in the case of integrated spectra of galaxies. Our diagnostics for the integrated spectra are the first which explicitly take into account the fact that when we observe a whole galaxy, we are observing an ensemble average of \\HII\\ regions of all ages, sizes and central cluster masses. The simplification we have been able to make by using the $R$ parameter has rendered this problem tractable to a more simple analysis other than treating each \\HII\\ region separately, and adding the whole lot together at the end.  In comparing the models to the observational data we have identified some potential problems with the way that observations of \\HII\\ regions have been made and analyzed. In particular: \\begin{itemize} \\item{Long-slit observations do not integrate over the whole of an \\HII\\ region, particularly in the case of observations of nearby, low metallicity dwarf irregular galaxies. By picking out the brightest regions of such \\HII\\ regions, we bias the data towards the high-excitation regions and the spectrum is not representative of the whole \\HII\\ region \\citep{Luridana01,Stazinska03} and } \\item{The use of a reddening correction using an extinction curve, rather than an attenuation curve (which is appropriate for spatially extended objects) will cause the line intensities of UV lines to be systematically overestimated, biasing the strong line diagnostics which use the \\OIILs\\ \\AA\\ lines.} \\end{itemize}  This second effect would reveal itself in the spectro-photometry as a failure to reproduce the correct intensities of the higher members of the Balmer Series, such as H$\\delta$ or H$\\epsilon$, and so should be quite easy to check for in observational data sets. Such an effect would also influence the measurement of the strength of the \\OIII\\ $\\lambda 4363$\\AA\\ line, again causing it to be overestimated. This could be part of the reason for the long-standing discrepancy between the abundances derived from the strong lines and those obtained from measurements of the electron temperature using ratios such as \\OIII\\ $\\lambda 4363$\\AA/\\OIIIL\\ , see \\citet{Bresolin04} for a recent discussion. We expect that the major part of the discrepancy will turn out to be caused by temperature fluctuations. However, this could be easily corrected for by assuming the Case B recombination line ratio between \\Hb\\ and H$\\gamma$, and measuring the forbidden line strengths relative to the nearest Balmer lines in the spectrum.  There is clearly a need for integral field spectrophotometry on extragalactic \\HII\\ regions. This would allow efficient construction of integrated spectra and assist in the removal of stellar flux from the spectrum. Such observations would not only help to remove the remaining discrepancies between theory and observation but also give insight into the use of theoretical integral spectra in the analysis of the chemical evolution history of the early Universe.  Our models have their deficiencies. In general, all the low ionization species are predicted weaker than observed. These include the \\OIL, the \\SIILs, and the \\OIILs, although these lines are probably affected by the reddening corrections made to the observations.  What could enhance the \\SIILs/\\Ha\\ ratio or the \\OIL/\\Ha\\ ratio? First, both the the \\ion{O}{1} and the  \\ion{S}{2} ion are very much more sensitive than the  \\ion{N}{2} ion to the diffuse radiation field in \\HII\\ regions. They are therefore much enhanced in the vicinity of ionization fronts, and in shadow regions behind elephant trunk features, which are obscured from direct stellar radiation and illuminated only by the diffuse nebular radiation field. Second, both are much enhanced in shock regions. Indeed high  \\SII/H$\\alpha$ ratios are used as a means of identifying supernova remnants either embedded in \\HII\\ regions or in the ISM at large, and \\OI/H$\\alpha$ ratios reach 0.5 or higher in shock-excited regions. Neither of these effects are taken into account in the models.  In addition, the \\ion{O}{1} ion is locked by charge-exchange reactions to the concentration of \\HI. Furthermore, the collision strength of \\OIL\\ increases with electron temperature. To produce a high  \\OIL/\\Ha\\ ratio in a purely photoionized region requires the presence of a hard radiation field. This could originate either as a thermal soft X-ray continuum from the shocked bubble of stellar wind material, or it could simply be an indication of a hotter sub-component in the global stellar radiation field.  If the radiation field produced by \\emph{Starburst 99 v5.} is too soft, this could be rectified in at least three ways, First, by the use of a flatter IMF in the upper mass ranges ($M/M_{odot} > 10)$. This could be possible because observationally we have very few constraints other than by use of the excitation of \\HII\\ regions. Second, by adding another stellar component such as mass-exchange massive binaries to the stellar synthesis.  A third - and quite likely possibility is that the EUV spectrum of the massive stars used in \\emph{Starburst 99 v5.} may be in error. With previous versions of \\emph{Starburst 99} based on the \\emph{CoStar} models \\citep{Schaerer97}, the EUV spectrum was, rather, too hard. However, \\citet{Morisset04} has demonstrated that the theoretical number of ionizing photons emitted by the central star and the shape of the ionizing spectrum is highly dependent upon the atmospheric model used. At one extreme are the \\emph{CoStar} models, at the the other extreme are found the `classical' \\emph{Kurucz} \\citep{Kurucz91, Kurucz94} plane-parallel LTE line blanketed models used in \\emph{Starburts99 v5.}. The Kurucz models display the softest EUV spectra and the lowest photoionising flux, and are restricted to $\\log g = 3.0$ for the models with higher effective temperature, $T_{\\rm eff}$.  In between these extremes are found three sets of models with fairly similar photoinising flux predictions \\citep{Morisset04}, the \\emph{TLUSTY} models of \\cite{Lanz03a, Lanz03b}, the \\emph{WM-Basic} models of \\cite{Pauldrach01} and the \\emph{CMFGEN} models \\cite{Hillier98}. Of these, the \\emph{TLUSTY} models are plane-parallel hydrostatic, while the other two are spherical dynamic atmospheres. All of this offers plenty of potential for error in the estimation of the the hardness of the EUV spectrum for a given cluster mass and metallicity.  Given that the \\OIIIL/\\Hb\\ ratio of the models also appears to be on the low side, we believe that an increase in the stellar effective temperature, either by a flattening of the IMF or by some other means offers the best chance to more properly match the theory with the observations. The models are too weak in their forbidden line strengths overall. The \\citet{Stoy33} method \\citep{Pottasch83} of deriving the effective temperature of the exciting stars relies upon measuring the total flux cooling lines (which is a measure of the mean energy per photoionization) to a recombination line (which counts the number of photoionizations). Thus, if we want to increase the absolute strengths of the forbidden lines, we have to increase the stellar effective temperature. This increase would also strongly increase the  \\OIL/\\Ha\\ ratio by both increasing the extent of the partially-ionized zonr and by increasing the electron temperature within it.  A niggling uncertainty with the modeling is the mis-match between the \\emph{Starburst 99 v5} abundance set and the abundance set given in table \\ref{table1}. The \\emph{Starburst 99 v5} has to rely on the older stellar evolution models which use the old value of the solar abundances. In an attempt to quantify this effect, we ran a test model with a $0.4Z_{\\odot}$ spectral synthesis cluster model from \\emph{Starburst 99 v5} embedded in a nebula with $1.0Z_{\\odot}$ (using the abundances given in table \\ref{table1}). This made surprisingly little difference, changing critical line ratios by 0.1~dex or less, except for the \\OIL/\\Ha\\ ratio, which increased by 34\\%. We therefore conclude that the results presented in this paper are reasonably secure against any future adjustment of stellar atmospheric abundances.  In conclusion, the major advances in modelling made in this paper are: \\begin{itemize} \\item{For individual HII regions, it becomes possible to estimate the ages  of the exciting stars from the positions of the observations on HII region isochrones. Previously, this was not possible, since the ionization parameter and the effective temperature of the cluster  (determined by the ageing of its stars) produce similar effects on the  emitted spectrum. The self-consistent treatment  we have developed which  includes the dynamical evolution of the HII region consistent with the properties of the stellar wind generated by the central cluster, and which recognises that the stellar wind determines the pressure in the HII region allows us to eliminate the use of arbitrary geometries.} \\item{ In addition, this work gives a methodology to construct much more ``realistic\" integrated spectra of HII galaxies, something that has not been possible before in the approximations which assume that a whole galaxy can be represented by a single spherical \\HII\\ region with a chosen abundance and ionization parameter.} \\end{itemize}"
    },
    "0608/astro-ph0608426_arXiv.txt": {
        "abstract": "I study a problem, whether the asymmetry of a $^{56}$Ni ejecta that results in the asymmetry of the H$\\alpha$ emission line at the nebular epoch of the type IIP supernova SN~2004dj is able also to account for the recently detected polarization of the supernova radiation. I developed a model of the H$\\alpha$ profile and luminosity with a nonthermal ionization and excitation taken into account adopting an asymmetric bipolar $^{56}$Ni distribution. On the background of the recovered distribution of the electron density I calculated the polarized radiation transfer. It is demonstrated that the observed polarization is reproduced at the nebular epoch around day 140 for the same parameters of the envelope and $^{56}$Ni distribution for which the luminosity and profile of H$\\alpha$ are explained. Yet the model polarization decreases slower compared to observations. The origin of an additional component responsible for the early polarization on day 107 is discussed. ",
        "introduction": "A problem of an explosion mechanism for type II supernovae remains unresolved, although this phenomenon is undoubtedly related with the gravitational collapse of an iron core. In this regard of great importance are the observational manifestations of the explosion that might hint the ways of the problem solution and to play a role of observational tests of models. One of the specific properties of the explosion mechanism is the asymmetry (Ardeljan et al. 2005; Blinnikov et al. 1990; Herant et al. 1992; Scheck et al. 2004; Burrows et al. 2005). It can manifest itself both in the asymmetry of the emission line profiles and in the polarization of radiation. Asymmetry effects were observed for the first time in the anomalous type IIP supernova SN~1987A (for H$\\alpha$ asymmetry see Phillips and Williams [1991] and for the polarization see Jeffery [1991]). Both effects have been interpreted in a model of asymmetric $^{56}$Ni ejecta embedded in the symmetric envelope (Chugai 1991, 1992). The line asymmetry was observed also in the normal type IIP supernova SN~1999em (Elmhamdi et al. 2003). Interestingly, in both supernovae the asymmetry was characterized by the redshift which indicated that the $^{56}$Ni was ejected predominantly toward the far hemisphere.  A standard type IIP supernova SN~2004dj (Nakano et al. 2004) generated a great interest because of its proximity (D=3.13 Mpc). However, more important is that its nebular spectra revealed a strong line asymmetry, particularly of the H$\\alpha$ emission. The asymmetry manifested itself as a pronounced blueshift of emission line maximuma by about 1500 km s$^{-1}$ (Chugai et al. 2005). Modeling the H$\\alpha$ in that paper led us to conclude that $^{56}$Ni is distributed as bipolar asymmetric ejecta with a more masive jet residing in the near hemisphere.  The recovery of the $^{56}$Ni distribution upon the basis of the profile modeling is, strickly speaking, ill-posed problem, so doubts may arise whether this procedure is unambiguous. An independent confirmation is therefore needed. Recently Leonard et al. (2006) reported results of the polarimetric observations of SN~2004dj in which the intrinsic variable polarization in the band of $6800-8200$ \\AA\\ was detected with maximum value of 0.56\\% at the end of the light curve plateau. According to authors suggestion the polarization in SN~2004dj arises because of the Thomson scattering, while the evolution of the polarization is interpreted as a result of the decrease of the Thomson optical depth of the supernovae with spherical envelope and asymmetric core. It is shown, simultaneously, that the change of position angle of the polarization contradicts to the model of axialy symmetric bipolar $^{56}$Ni ejecta.  A question arises whether the asymmetry of the $^{56}$Ni distribution responsible for the H$\\alpha$ profile asymmetry of the SN~2004dj is able also to account for the polarization at least at the nebular epoch? The present paper is concentrated on the answer to this important question. In second section I describe the model used to compute the electron density distribution caused by the asymmetric $^{56}$Ni distribution. The third section presents results of the modelling of H$\\alpha$ line and of the polarization produced by the Thomson scattering of continuum photons in the asymmetric electron distribution.  I adopt 2004 June 28 to be the explosion date in contrast to June 13 adopted in the previous paper. The revision is dictated by two arguments. First, plateau length of SN~1999gi, which was used as a template for SN~2004dj, is maximum ($\\approx125$ d) among type IIP supernovae, while the average duration is $\\approx110$ d. Second, from the spectral evolution of SN~2004dj Patat et al. (2004) estimate the explosion date to be about July 14. Adopting the plateau duration to be 110 d we come to a compromise explosion date of June 28, which is two days earlier than the explosion date estimated by Vinko et al. (2006) from the analysis of the photosphere evolution. ",
        "conclusions": "The primary goal of the paper was to provide an answer to the question whether the asymmetry of the $^{56}$Ni distribution responsible for the asymmetry of the H$\\alpha$ emission line at the nebular epoch of SN 2004dj is able to account for the observed polarization as well. I constructed a model of the H$\\alpha$ profile and luminosity with the nonthermal excitation and ionization of hydrogen for the asymmetric $^{56}$Ni distribution. I then calculated the transfer of a polarized radiation using Monte Carlo technique. The modelling shows that the model of the bipolar ejecta, which describes the H$\\alpha$ line, predicts that the model polarization is consistent with the observed polarization in SN~2004dj at the early nebular epoch. At the late stage $t\\sim 200$ d  the theoretical polarization is somewhat larger compared with observations. This disparity is possibly related with a more complicated structure of the matter distribution, including $^{56}$Ni, in the inner zone of the envelope.  The photospheric component of the polarization observed on day 107 is not described in principle by the model of the axisymmtric $^{56}$Ni ejecta. Leonard et al. (2006) attribute the early polarization to the deviation of the bipolar $^{56}$Ni ejecta from  the axial symmetry. According to their conjecture the recombination wave propagating toward the center of the supernova envelope uncovers in sequence layers with different position angles of $^{56}$Ni clumps. This might explain the observed change of polarization position angle between days 107 and 144. In this model the photosphere plays a role of a homogeneous screen, while asymmetry arise from the enhanced ionization originated from the asymmetric $^{56}$Ni distribution. The scenario proposed by Leonard et al. (2006) can be dubbed \"the model of non-axisymmetric $^{56}$Ni ejecta\".  An alternative conjecture that does not require abandoning the axial symmetry of $^{56}$Ni ejecta is conceivable either. Note, the photospheric component of the polarization emerged at the phase when the photosphere almost reached the center of the supernova envelope. On the other hand, the distribution of the chemical abundance in the inner layers of supernova should be essentially inhomogeneous due to an incomplete mixing of the hydrogen and helium enevelopes during the shock wave propagation (M\\\"{u}ller et al. 1991). The photosphere that forms in the chemically inhomogeneous material unavoidably should acquire large scale brightness variations that might result in a significant polarization of the photospheric radiation. Such a scenario one may dub a model of \"spotty photosphere\" in contrast to the model of non-axysymmetric $^{56}$Ni ejecta.  In both conjectures the early polarization is related to the final photospheric phase and in both mechanisms a \"flash\" of polarization should generally emerge at the end of the light curve plateau. The difference is that in the model of the spotty photosphere the polarization flash should be present even in SN~IIP with the almost spherical $^{56}$Ni distribution, while the model based on the $^{56}$Ni asymmetry predicts in this case essentially weak/no polarization. This signature can be used to discriminate between alternative conjectures about the origin of the photospheric component of polarization in SN~2004dj and other SN~IIP with the detected polarization flash.   \\bigskip  I am grateful to Douglas Leonard for the kind permission to use the spectrum of SN~2004dj. The work is partially supported by RFBR grant 04-02-17255.  \\newpage"
    },
    "0608/astro-ph0608610_arXiv.txt": {
        "abstract": "Based on results from cold dark matter N-body simulations we develop a dynamical model for the evolution of subhaloes within group sized host haloes. Only subhaloes more massive than $5\\times 10^8\\Msol$ are considered, because they are massive enough to possibly host luminous galaxies. On their orbits within a growing host potential the subhaloes are subject to tidal stripping and dynamical friction. At the present time ($z=0$) all model hosts have equal mass ($M\\vir = 3.9 \\times 10^{13} \\Msol$) but different concentrations associated with different formation times. We investigate the variation of subhaloe (or satellite galaxy) velocity dispersion with host concentration and/or formation time. In agreement with the Jeans equation the velocity dispersion of subhaloes increases with the host concentration.  Between concentrations $\\sim5$ and $\\sim20$ the subhaloe velocity dispersions increase by factor of $\\sim1.25$. By applying a simplified tidal disruption criterion, i.e. rejection of all subhaloes with a tidal truncation radius below $3\\kpc$ at $z=0$, the central velocity dispersion of the 'surviving' subhaloe sample increases substantially for all concentrations. The enhanced central velocity dispersions in the surviving subhaloe samples are caused by a lack slow tangential motions. Additionally, we present a fitting formula for the anisotropy parameter which does not depend on concentration if the group-centric distances are scaled by $r_s$, the characteristic radius of the NFW-profile. Since the expected loss of subhaloes and galaxies due to tidal disruption increases the velocity dispersion of surviving galaxies, the observed galaxy velocity dispersion can substantially overestimate the virial mass. ",
        "introduction": "Clusters of galaxies have been extensively studied in all available wave bands. They belong to the most massive and eye-catching structures in the Universe and are thought to closely resemble the cosmic composite of 85 per cent dark matter and 15 per cent baryonic matter \\citep{Spergel06}. According to the hierarchical model for structure assembly, these systems preferentially have formed late. For dark matter haloes it has been shown that the formation redshift is correlated with concentration (e.g.~\\citealt{Bullock01b,Wechsler02}) therefore clusters present a class of low concentration objects. In contrast, groups of galaxies span a whole range of formation times and concentrations. Some resemble the properties of clusters \\citep{Zabludoff98, Mulchaey00, Zabludoff00}, having low concentrations, others show an opposite behaviour.  Also the properties of the galaxy populations vary strongly amongst individual groups. Some groups have luminosity functions similar to those observed in clusters but 'fossil groups' display a central giant elliptical galaxy with only few intermediate-luminosity galaxies (e.g. \\citealt{Ponman94, Jones2003}). The optical properties of the central bright galaxies are consistent with an origin as the merged combination of a more typical, less over-luminous, galaxies. The variations amongst the satellite populations are thought to originate from varying formation times. Using N-body experiments \\cite{DOnghia05} demonstrated that the earliest formed group-sized dark matter haloes reveal structural similarity to observed fossil groups (see also \\citealt{SommerLarsen06}). The early formation times of fossil groups also impacts satellite galaxies that do not merge with the central giant. Based on the observed depletion of intermediate-luminosity galaxies in the fossil group RX J1552.2+2013, \\cite{MendesdeOliveira06} suggests that the members of this system may be affected by tidal disruption. \\cite{DaRocha05} find substantial amount of intra-group light in two of three compact groups which may, at least partly, come from tidally dissolved satellite galaxies. These observations indicate a substantial variation of the satellite population with time either by merging or by tidal dissolution.  In addition to the concentration-age and the luminosity-age relations, we describe here a concentration-velocity relation. According to the Jeans equation, galaxies orbiting in more concentrated host haloes have higher velocity dispersions. This effect is most pronounced in the central regions of the host haloe. As discussed above, early formed groups are more concentrated, leading to the conclusion that the central satellite galaxies in fossil groups move faster than central galaxies in less evolved groups of comparable size. Additionally, if tidal disruption reduces the number of satellite galaxies, the velocity dispersion of the remaining galaxies is expected to increase even higher.  The last statement is supported by the following consideration. Analysing N-body simulations, \\cite{Diemand04a} find an enhanced velocity dispersion of subhaloes compared to the dispersion of the diffuse dark matter component. Subsequently, \\cite{Faltenbacher06} argued that the loss of a slow moving, earlier accreted, subset of the subhaloe population is responsible for this bias. In analogy, if satellite galaxies are prone to tidal dissolution, then early accreted, slow moving, galaxies are preferentially dissolved. The lack of low velocity galaxies causes an enhanced velocity dispersion of the remaining galaxies. An observational hint for such a scenario may be the recent observation of the fossil system RX J1416.4+2315 by \\citep{Khosroshahi06}. Based on 18 member galaxies they find a velocity dispersion that is nearly twice as high as expected from the X-ray analysis of this system.  These findings motivated us to study the relation between the host concentration and dynamics of the subhaloe populations. The paper is organised as follows. \\S~\\ref{sec:model} introduces the dynamical model used to trace the evolution of substructure in host haloes of different mass accretion histories. \\S~\\ref{sec:point} deals with substructure as point-like particles, a reference for the subsequent more complex investigation. \\S~\\ref{sec:subhaloe} presents dynamical results when substructures are treated as extended units prone to tidal stripping and dynamical friction. Finally, \\S~\\ref{sec:summary} concludes with a summary. ",
        "conclusions": "\\label{sec:summary} Using a dynamical model for the evolution of subhaloes within a growing host potential in galaxy groups, we have investigated the relation between host concentration and properties of subhaloe population at $z=0$. The concentration is correlated with the group formation redshift. Groups with equal virial masses at $z=0$ but different formation times show significant differences in their subhaloe properties. As a base point we have computed the dynamical evolution of point mass samples and their dependence on the host concentration. The model then has been adapted to cope with the spatial extent of subhaloes and the impact of dynamical friction and tidal disruption. We only focus on the dynamical evolution of subhaloes that are massive enough to possibly host galaxies. Every substructure analysis is subject to detection limits whether it is the sensitivity limit in connection with observations or the resolution limit in numerical simulations. To study the effects of detection limits, we create a {\\it reduced} subhaloe sample comprising only those subhaloes with current radii larger than a rejection radius of $r\\rej = 3\\kpc$. The reduced sample is a sub-sample of the {\\it complete} sample, comprising all subhaloes independent of the actual truncation radius. The rejection radius can be considered as a parameter of the model and has been calibrated by observations as well as numerical simulations (see \\citealt{Faltenbacher05}). Originally the model is set up to trace the evolution of the dark matter subhaloes in a group environment, however if the subhaloes host luminous galaxies, a tidal truncation below $3\\kpc$ most probably would affect the stellar component of the galaxy as well.  Our main findings are: (1) The velocity dispersion of subhaloes or satellite galaxies in groups depends on the concentration of the underlying dark matter distribution of the host haloe. Equal mass hosts with higher concentrations have enhanced  velocity dispersions of the subhaloes, in particular at the centre. Since higher host concentrations are caused by earlier formation times, this relation implies that subhaloe populations with early formed hosts exhibit higher velocity dispersions than those residing in more recently formed hosts of equal mass. (2) We propose a fitting formula for the radial dependence of the anisotropy parameter $\\beta$, which is independent of concentration if the group-centric distances when scaled by the characteristic radius of the NFW-profile $r_s$. (3) The increase of the central velocity dispersion with concentration is amplified in the reduced sample. This is caused by an enhancement of the tangential velocity dispersion and a lack of slow radial motions compared to the complete sample. (4) The dynamical evolution of subhaloes that experience tidal stripping and dynamical friction does not alter the slope of the mass function, in agreement with the results from N-body simulations (see e.g. \\citealt{Reed05}), but dynamical evolution reduces the amplitude by a factor of $\\sim2$. In the complete sample the amplitude of the mass function is only weakly  dependent upon the host concentration. However, the reduced sample shows a stronger dependence on host concentration with amplitudes decreasing with increasing host concentrations. For a host concentration of 20 the amplitude of the initial mass function decreases by a factor of $\\sim5$, however this value depends strongly on the choice of the rejection radius.  The present investigation reveals the difficulties associated with mass estimates derived from the velocity dispersion of satellites in groups. The concentration of the host haloe and the effects of tidal reduction of the satellite galaxies in groups may alter the central velocity dispersions. The appearance of intra-group stars is strong evidence for the reduction of the satellite luminosities. According to \\cite{Osmond04} the velocity dispersion provides a very unreliable measure of system mass. The concentration and its influence on tidal reduction add more concerns on the accuracy of such mass estimates.  Combined observations of X-ray temperatures and velocity dispersions of satellite galaxies have recently become available for groups. As a application of the present analysis these observations my be used to infer the impact of tidal forces onto the central satellite population. If tidal forces in fossil groups efficiently reduce the number of detectable satellites, we expect the central velocity dispersion to be high compared to the X-ray temperature. Some support for this scenario has been found by \\cite{Khosroshahi06}.  A comparison of the X-ray temperatures of the hot intra-group gas with the velocity dispersions found for the satellite galaxies can be used to infer the degree of tidal disruption. A flattening of the central number density profiles along with an increasing velocity dispersion are the signatures of tidal reduction of the central subhaloe or satellite population.  \\cite{SommerLarsen06} considers radial velocity dispersions in fossil systems at a fixed radius $r\\gtrsim30\\kpc$. We have presented a uniform fitting formula for the anisotropy parameter provided the group-centric distances are scaled by $r_s$, the characteristic radius of the NFW-profile of the host. Fossil groups are assumed to be early formed systems (see e.g.~\\citealt{DOnghia05}) with high concentrations. Our fitting formula implies that at a fixed radius more concentrated hosts generate more radially anisotropic velocity dispersions. This result has to be taken into account if the velocity dispersions of different groups are compared at physically similar radii."
    },
    "0608/astro-ph0608268_arXiv.txt": {
        "abstract": "We present predictions for the cosmic metal budget in various phases of baryons from redshift $z=6\\rightarrow 0$, taken from a cosmological hydrodynamic simulation that includes a well-constrained model for enriched galactic outflows.  We find that substantial amounts of metals are found in every baryonic phase at all epochs, with diffuse intergalactic gas dominating the metal budget at early epochs and stars and halo gas dominating at recent epochs.  We provide a full accounting of metals in the context of the missing metals problem at $z\\approx 2.5$, showing that $\\sim40$\\% of the metals are in galaxies, and the remainder is divided between diffuse IGM gas and shocked gas in halos and filamentary structures.  Comparisons with available observations of metallicity and metal mass fraction evolution show broad agreement. We predict stars have a mean metallicity of one-tenth solar already at $z=6$, which increases slowly to one-half solar today, while stars just forming today have typically solar metallicity.  Our \\ion{H}{i} column density-weighted mean metallicity (comparable to Damped \\lya\\ system metallicities) slowly increases from one-tenth to one-third solar from $z=6\\rightarrow 1$, then falls to one-quarter solar at $z=0$. The global mean metallicity of the universe tracks $\\sim 50$\\% higher than that of the diffuse phase down to $z\\sim 1$, and by $z=0$ it has a value around one-tenth solar.  Metals move towards higher densities and temperatures with time, peaking around the mean cosmic density at $z=2$ and an overdensity of 100 at $z=0$.  We study how carbon and oxygen ions trace the path of metals in phase space, and show that \\ion{O}{iii}-\\ion{O}{vii} lines provide the most practical option for constraining intergalactic medium metals at $z\\la 2$. ",
        "introduction": "Heavy elements are produced exclusively in stars, and then distributed throughout the Universe through supernovae and stellar winds.  As such, they provide a unique tracer of the star formation and feedback processes that are central to understanding how galaxies form and evolve. Metals are present at the earliest observable epochs: The farthest known quasar at $z=6.4$ shows substantial enrichment despite being less than 1~Gyr old~\\citep{bar03}, and metals are seen in the intergalactic medium (IGM) at $z\\ga 6$~\\citep{bec06,rya06,sim06}.  At lower redshifts, the galaxy mass-metallicity relation has been quantified out to $\\sim 2$~\\citep{tre04,sav05,erb06}, intracluster gas metallicities have been constrained out to $z\\sim 1$~\\citep{ros04} and the IGM is observed in metal lines nearly continuously from $z\\sim 6\\rightarrow 0$.  \\citet{pet99} first noted the interesting fact that the metals seen in Lyman break galaxies at $z\\sim 2-3$ fall well short of the amount expected to be produced by the stars in those galaxies.  This is known as the {\\it missing metals problem}.  It implies that most metals must either reside outside of galaxies, or be in some non-observable form in or around galaxies.  Attempts to quantify metals in the IGM \\citep[][hereafter FSB05]{fer05} or other galaxy populations such as sub-millimeter galaxies~\\citep{bou06a} or damped Lyman alpha (DLA) systems~\\citep{pro06} have not clearly been able to account for the shortfall, though uncertainties remain substantial~\\citep[hereafter BLP06]{bou06b}.  Virtually every galaxy class and environment examined seems to contain a small but non-trivial amount of metals, so it appears that metal pollution has been curiously democratic.  As the distribution of metals in various environments and at various epochs becomes better measured, this provides important constraints into galactic feedback processes.  A theoretical framework for understanding these wide ranging observations is provided by simulations of galaxy formation that include metal production and distribution mechanisms.  The pioneering work of \\citet{cen99b} using cosmological hydrodynamic simulations highlighted some basic trends arising in such models, such as the early enrichment of galaxies and the natural establishment of a metallicity gradient with cosmic overdensity.  Using more advanced, higher-resolution cosmological simulations, \\citet{tha02} studied global metal enrichment from dwarf galaxies down to $z=4$, and showed that it is possible to enrich the IGM at early times to a value broadly consistent with data while producing dwarf galaxies with enrichment levels similar to local dwarfs.  An alternative approach was taken by \\citet{cal04}, who used detailed chemical evolution models for various galaxy types together with observations of luminosity functions in those types to construct the evolution of metals in galaxies, intracluster gas, and IGM gas. While their models make detailed predictions for individual elements in various galaxy classes, they do not embed the predictions within an ab initio hierarchical structure formation framework, relying instead on passively evolving observed luminosity functions back in time.  Still, this provides some interesting predictions that are not dissimilar to those presented here.  In \\citet{cal06} they extend their approach to study IGM enrichment, and suggest that winds with velocities exceeding 1000~km/s are required to pollute the IGM to observed levels.  In \\citet[][hereafter OD06]{opp06} we took such studies to the next level of sophistication, using cosmological simulations that incorporate enriched kinetic feedback in a full hierarchical structure formation setting.  Our models, employing outflow scaling relations taken from local starburst data, showed remarkable success at reproducing observations of \\ion{C}{iv} enrichment in the IGM from $z\\sim 6\\rightarrow 2$. Matching these data generally required highly mass loaded but low velocity winds from early galaxies, so as to eject a substantial metal mass into the IGM without overheating it.  Such scalings naturally arise in momentum-driven wind scenarios~\\citep{mur05} that are also favored by local outflow observations~\\citep{mar05,rup05}, providing an intriguing connection between outflows at all epochs.  In \\citet{dav06} we show that significant early feedback is also necessary in order to suppress reionization-epoch star formation.  In Finlator, Dav\\'e \\& Oppenheimer (in preparation) we show that the resulting galaxy mass-metallicity relation from our favored outflow models agrees well with observations at $z\\sim 2$.  These successes suggest that we now have a plausible structure formation framework with which to study the global evolution of metals in the Universe.  In this paper we use cosmological hydrodynamic simulations incorporating our most successful feedback model from OD06 (described in \\S\\ref{sec:sims}) to study the evolution of metals in various baryonic phases in the Universe from $z=6\\rightarrow 0$ (\\S\\ref{sec:metphase}) and provide quantitative answers for the missing metals problem (\\S\\ref{sec:missing}).  We go on to compare the metallicities in those phases with observations (\\S\\ref{sec:metallicity}), showing broad agreement.  We make predictions for how metals evolve in cosmic phase space (\\S\\ref{sec:phasespace}), and study how different oxygen and carbon ions may be used to probe various regimes of phase space (\\S\\ref{sec:tracers}).  We present our conclusions in \\S\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  Using cosmological hydrodynamic simulations of galaxy formation including galactic outflows, we have studied the distribution and evolution of metals in various baryonic phases in the Universe.  Our main conclusions are summarized as follows:  \\begin{itemize}  \\item Metals are present in substantial amounts in a wide range of environments from the earlier epochs studied ($z\\sim 6$) until the present.  Diffuse IGM gas dominates the metal budget early on, while stars and shocked gas hold most of the metals today.  \\item The missing metals (outside of galaxies) at $z\\sim 2.5$ are mainly in intergalactic gas that is mildly hotter than pure photoionization temperatures, having $10^4\\la T\\la 10^5$~K.  However, this gas is not confined to halos, but is spread across a wide range of overdensities, with a peak around the cosmic mean density.  \\item Comparisons of our simulation to observations broadly show agreement in the amount of metals in various baryonic phases, and show that metals are spread relatively evenly across all phases.  \\item The global mean metallicity tracks about 50\\% higher than the diffuse IGM metallicity from $z=6\\rightarrow 1$, after which it continues increasing to about one-tenth solar at the present epoch.  Hot halo gas, which is predicted to have a luminosity-weighted metallicity of around one-fifth solar today, does not track the global metallicity, suggesting that clusters are unlikely to be representative samples of the universe at large.  This conclusions comes with a significant caveat, however, as our simulation volume is too small to produce actual clusters, so improved simulations are needed to confirm this.  \\item Stars and star-forming gas quickly obtain a mean metallicity of $\\ga 0.1Z_\\odot$ solar at $z\\sim 6$, indicating early enrichment in galaxies.  Today stars have a mean metallicity of one-half solar, while stars just forming now have typically solar metallicity.  \\item Our models predict a global \\ion{H}{i} column density-weighted metallicity that increases slowly with redshift from $z=6\\rightarrow 1$, and then actually decreases to $z=0$, never exceeding around one-third solar.  This is broadly consistent with observations of DLA metallicities, and does not require a substantial contribution from sub-DLAs.  \\item Metals generally move from lower temperature and more diffuse regions at early epochs towards hotter, more overdense regions at low redshifts.  This occurs because early blowout is very efficient at enriching diffuse gas, while at lower redshift outflows subside and large potential wells reaccrete gas.  \\item Oxygen, particularly \\ion{O}{iii}--\\ion{O}{vii}, provides the most effective tracer of metals across a wide range of cosmic phase space at $z\\la 2$, but a full metal accounting in the IGM will probably require next-generation instruments.  \\end{itemize}   Obtaining a complete census of metals in the Universe will require a concerted effort in studying a variety of systems and environments at various epochs.  Still, there is a realistic possibility of providing a full census of metals in the Universe in the foreseeable future.  This is certainly a formidable challenge for upcoming facilities, but the payoff is large.  By comparing such observations with cosmological simulations of outflows that track the interplay between various baryonic phases in mass, metals, and energy, it will be possible to stringently constrain models of galaxy formation and feedback.  This will dramatically advance our understanding of how galaxies and intergalactic gas interact in the cycle of accretion and feedback, and provide for the first time a complete view of how baryons evolve in cosmic phase space from the dark ages until the present."
    },
    "0608/astro-ph0608097_arXiv.txt": {
        "abstract": "We present the results of Giant Metrewave Radio Telescope \\ion{H}{i} 21-cm line observations of the eXtremely Metal Deficient (XMD) blue compact galaxy (BCG) HS~0822+3542. HS~0822+3542 is the smallest known XMD galaxy; from HST imaging it has been suggested that it actually  consists of two still smaller ($\\sim$100~pc sized) ultra-compact dwarfs that are in the process of merging. The brighter of these two putative ultra compact dwarfs has an ocular appearance, similar to that seen in galaxies that have suffered a penetrating encounter with a smaller companion. From our \\ion{H}{i} imaging we find that the gas distribution and kinematics in this object are similar to that of other low mass galaxies, albeit with some evidence for tidal disturbance. On the other hand, the \\ion{H}{i} emission has an angular size $\\sim$25 times  larger than that of the putative ultra-compact dwarfs. The optical emission is also offset from the centre of the \\ion{H}{i} emission. HS~0822+3542 is located in the nearby Lynx-Cancer void, but has a nearby companion LSB dwarf galaxy  SAO~0822+3545. In light of all this we also consider a scenario where the optical emission from HS~0822+3542 comes not from two merging ultra-compact dwarfs but from multiple star forming regions in a tidally disturbed galaxy. In this model, the ocular appearance of the brighter star forming region could be the result of triggered star formation. ",
        "introduction": "\\label{sec:intro}  The metallicity of the ISM of external galaxies is often quantified in terms of the oxygen abundance,  which, in turn, is generally expressed in terms of the quantity 12+$\\log$(O/H). A very small fraction of gas-rich dwarfs have 12+$\\log$(O/H) $\\lesssim$7.65, i.e. an oxygen abundance of less than $\\sim$1/10 solar. These dwarfs, (the majority of which are also members  of Blue Compact Galaxy (BCG) class), are termed eXtremely Metal Deficient (XMD) galaxies (see \\cite{kunth00} for a review). Low oxygen abundances are naturally expected for galaxies that have only recently experienced their first episode of star formation. For example, metallicities of $\\lesssim$1/10 solar are typical of high redshift Damped Lyman-$\\alpha$ systems (e.g., \\citealp{pettini97}), which are widely regarded as proto-galactic objects. It is now well established that the majority of (but not all) well studied XMD galaxies have a substantial mass fraction of their stellar mass in old stars (see e.g. a brief review in \\citealp{pustilnik06}) and are hence not forming stars for the first time. Nonetheless, their very low ISM metallicity still makes them the best local analogs of young galaxies at high redshift. Detailed studies of XMD galaxies should hence give insight into galaxy evolution in the early universe.  While the {\\it fraction} of XMD BCGs in the Local Universe is very small ($\\lesssim$2\\% of all known BCGs, e.g., \\citealp{pustilnik05}), the {\\it number} of known XMD galaxies has seen a dramatic increase in the recent past, via new large surveys for emission line galaxies (e.g. HSS, KISS and SDSS), and $\\sim$60 such galaxies are now known. However, detailed optical and radio observations are currently available only for a very small fraction of them.  As less evolved galaxies, XMDs are  expected  to be very gas-rich, and indeed for some objects (e.g., SBS 0335$-$052;  \\citealp{pustilnik01,pustilnik04}), the gas makes up $\\sim 99$\\% of the total baryonic mass. \\ion{H}{i} observations are hence critical in understanding the nature of these rare galaxies. For some of XMD BCGs from the Second Byurakan Survey (SBS) the integrated \\ion{H}{i} parameters are available in \\cite{thuan99}. More recently, \\cite{pustilnik06} (Paper I) conducted single dish HI observations for 22 XMD galaxies. Follow-up aperture synthesis \\ion{H}{i} observations of these galaxies are now in progress (e.g. BCG SBS 1129+576 \\cite{ekta06}, (Paper II)).  In this paper we present Giant Metrewave Radio Telescope (GMRT) \\ion{H}{i}~21cm observations of one of the nearest known XMD BCGs, HS 0822+3542 (12 + log(O/H) = 7.35; \\citealp{kniazev00}). \\cite{pustilnik03} found that this galaxy has a companion at a projected separation of 3.8\\arcmin\\ (i.e. $\\sim$11.4 kpc at the distance to the pair\\footnote{ We assume a distance of 11~Mpc. At this distance 1\\arcsec\\ corresponds to a linear separation of $\\sim$53~pc.}) -- the Low Surface Brightness (LSB) dwarf galaxy SAO 0822+3545. The separation between HS~0822+3542 and SAO~0822+3545 is much smaller than that of the primary beam of the GMRT antennas, and our observations are sensitive to \\ion{H}{i} emission from both of them. In Sec.~\\ref{sec:obs} we describe the GMRT observations and data reduction. Results are presented in Sec.~\\ref{sec:res} and discussed in Section~\\ref{sec:dis}. ",
        "conclusions": "\\label{sec:dis}  The main \\ion{H}{i} and optical parameters of HS~0822+3542 and SAO~0822+3545 are summarised in Tab.~\\ref{tab:param}. Using deconvolved Nordic Optical Telescope images, Pustilnik et al. (2003) found that the bright knot of optical emission in HS~0822+3542 consists of two components. Based on  HST images, \\cite{corbin05} state that these two components have a  surface brightness that is $\\sim 100$ times larger than that of the diffuse emission extending to the north west. The brighter of these two components (component ``A'' in the terminology of \\cite{corbin05}; see their Fig~1) has an ocular morphology, with a central blue star forming region surrounded by a ring of redder stars. The fainter component (component ``B'') is $\\sim$5 times less luminous and more irregular in appearance. Morphologies similar to that of component ``A'' are seen in galaxies where there has been a penetrating encounter with a smaller companion (e.g., the Cartwheel galaxy). Based on this, \\cite{corbin05} suggest that HS~0822+3542 is actually a system of two merging ultra-compact dwarfs, and that it represents a dwarf galaxy in the process of formation.  \\begin{table} \\caption{Main parameters of the studied galaxies} \\label{tab:param} \\begin{tabular}{lcc} \\\\ \\hline \\hline Parameter                           & HS~0822+3542            & SAO~0822+3545    \\\\ \\hline R.A.(J2000.0)(opt)                  & ~~08 25 55.47           & ~~08 26 05.59     \\\\ DEC.(J2000.0)(opt)                  & $+$35 32 32.9           & $+$35 35 25.7     \\\\ B$_{\\rm tot}$                       & 17.92$\\pm$0.05$^{(1)}$  & 17.56$\\pm$0.03$^{(3)}$    \\\\ V$_{\\rm hel}$(\\ion{H}{i})(\\kms)     & 726.6$\\pm$2$^{(2)}$     & 740.7$\\pm$1.3$^{(2)}$  \\\\ Dist$^{(3)}$ (Mpc)                  & 11.0                    & 11.0            \\\\ M$_{\\rm B}^0$ $^{(4)}$              &  --12.49                & --12.85          \\\\ Opt. size (\\arcsec)$^{(5)}$         & 14.8$\\times$7.4$^{(1)}$ & 28.2$\\times$15.5$^{(3)}$  \\\\ Opt. size (kpc)                     & 0.79$\\times$0.39        & 1.50$\\times$0.83$^{(3)}$  \\\\ Inclination (deg)$^{(6)}$           & 62                      & 59               \\\\ 12+$\\log$(O/H)  \\                   & 7.44$^{(3)}$          & ---               \\\\ \\ion{H}{i} int.flux$^{(7)}$         & 0.27$\\pm$0.06$^{(2)}$   & 0.83$\\pm$0.08$^{(2)}$      \\\\ W$_\\mathrm{50}$ (km s$^{-1}$)       & 21.7$\\pm$4$^{(2)}$      & 35.5$\\pm$3$^{(2)}$          \\\\ V$_\\mathrm{rot}$ (\\ion{H}{i})(\\kms) & $\\le$5$^{(2)}$         & ~14$^{(2)}$          \\\\ M(\\ion{H}{i}) (10$^{8} M_{\\odot}$)  & 0.77$^{(2)}$            &2.36$^{(2)}$       \\\\ M$_{\\rm dyn}$ (10$^{8} M_{\\odot}$)  & 4$^{(2)}$             & 14$^{(2)}$       \\\\ M(\\ion{H}{i})/L$_{\\rm B}$$^{(8)}$   & 0.50$^{(2)}$            & 1.10$^{(2)}$       \\\\ \\hline\\hline  \\multicolumn{3}{l}{(1) -- from \\cite{kniazev00}} \\\\ \\multicolumn{3}{l}{(2) -- derived in this paper; velocities are from a gaussian fit } \\\\ \\multicolumn{3}{l}{(3) -- from \\cite{pustilnik03}} \\\\ \\multicolumn{3}{l}{(4) -- corrected for Galactic extinction A$_{\\rm B}$=0.20} \\\\ \\multicolumn{3}{l}{(5) -- $a \\times b$ at $\\mu_{\\rm B}=$25\\fm0~arcsec$^{-2}$} \\\\ \\multicolumn{3}{l}{(6) -- from the optical diameters and assuming an intrinsic}\\\\ \\multicolumn{3}{l}{\\hskip 0.3in axis ratio of 0.2} \\\\ \\multicolumn{3}{l}{(7) -- in units of Jy$\\cdot$\\kms}\\\\ \\multicolumn{3}{l}{(8) -- in units of ($M/L_\\mathrm{B}$)$_{\\odot}$}  \\\\ \\end{tabular} \\end{table}  Our \\ion{H}{i} observations show that the gas in this system appears to have settled into a disk with fairly regular kinematics. We note again however, that this may be a consequence of our marginal angular resolution. In any case, a regular HI disk may be consistent with the merger scenario -- recent numerical simulations show that in gas-rich mergers, the gas quickly settles into a disk (\\citealp{springel05}). On the observational front, there are cases of merging galaxies where a large fraction of the gas does seem to have settled into a  disk (e.g., IC~2554, \\cite{koribalski03}; NGC~3310 \\citealp{kregel01}). HI observations of the Cartwheel galaxy itself (\\citealp{higdon96}) show that the gas is in a disk, albeit with most of the gas in an expanding ring at the periphery of the disk. While the high resolution image of HS~0822+3542 (Fig.~\\ref{fig:m0ovhr}) could be interpreted as emission from a ring, such an intepretation would require that the HI disk should be viewed almost perfectly edge on, i.e. the HI disk inclination must be much higher than that inferred from the optical image (see ~Tab.\\ref{tab:param}). On the other hand, the size scale of  the gas disk is much larger than that of the optical components seen in the HST image. The larger component~A has an angular size of only $\\sim$1.6$^{''}$; from Fig.~\\ref{fig:m0ovhr}[A]) this is $\\sim$26 times less than the deconvolved size (at a column density of 1 M$_\\odot$/pc$^{2}$) of the \\ion{H}{i} disk.  The ratio of \\ion{H}{i} mass to the optical luminosity in these two components is however not so extreme, with M(HI)/(L$_{\\rm B}^A + L_{\\rm B}^B) \\sim$1.7 (where we use the values for L$_{\\rm B}^A + L_{\\rm B}^B$ estimated by \\citealp{corbin05}). For the system as a whole we have (from Tab.~\\ref{tab:param}), M(HI)/L$_{\\rm B} \\sim$0.50. Note that the L$_B$ value listed in Tab.~\\ref{tab:param}  has (as is generally the case for BCGs) a substantial contribution from nebular emission (including line emission) while the blue  luminosities given by \\cite{corbin05} for the two components are corrected for the contribution from strong lines.  SAO~0822+3545 has a radial velocity difference of only $\\sim$16~\\kms\\ and a projected separation of only $\\sim$11~kpc from HS~0822+3542. The charcteristic dynamical time for this pair $\\tau_{\\rm dyn}\\sim r_p/\\Delta V \\sim$700~Myr, comparable to the rotation periods of the galaxies in the pair. It would hence seem unreasonable to treat HS~0822+3542 as an isolated system. The peculiar morphology of HS~0822+3542, with its large scale \\ion{H}{i} tails and the offset between the optical and \\ion{H}{i} emission could be due to interaction with  SAO~0822+3545. The \\ion{H}{i} morphology of SAO~0822+3545, as seen on the high-resolution map in Fig.~\\ref{fig:m0ovhr}[D], is also very irregular, with a tail like feature to the north of the main body and an \\ion{H}{i} `hole' the SE part of the optical body.  In view of the strong possibility of an external tidal trigger to the star formation in HS~0822+3542, it is worth taking a relook the scenario  (earlier suggested by \\citealp{pustilnik03}) where components ``A'' and ``B'' are star forming regions in a tidally disturbed galaxy.  In this scenario, the ``ocular'' appearance of component ``A'' could be the result of triggered star formation (\\cite{elmegreen77};\\citealp{elmegreen02}), where feedback from an older central star cluster has resulted in star formation in a surrounding ring. In this context, it is interesting to note that \\cite{pustilnik03} found evidence for an expanding superbubble in HS~0822+3542. Our \\ion{H}{i} data have too low a signal-to-noise ratio to search for corresponding signatures in the \\ion{H}{i}  maps. However, we note that at in the highest resolution image (Fig.~\\ref{fig:m0ovhr}[C]) the \\ion{H}{i} column density peaks on either side of the optical emission. Triggered star forming rings around an existing star cluster have been observed both in our own galaxy (\\cite{deharveng03}, \\citealp{deharveng05}) as well as in external dwarf galaxies (Sextans A; \\citealp{vandyk98}).  Further, the evidence presented by \\cite{pustilnik03} for recent star formation elswhere in HS~0822+3542 (e.g. the arc at the  northwestern edge (or the tip of the ``tidal plume'' in the terminology of \\citealp{corbin05})) is also probably easier to understand in this scenario than in one where HS~0822+3542 is an ongoing merger of two ultra-compact dwarfs. Finally, we note that though \\cite{corbin05} state that the surface brightness of components ``A'' and ``B'' are a factor of $\\sim 100$ higher than that of the stellar ``tidal'' plume, the values for the surface brightness given in their paper  (23 mag arcsec$^{-2}$  for the plume and 20.6 mag arcsec$^{-2}$ for component ``B'' and the ring in component ``A'') imply that the difference in surface brightness is only $\\sim 9$. There hence appear to be good reasons for regarding the optical knots in HS~0822+3542 as star forming regions in a single galaxy, as opposed to merging ultra-compact dwarfs. Higher quality data and detailed numerical modelling would however be required to conclusively distinguish between these two models.  As noted by \\cite{pustilnik03}, the LSB galaxy SAO~0822+3545 is unusually blue. Further, despite a small downward correction of its \\ion{H}{i} flux relative to that given in \\cite{pustilnik03} SAO~0822+3545 remains one of the most gas-rich known LSB galaxies, with the gas-mass fraction M$_{\\rm gas}$/(M$_{\\rm gas}$+M$_{\\rm star}$) in the range of 0.77 to 0.88, depending on the age/mass of the old stellar population.  Comparison of several evolutionary models for SAO~0822+3545 with the constraints provided by its integrated colours, EW(H$\\alpha$), M$_{\\rm gas}$ and M$_{\\rm dyn}$ favours the scenario where the optical emission comes from a mixture of a young stellar population ($t_{1}$ = 10~Myr, M$_{\\rm young}$ = (2.4--4.4)$\\times$10$^5$ M\\sunn) and an `old' stellar population (with age in the range of 0.25 to 10 Gyr and mass M$_{\\rm old}$ of 10-30 M$_{\\rm young}$). Such a `starburst' event with the instantaneous formation of $\\sim$3$\\times$10$^5$M\\sunn\\ and a release of a large amount of kinetic energy in the form of stellar winds and SNR should have a measurable effect on the neutral gas; in particular one might expect to see \\ion{H}{i} ``holes''. Indeed, in Fig.~\\ref{fig:m0ovhr}[D]  we see that in the optically brighter south-east edge there is a clear deficit of \\ion{H}{i} gas. We note however that this deficit is seen mainly in the high resolution image, and that it is possible that the ``hole'' is filled with low column density gas. If we do treat the feature as a ``hole'', it corresponds to a linear size of $\\sim$500~pc. \\cite{mccray87} give a relation between the hole size $R$ and the number of supernovae ($N_*$) producing the hole, viz.  $$R ~\\sim 97 (N_* E_{51}/n_0)^{1/5} {t_7}^{3/5}$$  where $E_{51}$ is the energy per supernovae in units of $10^{51}$ erg, $n_0$ is the ambient density in atoms cm$^{-3}$ and $t_7$ is the age in units of $10^7$ yr. If we take $E_{51},n_0$ and $t_7$ to be $\\sim 1$,  we find that the energy output from $\\sim 5$ supernovae would be sufficient to produce a hole this size. This number of supernovae is easily produced by a young stellar population with mass $\\sim 2\\times 10^5$M\\sunn ~(from e.g. PEGASE2; \\citealp{fioc99}). While detailed modelling and better HI and optical data would be essential to see if the HI ``hole'' that we see has been produced by feed back from the star burst, the above calculation does indicate that a young stellar population with mass similar to that inferred by \\cite{pustilnik03} should in principle be able to produce a hole comparable in size to the galaxy itself.  Finally we note  that the peak inclination corrected \\ion{H}{i} column densities in HS~0822+3542 are comparable to the observed threshold for star formation in dwarf galaxies ($\\NHI \\sim \\times 10^{21}$ atoms~cm$^{-2}$; \\citealp{skillman87}). While star formation is indeed ongoing in HS~0822+3542, the peak \\ion{H}{i} column density is  offset from the sites of current star formation (~\\ref{fig:m0ovhr}[C]). This may be a consequence of feedback (either in the form of ionization or evacuation). Similarly, the \\ion{H}{i} column density in SAO~0822+3545 (see Fig.~\\ref{fig:m0ovhr}[D]) is above the threshold value over a large fraction of its area. In particular, the prominent double \\ion{H}{i} blob on the east edge of the galaxy has high density, but no traces of star formation. There are several other known cases of dwarf galaxies where current star formation is  offset from the location of the peak HI column density (e.g. SBS~0335$-$052; \\citealp{pustilnik01}), and, in a detailed study of the relation between star formation and HI column density in a sample of nearby extremely faint dwarf galaxies \\cite{begum06} found no one to one correspondence between high \\ion{H}{i} column density and ongoing star formation.  In summary, our GMRT observation of the galaxy pair HS~0822+3542/SAO~0822+3545 shows (i)~some evidence for tidal interaction, (ii)~that the HI emission from HS~0822+3542 has an angular size $\\sim 25$ times larger than that of the optical components identified in the HST imaging, and (iii)~that the HI properties of the galaxies are in general similar to those of comparably sized gas rich dwarf galaxies. We suggest that the optical knots in HS~0822+3542 are star forming regions in a tidally disturbed dwarf galaxy as opposed to being merging ultra-compact dwarfs.  {\\bf Acknowledgements} The observations presented in this paper were obtained using the GMRT which is operated by the National Centre for Radio Astrophysics (NCRA) of the Tata Institute of Fundamental Research (TIFR), India. SAP acknowledges partial support from the Russian state program \"Astronomy\" and from the RFBR grant No.06-02-16617. This research has made use of the NASA/IPAC Extragalactic Database (NED), which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. The authors acknowledge the SDSS database for the images used for this study. Funding for the Sloan Digital Sky Survey (SDSS) has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S. Department of Energy, the Japanese Monbukagakusho, and the Max Planck Society. The SDSS Web site is http://www.sdss.org/. The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions. The Participating Institutions are The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, The Johns Hopkins University, the Korean Scientist Group, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington. We thank the referee for comments that resulted in a substantial improvement in the quality of this paper."
    },
    "0608/astro-ph0608574_arXiv.txt": {
        "abstract": "We imaged five objects near the star forming clouds of Ophiuchus with the Keck Laser Guide Star AO system.  We resolved Allers et al. (2006)'s \\#11 (Oph 16222-2405) and \\#16 (Oph 16233-2402) into binary systems. The \\#11 object is resolved into a $243$ AU binary, the widest known for a very low mass (VLM) binary. The binary nature of \\#11 was discovered first by Allers (2005) and independently here during which we obtained the first spatially resolved $R\\sim2000$ near-infrared (J \\& K) spectra, mid-IR photometry, and orbital motion estimates. We estimate for 11A and 11B gravities (log(g)$>3.75$), ages ($5\\pm2$ Myr), luminosities (log(L/$L_{\\sun}$)=$-2.77\\pm0.10$ and $-2.96\\pm0.10$), and temperatures ($T_{eff}=2375\\pm175$ and $2175\\pm175$ K). We find self-consistent DUSTY evolutionary model (Chabrier et al. 2000) masses of $17^{+4}_{-5} M_{Jup}$ and $14^{+6}_{-5} M_{Jup}$, for 11A and 11B respectively. Our masses are higher than those previously reported (13--15 $M_{Jup}$ and 7--8 $M_{Jup}$) by Jayawardhana \\& Ivanov (2006b). {\\it Hence, we find the system is unlikely a ``planetary mass binary'', (in agreement with Luhman et al. 2007) but it has the second lowest mass and lowest binding energy of any known binary}. Oph \\#11 and Oph \\#16 belong to a newly recognized population of wide ($\\ga100$ AU), young ($<10$ Myr), roughly equal mass, VLM stellar and brown dwarf binaries. We deduce that $\\sim 6\\pm3\\%$ of young ($<10$ Myr) VLM objects are in such wide systems. However, only $0.3\\pm0.1\\%$ of old field VLM objects are found in such wide systems.  Thus, young, wide, VLM binary populations may be evaporating, due to stellar encounters in their natal clusters, leading to a field population depleted in wide VLM systems. ",
        "introduction": "\\subsection{Imaging Oph \\#11 and Oph \\#16 at J, H, Ks, \\& Ls}  We independently discovered the binary nature of the Oph 11 system with the Gemini NIRI camera (Hodapp et al. 2003) on July 1, 2006 (UT) in the J, H, \\& Ks bands (all other objects in our survey were observed with the Keck II LGS AO system).  We then obtained Ls ($\\sim3.0\\mu$m) images with the NIRC camera on the Keck I telescope on July 7, 2006 (UT). In all cases seeing was excellent, yielding images with FWHM$\\sim$0.3'' in the Ks band. We reduced these data in the standard manner (Close et al. 2002, 2003). All images were fully flat-fielded, sky and dark subtracted, bad pixel masked, aligned, and medianed.  Figures 1 and 2 and Tables 1-3 present the measured and derived characteristics of the Oph 11 and 16 binaries.  Since the northern component is brighter (in the NIR) we name it Oph 11A and the apparently cooler southern component Oph 11B.  The system is also known as 2MASS J16222522-2405138 and Jayawardhana \\& Ivanov (2006b) refer to it as Oph 162225-240515 (and Oph 1622). We will refer to this system as Oph \\#11 (and as Oph 1622-2405) since it was first mentioned in the widely available literature as \\#11 by Allers et al. (2006) to be an interesting low-mass object (in fact it was 11B that Allers was referring to; Allers et al. 2007). The designation of Oph 1622 is inappropriate since Allers et al.'s \\#10, \\#11, \\#12, and \\#13 objects would all also be ``Oph 1622'' due to their similar RA.  The night of July 1, 2006 (UT) was photometric and zeropoints from photometric standards were used for calibration at J, H, and Ks. There are small ($2MASS-Gemini=$ $-0.12\\pm0.09$, $0.16\\pm0.10$, $0.06\\pm0.10$ mag) differences between our calibration and the J, H, and Ks fluxes of \\#11 in the 2MASS point source catalog. However, the $1.9\\arcsec$ separation of the binary likely leads to some errors in the 2MASS $4\\arcsec$ aperture photometry, and so we adopt our values (which take the separations of the binary into account) throughout the present paper.   We also analyzed the public ${\\it Spitzer}$ [3.6], [4.5], [5.8], [8.0] $\\mu$m IRAC images of \\#11. Utilizing a custom IRAC PSF fitting program we were able to detect (for the first time) both components of the binary in all IRAC passbands (see, for example, Figure 3). In all cases a double PSF model of a 1.94$\\arcsec$ binary with PA=$176^{\\circ}$ was a significantly better fit to the IRAC images than a single IRAC PSF.  \\subsection{NIRSPEC J and H band Spectra of 11A and 11B}  Spatially resolved $J$- and $K$-band spectra of the \\#11 system components were obtained on August 9, 2006 with NIRSPEC in low resolution mode on the Keck~II telescope (Mclean et al. 1998, 2000).  We used a 0$\\farcs$57$\\times$42\\arcsec~slit and the image rotator to maintain the $\\sim$180\\degr~ position angle with both components on the slit. The 0$\\farcs$57-wide slit corresponds to 3 pixels on the detector and a nominal resolving power of R$\\sim$2000.  $J$-band observations utilized the NIRSPEC-3 order-sorting filter (1.143$-$1.375~$\\micron$), and $K$-band observations used the $K$ filter (1.996$-$2.382~$\\micron$).  For each filter a set of four frames was obtained in an ABBA nod pattern, nodding $\\sim$20\\arcsec~along the slit between the A and B positions.  Integration times for each exposure were 360~seconds in the $J$-band and 180~seconds in the $K$-band.  An A0V calibrator star, HIP 80224, was observed in each band for a 60 second integration at each nod position and at a similar airmass to facilitate the removal of telluric absorption features from the target spectra.  Flat, dark, and arc lamp exposures were obtained with the calibration unit internal to NIRSPEC.   The NIRSPEC data were reduced using the REDSPEC code,\\footnote{See http://www2/keck.hawaii.edu/inst/nirspec/redspec} a package of IDL procedures developed specifically for the reduction of NIRSPEC spectra by S. Kim, L. Prato, and I. S. McLean.  The reduction of low resolution NIRSPEC data with REDSPEC is described in detail in McLean et al. (2003), and summarized here.  REDSPEC uses arc lamp frames to spatially and spectrally rectify the raw data and to determine a wavelength solution from observed neon and argon emission lines of known wavelengths.  Following rectification, pairs of target nods are subtracted to remove background and divided by the rectified, dark-subtracted flat field.  The target and calibrator spectra are extracted by summing over 9-11 rows of data centered on the peak of the trace.  For the Oph~11 observations, the peaks of the traces are separated by 18 pixels, resulting in minimum contamination between components.  Intrinsic spectral features of the calibrator (Pa$\\beta$ at 1.282~$\\micron$ in the $J$-band and Br$\\gamma$ at 2.166~$\\micron$ in the $K$-band) are removed by linear interpolation over the line. The calibrator spectrum is divided from the target spectrum, and the target spectrum is multiplied by a 9770~K blackbody (the $T_{eff}$ of HIP~80224) to restore the spectral slope.  To produce the final spectrum, the extracted spectra from all four nods are averaged and normalized to the continuum level at a given wavelength.   The actual resolution of the observations was estimated by fitting a Gaussian to several emission lines in the raw arc lamp frames, which resulted in R$\\sim$1500 for $J$-band and R$\\sim$1900 for $K$-band. The signal-to-noise is calculated from the maximum peak-to-peak variation and estimated to be S/N~$\\gtrsim$~100 for each spectrum. ",
        "conclusions": "We have obtained the first spatially resolved $R\\sim2000$ near-infrared (J \\& K) spectra, mid-IR photometry, and orbital motion estimates for the wide brown dwarf binary Oph \\#11 (Oph 1622-2405). We estimate for 11A and 11B gravities (log(g)$>3.75$), ages ($5\\pm2$ Myr), luminosities (log(L/$L_{\\sun}$)=$-2.77\\pm0.10$ and $-2.96\\pm0.10$), and temperatures ($T_{eff}=2375\\pm175$ and $2175\\pm175$ K). We find self-consistent DUSTY evolutionary model (Chabrier et al. 2000) masses of $17^{+4}_{-5} M_{Jup}$ and $14^{+6}_{-5} M_{Jup}$, for 11A and 11B respectively. Our masses are higher than the previously reported 13--15 $M_{Jup}$ and 7--8 $M_{Jup}$ masses of Jayawardhana \\& Ivanov (2006b). {\\it Hence, we find the system is unlikely a ``planetary mass binary'', but it has the second lowest mass and lowest binding energy of any known binary}.  The \\#16 binary (Oph 1623-2402) is also an unusually wide (projected separation of 212 AU), very low-mass (VLM), binary composed of a $\\sim100 M_{Jup}$ primary (16A) and a $\\sim73 M_{Jup}$ (16B) secondary.   While young VLM datasets are dominated by small number statistics, they are suggestive of a moderately common ($6\\pm3\\%$) young ($<10$ Myr) population of wide ($>100$ AU), very low-mass ($M_{tot}<0.2M_{\\sun}$), binary systems like Denis PJ161833.2 -251750.4, 2MASS J11011926-7732383, Oph\\#11 (Oph 1622-2405), and Oph\\#16 (Oph 1623-2402). Our estimate of the wide young binary fraction agrees well with the $5^{+6}_{-2}\\%$ found in the Upper Sco by Bouy et al. (2006). We argue that, over time, the majority of these wide VLM systems become unbound leaving just the surviving (most tightly bound) fraction observable in the field today.  We suspect these wide systems likely become unbound since only $\\sim 0.3\\pm0.1\\%$ of field VLM objects are wider than 100 AU today. Hence it appears that the field population is depleted in wide VLM binaries by a (selection effect corrected) factor of $X_{evap} = 15^{+20}_{-10}$. The exact mechanism for this evaporation of wide VLM systems is unlikely to be interactions with stars in the field or with molecular clouds and their sub-clumps. However, from the Fokker-Planck solutions of Weinberg et al. (1987), we find stellar encounters with ``birth cluster stars'' may be efficient in dissolving these wide binaries. Indeed such encounters may help boost the observed $sep$ past the $sep^{catastrophic*}_{cluster} \\sim 80\\left(\\frac{M_{tot}}{0.1 M_{\\sun}}\\right)$ AU limit where catastrophic cluster encounters can occur. More detailed numerical studies will be required to better constrain the evolution of this evaporating population."
    },
    "0608/astro-ph0608604_arXiv.txt": {
        "abstract": "We compute f-mode travel-time sensitivity kernels for flows.  Using a two-dimensional model, we show that it is important to account for several systematic effects, such as the foreshortening and the projection of the velocity vector onto the line of sight. Correcting for these effects is necessary before any data inversion is attempted away from the center of the solar disk. ",
        "introduction": "The solar f-modes are useful tools for studying perturbations that reside within the top $2$~Mm of the surface of the Sun, and have been used as a diagnostic as such before, e.g., \\citet{duvall2000}. The linear forward problem of time-distance helioseismology is to calculate functions, or kernels, which give the sensitivity of travel-time measurements to these particular perturbations of the f-modes. For a discussion of this general procedure, which is the one that we will follow  here, see \\citet{gizon2002}.  In this study we are interested in \\textit{flow} kernels, i.e., functions which give the two-dimensional sensitivity of travel-time measurements to small amplitude, two-dimensional, spatially varying flows. The flow kernel, $\\bK$, is defined as \\begin{equation} \\delta\\tau=\\int\\!\\!\\int_S d^2\\br\\,\\bu(\\br)\\cdot\\bK(\\br), \\label{kern-def} \\end{equation} where $\\delta\\tau$ is the perturbation to the travel time, $\\bu=(u_x,u_y)$ is the two-dimensional flow perturbation, and the integral is over the position vector $\\br=(x,y)$ over the whole solar surface. The kernel $\\bK=(K_x,K_y)$ is a two-dimensional vector which has sensitivity to flows in the $\\bx$ and $\\by$ directions.  We calculate the kernels from an approximate solar model, which utilizes the Born approximation in estimating the scattering from solar inhomogeneities. This has been shown to be more accurate than calculating kernels from a ray approximation \\citep{birch2004}, since it includes finite-wavelength effects.  Ultimately, the goal is to perform inversions with real time-distance measurements to extract the unknown $\\bu$ from Eq.(\\ref{kern-def}), given $\\delta\\tau$ and $\\bK$.  However, there are several difficulties in doing this since the travel-time measurements incur systematic effects from the observing process, particularly due to the center-to-limb variations. To overcome these problems, one strategy is to calculate the sensitivity kernels from the outset accounting for the systematic effects so that the whole procedure - the forward calculation and the inversions - remains consistent.  We have accomplished this for two important effects: the foreshortening of the data pixel, and the projection of the solar velocity field onto the line of sight, both of which will be defined in detail below. We present an analysis of the contribution that these two effects have on the power spectrum of f-modes and on the kernels. For brevity, we leave the details of the full calculation of the power spectra and kernels for a future publication \\citep{jackiewicz2006}, and only focus here on the role of the systematic effects mentioned above on the final results. ",
        "conclusions": ""
    },
    "0608/astro-ph0608432_arXiv.txt": {
        "abstract": "We present optical photometry and spectroscopy of the type IIP supernova SN 2004et that occurred in the nearby galaxy NGC 6946. The observations span a time range of 8 days to 541 days after explosion. The late time bolometric luminosity and the H$\\alpha$ luminosity in the nebular phase indicate that $0.06\\pm0.02{M_\\odot}$ of $^{56}$Ni was synthesised during the explosion. The plateau luminosity, its duration and the expansion velocity of the supernova at the middle of the plateau indicate an explosion energy of $E_{\\rm{exp}}= 1.20^{+0.38}_{-0.30} \\times 10^{51}$ ergs. The late time light curve and the evolution of the [OI]  and H$\\alpha$ emission line profiles indicate the possibility of an early dust formation in the supernova ejecta. The luminosity of  [OI] 6300, 6364~\\AA \\, doublet, before the dust formation phase, is found to be comparable  to that of SN 1987A at similar epochs, impling an oxygen mass in the range $1.5 - 2{M_\\odot}$, and a main sequence mass of $20{M_\\odot}$ for the  progenitor. ",
        "introduction": "\\label{intro} Supernovae have been classified mainly based on their spectrum near maximum, the events which show the presence of hydrogen lines have been termed as supernovae type II (SNe II) and those which do not show hydrogen are classified as type I. Type II events are thought to arise from the gravitational collapse of stars more massive than $8\\,M_\\odot$. SNe II have been further divided based on their light curve \\citep{baronr} as type II-L (linear) and type II-P (plateau). Supernovae type IIP (SNe IIP) are characterized by a plateau of nearly constant luminosity in their light curve, originating because of propagation of a cooling and recombination wave through the supernova envelope.  \\cite{hamuy03} has shown that SNe IIP form a sequence from low-luminosity, low-velocity, nickel-poor events to bright, high velocity, nickel-rich objects. Further, there is an indication that more massive progenitors produce more energetic explosion and supernovae with greater energies produce more nickel. The direct identification of the progenitor of a few SNe IIP indicates these events arise from stars with masses in the range $\\sim8\\,-15\\,{\\rm{M}}_\\odot$ (\\citealt{smartt}, \\citealt{van},  \\citealt{li3}).  SN 2004et was discovered by Moretti on Sept. 27, 2004 in the nearby starburst galaxy NGC 6946, which has already produced seven supernovae during  1917 to 2003. Based on a high resolution spectrum that showed a relatively featureless spectrum with a very broad, low contrast H$\\alpha$ emission, \\cite{zwitter} classified the supernova as a type II event, which  was later confirmed by a low resolution spectrum taken on 2004 Oct 01 by \\cite{filippenko04}. The P-Cygni profile of H$\\alpha$ was greatly dominated by the emission component, while the other hydrogen Balmer lines had a more typical P-Cygni profile. The continuum was quite blue, although there was very low flux shortward of 4000\\AA. The supernova was detected in the radio frequencies at 22.460 GHz and 8.460 GHz on 2004 Oct 5.128 \\citep{stockdale}, suggesting the presence of appreciable circumstellar material  around SN 2004et.  Based on pre-explosion images of NGC 6946, \\cite{li1} identified the candidate progenitor as a yellow supergiant with an estimated zero age main sequence  mass of ${15}^{+5}_{-2}{M_\\odot}$. The identification of the progenitor was confirmed with post-outburst Hubble Space Telescope images obtained on 2005 May 02 \\citep{li2}. This makes SN 2004et one of the few core-collapse supernovae with a directly identified progenitor.  The proximity and brightness of SN 2004et made it an ideal target for an intensive monitoring. We present in this paper the results based on extensive photometric and spectroscopic observations of SN 2004et during $\\sim$ 8 - 541 days since explosion. ",
        "conclusions": ""
    },
    "0608/hep-ph0608041_arXiv.txt": {
        "abstract": "We present results of a microscopic calculation using NJL-type model of possible spin-one pairings in two flavor quark matter for applications in compact star phenomenology. We focus on the color-spin locking phase (CSL) in which all quarks  pair in a symmetric way, in which color and spin rotations are locked. The CSL condensate is particularly interesting for compact star applications since it is flavor symmetric and could easily satisfy charge neutrality. Moreover, the fact that in this phase all quarks are gapped might help to suppress the direct Urca process, consistent with cooling models. The order of magnitude of these small gaps ($\\simeq 1$ MeV) will not influence the EoS, but  their also small critical temperatures ($T_c \\simeq 800$ keV) could be relevant in the late stages neutron star evolution, when the temperature falls below this value and a CSL quark core could form. ",
        "introduction": "\\label{intro}  The most favorable places in nature where color superconducting states  of matter are expected to occur are the interiors of compact stars, with temperatures well below 1 MeV and central densities exceeding the nuclear saturation density $\\rho_0$ by several times.  Color superconductivity has been widely studied from non-perturbative low-energy QCD models where gaps of the order of magnitude of $\\simeq 100$ MeV has been calculated (\\cite{Rapp:1997zu,Alford:1997zt}). One of the effective models most used is the Nambu Jona-Lasinio (NJL) model that considers that the quarks  interact locally by a 4-point vertex effective force and  disregards the gluon degrees of freedom. The model uses an attractive interaction in the scalar meson channel that causes spontaneous chiral symmetry breaking if the coupling is strong enough. At the mean field level, the model shows how quarks acquire a dynamical constituent mass, which is proportional to the vacuum expectation value of the scalar field. The NJL model has then been extended and widely used to described successfully chiral restoration at finite temperatures and color superconductivity at finite density (for a review see \\cite{Buballa:2003qv}).  For compact star applications, color super\\-conduc\\-ting quark matter phases enforcing color and charge neutrality has been widely studied (\\cite{Alford:2000sx,Steiner:2002gx}). It has been shown that local charge neutrality may disfavor the occurrence of phases with large gaps where quarks with different flavor pair in a spin-0 condensate, like the 2SC phase (\\cite{Alford:2002kj}). On the other hand, NJL-type model calculations show that the intermediate density region of the neutral QCD phase diagram, where the quark chemical potential is not sufficiently large to have the strange quark deconfined ($\\mu \\geq 430-500$ MeV, see \\cite{Buballa:2001gj,Neumann:2002jm}), might be dominated by $u,d$ quarks (\\cite{Ruster:2005jc,Blaschke:2005uj}).  If this is the case, 2-flavor quark matter phases may occupy a large volume in the core of compact stars (\\cite{Grigorian:2003vi,Shovkovy:2003ce}).  While the stability of  neutral 2SC pure phase is rather model dependent and might be unlikely for moderate coupling constants (\\cite{Aguilera:2004ag,GomezDumm:2005hy}), phases with other pairing patterns (or none at all) become important for the phenomenology of neutron stars. Then, besides other possibilities% \\footnote{no pairing (normal quark matter), pairing with displacement of the Fermi surfaces (LOFF or crystalline structure) or deformation of the Fermi surfaces , interior gap structure, gapless 2SC or gCFL, etc.}, quarks could pair in spin-one condensates (\\cite{Schafer:2000tw,Alford:2002rz}).  These condensates with small pairing gaps ($\\Delta \\simeq 1$ MeV), are not expected to have influence on the equation of state but could strongly affect the transport and thermal properties of quark matter and therefore leads to consequences for the phenomenology of compact stars.  An energy gap in the quasiparticle excitation spectrum introduces a suppression of the neutrino emissivity and the specific heat of the paired fermions by a Boltzmann factor $\\exp(-\\Delta/kT)$ for $T$ smaller than the critical temperature  $T_c$ for pair formation. Thus, since the neutrino luminosity is strongly dominated by the neutrino emission from the core, the cooling of a neutron star during its early life  will be affected by the dense matter pairing pattern. On the other hand, most of the specific heat of the star is provided by the core and it shows a discontinuity when the star temperature crosses $T_c$. A comparison of cooling curves with observations could provide a hint to discriminate dense matter phases.  Several studies have shown that the occurrence of unpaired quarks in the core leads to rapid cooling via the direct Urca process , incompatible with the observations. Thus, phases that present no gapless modes with small pairing gaps ($\\Delta \\simeq 1$ MeV) prevent the direct Urca to work uncontrolled and  might help to give a consistent picture of the observed data (see \\cite{Page:2000wt,Page:2004,Yakovlev:2004,Grigorian:2004jq}).  On the other hand,  it has been demonstrated that the CSL phase exhibit a Meissner effect and the magnetic field of a neutron star would be  expelled from a CSL quark core if it does not exceed the critical magnetic field $B_{c}$ that destroy the superconducting phase \\cite{Schmitt:2003xq}. This hypothesis might be consistent with recent investigations that indicate the crust confination of the magnetic field \\cite{Geppert:2004sd,Perez-Azorin:2005ds}.  To study this qualitatively, we present in this work a summary of the results obtained in previous works of NJL model calculations of two spin-one pairing patterns, focusing in the CSL phase, and a new analysis of features that are relevant to apply the CSL in phenomenological applications, like compact star cooling or protoneutron star evolution. ",
        "conclusions": "We summarize here the interesting features that the  CSL phase presents for compact star applications: \\begin{itemize} \\item since Cooper pairs are single flavor they are not  affected by charge neutrality \\item color neutrality is automatically fulfilled \\item there are no gapless modes. This has important consequences e.g. in neutrino emissivities, suppressing the direct Urca process that leads to a rapid cooling \\item their small gaps ($\\Delta \\simeq 1$ MeV) will not influence on the EoS, and the pressure is expected to be approximately the one of the normal $ud$-quark matter without pairing. Similar to the latter, stable hybrid star configurations could be obtained with a relatively large quark matter core. \\item it exhibits a Meissner effect and the magnetic field of a neutron star would be  expelled from a CSL quark core if it does not exceed the critical magnetic field $B_{c}$ that destroy the superconducting phase. This hypothesis might be consistent with recent investigations that indicate the crust confination of the magnetic field. Moreover, the hypothesis of a Type I superconductor is supported by the inferred precession of some compact objects. \\end{itemize}  This study presents a qualitative analysis of the spin-1 pairing and cannot be taken as conclusive to decide whether this phase could be realized in neutron star cores. However,  the features listed show that CSL could be consistent with the phenomenology of compact stars in many aspects."
    },
    "0608/astro-ph0608354_arXiv.txt": {
        "abstract": "We report the discovery of an eclipsing X-ray binary in M101, the first such system to be discovered outside the Local Group.  Based on a sequence of 25 Chandra observations that sample a wide range of orbital phases, we find a period of 32.688 $\\pm$ 0.002 hours, which we interpret as an orbital period. The folded light curve exhibits an eclipse lasting about 8 hours, suggesting a compact orbit in an nearly edge-on configuration.  The X-ray binary has an average luminosity of $L_X(0.3-8 {\\rm keV}) \\approx 1.3\\times10^{38}$ erg/sec, with only one out of the 25 observations significantly lower in flux than the average light curve.The presence of the eclipse and the $\\sim$ 1.4-day orbital period suggests that this source is an analog of the well studied eclipsing X-ray binary pulsars Her X-1 or LMC X-4.  Combining the Chandra data and the HST ACS/WFC images, we have identified several possible optical counterparts, including an O5-O3 star with V = 25.0.  Follow-up optical monitoring observations should be able to identify the donor and further constrain the orbital properties. ",
        "introduction": "The detections of X-ray binaries have recently begun in galaxies outside the Local Group with the advent of the Chandra X-ray observatory. Owing to its unprecedented sensitivity and spatial resolution, it is possible to resolve individual X-ray point sources in these galaxies. In the past six years, Chandra has studied about one third ($\\sim$70/220) of the galaxies with $D_{25}>1^\\prime$ within 10 Mpc (Liu 2005); these galaxies span the whole spectrum of  galaxy types, including spirals (e.g., M81, Tennant et al.  2001; M51, Terashima \\& Wilson 2004), ellipticals (e.g., NGC1399, Angelini et al.  2001; NGC4697, Sarazin et al. 2000), starburst galaxies (e.g., Antenna and M82, Zezas et al. 2000), and colliding galaxies (NGC7714/7715, Smith et al. 2005). The distributions and luminosity functions of the X-ray binaries in these galaxies, as compared to those in the Milky Way (Grimm et al. 2002) and M31 (Kong et al. 2003), show significant dependence on the galaxy type and star formation properties.    An ultra-deep Chandra ACIS survey (PI: Kuntz) has been executed for the face-on spiral M101 (NGC5457) at a distance of 6.8 Mpc (Freedman et al. 2001). Twenty four Chandra observations were tiled to achieve a full coverage of M101, with a cumulative exposure of one mega second within 4$^\\prime$ of the nucleus. Such a deep exposure can reach a detection limit of 4$\\times$$10^{35}$ erg/sec, only  a factor of two more luminous than current Chandra studies of M31 (Kong et al.  2002). This survey allows detection of most low-mass X-ray binaries, most high-mass X-ray binaries, most young supernova remnants and many middle-aged ones, super-bubbles and giant HII regions. The observations were carried out over the year from January 2004 to January 2005, allowing studies of the variability and transient behaviors of these sources.   Here we report the discovery of a 32.688$\\pm$0.002 hour period for an eclipsing X-ray binary from the M101 ultra-deep survey and an early observation in 2000. The Chandra observations and the search for periodicities are described in section 2. Additional optical observations with HST ACS are described in Section 3. In Section 4, we describe the constraints the 32.69 hour period places on the accretor and donor masses, and discuss the prospects to monitor and identify the donor at optical wavelengths. This X-ray source has been detected previously in ROSAT HRI and PSPC observations and named as M101-X7 (Liu \\& Bregman 2005; H36/P21 in Wang, Immer \\& Pietsch 1999), but the periodicity could not have been detected given the low photon counts and the small number of observations. ",
        "conclusions": "M101 X7 (CXO J140336.0+541925)  was observed in 25 Chandra observations spanning five years, and exhibited significant modulations in twelve observations, with apparent minima present in seven observations. The Lomb periodogram and the phase dispersion minimization were applied in a search for periodicities, revealing a period of 32.688 $\\pm$ 0.002 hours for this source. The average light curve folded with this period shows a deep eclipse with an eclipse full angle of $\\theta_e \\sim 88^\\circ \\pm 11^\\circ$, suggesting that this source is an eclipsing X-ray binary in a nearly edge-on configuration. At a distance of 6.8 Mpc, this is the first eclipsing X-ray binary discovered outside the Local Group. This discovery can only have been made with the superb sensitivity and spatial resolution of the Chandra X-ray Observatory, and with the M101 ultra-deep sequence of observations sampling different orbital phases.   Residual emission was present in the eclipse, indicating that the X-rays are scattered or otherwise processed by, e.g., the strong wind as in the high mass X-ray binaries Cen X-3 (e.g., Schreier et al. 1972) and Vela X-1 (Becker et al. 1978). The presence of strong winds may lead to a slight over-estimate of the eclipse full angle, and its possible variability may have contributed to the mis-alignment of the ingress and egress in some observations. The residual emission from stellar winds will make the spectrum in the eclipse appear softer, while the additional absorption by the blocking secondary will make it harder. A preliminary study, limited by the small number of photons, showed that the hardness of the spectrum did not change significantly with the orbital phases, except for slight hardening present in the eclipse.    Much information can be inferred for this binary from the light curves. The edge-on viewing geometry as suggested by the presence of eclipse in the light curve renders inapplicable the beaming mechanism which enhances the radiation along the polar axis. Thus, its average luminosity of $\\bar L_X$=1.3$\\times10^{38}$ erg/sec constrains the accretor mass through the Eddington limit to be larger than 1 $M_\\odot$. Despite the differences in the average fluxes in individual observations, 24 out of the 25 light curves were roughly consistent with the average light curve, and only one (i.e., 4\\%) was significantly lower in flux than the average light curve. This is distinctly different from Be X-ray binaries with typical periods of 100-300 days, which are in high flux states only briefly when the accretors are near the perigee of their highly elliptical orbits. The relatively shorter period, the wide eclipse, and the nearly uniform distribution of the high flux states suggest that this binary has a much more compact and highly circularized orbit.   The presence of the 32.69 hour period and the X-ray eclipse place strong constraints on the binary properties if we assume the period is the orbital period and the donor fills its Roche lobe. The period $P$ can be computed with Kepler's law $P = 2\\pi a^{3/2}/\\sqrt{G(M_1+M_2)}$, in which the separation $a$ can be related to the Roche lobe radius $R_{cr}$ by $R_{cr} = a\\cdot f(q) = a \\cdot 0.49 q^{2/3}/[0.6 q^{2/3} + ln(1+q^{1/3})]$ (Eggleton 1983).  Here $M_1$ is the accretor mass, $M_2$ is the donor mass, and $q = M_2/M_1$ is the mass ratio. If the donor is filling its Roche lobe, i.e., $R_2 = R_{cr}$, the period will be determined solely by the accretor mass and the donor mass and radius that can be determined empirically by the donor's spectral type. A discrete set of donors of the spectral types as tabulated in {\\it the Astrophysical Quantities} (Allen 2000) are tried to obtain this period for an accretor mass range 1-10$^4$ $M_\\odot$. This period can only be obtained for some dwarf stars; all giants and supergiants give much longer periods.   The donor star can have evolved off the main sequence even if it is not yet a giant/supergiant star. To consider such donors, we first considered main sequence dwarf donors spanning a continuous mass range of 0.5-120 $M_\\odot$, with radii $R_2 = r_0$ interpolated from the empirical mass-radius relation based on dwarf stars of all spectral types. We then introduced a helium core mass $M_c$ for each star, using a standard relationship between core mass and stellar radius: $R_2 = r_0 + 3700 M_c^4 /(1+ M_c^3 + 1.75 M_c^4)$. In the donor mass-accretor mass phase plane the viable binary systems with donors of $M_c=0$ are shown as the thickest line, while the viable binary systems with donors of $M_c>0$ populate the regions below and to the left of this line (Figure 9). The maximum eclipse full angles $\\theta_E$ are computed for the viable binary systems with $sin(\\theta_E/2) = R_2/a$ assuming the inclination angle $i=90^\\circ$, and the observed large $\\theta_e$ ($=\\theta_E sini$) favors those systems with massive donors (Figure 9). Such systems have a total mass $\\le 10^2 M_\\odot$, and the radii of the orbits are less than 1 light minute, rendering barycenter corrections unnecessary.      The accretion rates are calculated for the viable binary systems assuming donors overflow the Roche lobe under conservative mass transfer. Invoking conservation of total angular momentum, the accretion rate can be expressed as $\\dot M_2/M_2 = N/D$. Here the denominator can be expressed as $4[1-(M_c/\\bar M_c)^2] + 2(1-q) - (1+q) dlnf/dlnq$ for conservative mass transfer, and the numerator $N$ can be approximated as ${2.3\\times10^{-7} \\over R_2 } [{(M_1+M_2)^2 \\over M_1} {R_2^4 \\over a^5}]$ when magnetic braking dominates the orbital angular momentum loss. Calculations show that the accretion rate decreases for smaller mass ratio $q$'s and larger $M_c$'s; the contours are plotted in Figure 9 for $\\dot M_2$ = 1, 0.1, 0.01, $10^{-3}$, $10^{-4}$ and $10^{-5}$ in unit of $\\bar L_X/c^2$. For a region with $q>1$ to the right of $\\dot M_2$ = $\\bar L_X/c^2$, $D\\le0$, indicating the breakdown of the adopted expression and very large mass transfer rates that likely lead to common envelopes unless the accreted materials can be dissipated effectively.   Similar orbital periods and eclipses have been detected in two well studied X-ray binaries, i.e., the intermediate-mass X-ray binary Her X-1 and the high-mass X-ray binary LMC X-4 as shown in Figure 9. Her X-1 is an eclipsing binary pulsar with a pulse period of 1.24 seconds and an orbital period of 1.7 days discovered by Uhuru (Tananbaum et al. 1972). Its companion, HZ Her, is a $\\sim2M_\\odot$ A/F star (Bahcall \\& Bahcall 1972) filling its Roche lobe in a compact circular orbit. The X-ray light curve for Her X-1 consists of alternating main on-state, off-states, and short on-states with a super-orbital period of 35 days, which is believed to be caused by the precession of a tilted accretion disk, and has been studied by almost all X-ray missions (cf, Shakura et al.  1998). LMC X-4 was first discovered by Uhuru (Giacconi et al. 1972), and was identified with an O7 III-V star exhibiting a binary period of 1.4 days (Chevalier \\& Ilovaisky 1977). The X-ray pulsations were discovered with a period of 13.5 seconds, the Doppler variations of which revealed the system as a neutron star in circular orbit with a 17 $M_\\odot$ donor nearly filling its Roche lobe (Kelley et al. 1983). The X-ray light curve for LMC X-4 exhibits eclipses, occasional flares, and a 30.5 day super-orbital modulation as prototyped by Her X-1 (Lang et al. 1981).    M101 X7 may be an M101 analog for Her X-1 or LMC X-4 given the similarities in the orbital period and the presence of eclipses. There is, however, no evidence for a super-orbital period of about 30 days for M101 X7, and the pulse period cannot be detected given the low count rate of about 0.005 cnt/sec. Furthermore, its duty cycle for the low flux state is $\\sim$ 4\\%, while the duty cycle for the low flux state (off-state) is $\\ge$50\\% in Her X-1 and LMC X-4. In addition, its eclipse full angle is much larger than those for Her X-1 and LMC X-4 (both $\\sim50^\\circ$), suggesting a higher inclination angle and/or a more massive donor than for Her X-1 or even LMC X-4 (Figure 9). Note, however, that the uncertainties in the duty cycle and eclipse angle are large, given the small number of photons collected and the temporal sampling. This makes direct comparison with Her X-1 and LMC X-4 difficult.    The donor and accretor masses can be further constrained by monitoring and identifying the donor in the optical. Our analyses of the HST ACS/WFC archive observations have resulted in several candidates for the donor, which span a large range of magnitude and mass. The donor mass, thus the primary mass, can be further constrained if the donor can be identified with a single counterpart. For example, were the O5-O3 star A1 (60-120 $M_\\odot$) identified as the secondary, the binary system would be constrained to the lower right corner in Figure 9, with the primary mass lower than 10 $M_\\odot$. This identification can be achieved by photometric monitoring of these counterparts over an orbital period, because the optical light curve should exhibit eclipse and ellipsoidal modulations given the edge-on viewing geometry of the binary. These observations can be carried out with either HST or large ground-based telescopes with adaptive optics such as Gemini-North in the north hemisphere."
    },
    "0608/astro-ph0608448_arXiv.txt": {
        "abstract": "We present five new satellites of the Milky Way discovered in Sloan Digital Sky Survey (SDSS) imaging data, four of which were followed-up with either the Subaru or the Isaac Newton Telescopes.  They include four probable new dwarf galaxies -- one each in the constellations of Coma Berenices, Canes Venatici, Leo and Hercules -- together with one unusually extended globular cluster, Segue 1. We provide distances, absolute magnitudes, half-light radii and color-magnitude diagrams for all five satellites. The morphological features of the color-magnitude diagrams are generally well described by the ridge line of the old, metal-poor globular cluster M92. In the last two years, a total of ten new Milky Way satellites with effective surface brightness $\\mu_v \\gtrsim 28 \\magsq$ have been discovered in SDSS data. They are less luminous, more irregular and appear to be more metal-poor than the previously-known nine Milky Way dwarf spheroidals. The relationship between these objects and other populations is discussed. We note that there is a paucity of objects with half-light radii between $\\sim 40$ pc and $\\sim 100$ pc. We conjecture that this may represent the division between star clusters and dwarf galaxies. ",
        "introduction": "The known satellite galaxies of the Milky Way all lie within $\\sim 300$ kpc and their brightest stars are resolvable from ground-based telescopes. So, it is possible to acquire an enormous wealth of data on their stellar populations, making the satellite galaxies important objects in many fields of astrophysics~\\citep[see e.g.,][]{Do97,Sh03,To04,Pr05}. They have also emerged as a battle-ground in Near-Field Cosmology~\\citep{Fr02}. A fundamental prediction of cold dark matter (CDM) theories is an abundance of substructure in the non-linear regime. As noted by \\citet{Kl99} and \\citet{Mo99}, galaxy assembly in CDM cosmogonies typically yields an order of magnitude more dark haloes than there were known satellites around the Milky Way.  Prior to the Sloan Digital Sky Survey~\\citep[SDSS;][]{Yo00}, there were 9 widely-accepted Milky Way dwarf spheroidals (dSphs), namely Draco, Ursa Minor, Fornax, Carina, Sculptor, Leo I, Leo II, Sextans and Sagittarius.  Seven of the Milky Way dSphs were discovered by eye using photographic plates. The eighth, Sextans, was found by \\citet{Ir90} as part of a search of automated scans of photographic plates, whilst the ninth, Sagittarius, was first identified kinematically from radial velocity surveys of the Galactic bulge~\\citep{Ib95}.  The number of known Milky Way dSph satellites had been increasing at a rate of one or two per decade before the advent of SDSS.  The impact of SDSS has been dramatic. Four new Milky Way dSph satellites have been discovered in SDSS data in quick succession: Ursa Major~\\citep{Wi05a}, Canes Venatici~\\citep{Zu06a}, Bo{\\\"o}tes~\\citep{Be06b} and Ursa Major II~\\citep{Zu06c,Gr06}, together with what appears to be an unusually extended globular cluster~\\citep{Wi05b}. None are apparent in SDSS images, but all are very clearly identifiable as overdensities of resolved stellar objects.  This paper presents a further five new satellites found in SDSS data, one each in the constellations of Coma Berenices, Canes Venatici, Hercules, and two in Leo. We have confirmed four of these discoveries with follow-up photometry on the Subaru telescope on Mauna Kea and the Isaac Newton Telescope on La Palma. This brings the total number of Milky Way companions found with SDSS data to ten, eight of them probable dSphs.  This roughly doubles the number known prior to SDSS.  They have eluded previous discovery because they are all of low surface brightness ($\\mu_v \\gtrsim 28 \\magsq$).  In fact, recent years have seen the discovery of a number of objects that blur the hitherto clear distinction between star clusters and dwarf galaxies. These include the ultra compact dwarf galaxies in the Fornax cluster~\\citep[e.g.,][]{Hi99,Dr00,Mi02}, the globular clusters with unusually large half-light radii in M31~\\citep{Hu05} and the faint dSphs around M31~\\citep{Zu06b,Ma06}. The 10 new SDSS discoveries all lie in this poorly charted territory, where -- in the absence of kinematic data -- the distinction between star clusters and dwarf galaxies is hazy.  The paper is organised as follows: \\S 2 provides a summary of the SDSS and follow-up photometry on our 5 new discoveries, together with a table of their properties.  \\S3 reviews the relationship between globular clusters and dwarf galaxies in the light of our new data, and considers the implications of our discoveries for Near-Field Cosmology. \\S 4 summarizes our conclusions.  \\begin{figure*}[t] \\begin{center} \\includegraphics[height=17cm]{f1.sm.eps} \\caption{Discovery panels for the 5 new satellites. The first column is a cut-out of the SDSS, with a box showing the location of the Subaru field ($34^\\prime \\times 27^\\prime$) or INT field ($34^\\prime \\times 34^\\prime$) and a circle marking the central part of the object. The second column shows the pixellated stellar density. The pixels are $4^\\prime$ on each side. For each object, 3 circles are shown of radii $r_1, r_2$ and $r_3$. The CMD of stars lying within a circle of radius $r_1$ is given in the third column. The CMD of stars lying in the annulus defined by the outer radii ($r_2$ and $r_3$) is given in the fourth column. [$r_1,r_2,r_3$ for Coma are $0.15^\\circ,0.4^\\circ, 0.43^\\circ$, for CVn II ($0.12^\\circ, 0.3^\\circ, 0.32^\\circ$), for Segue 1 ($0.12^\\circ, 0.5^\\circ, 0.51^\\circ$), for Her ($0.1^\\circ, 0.3^\\circ, 0.32^\\circ$) and for Leo IV ($0.1^\\circ, 0.3^\\circ, 0.32^\\circ$]. \\label{fig:objsa}} \\end{center} \\end{figure*} \\begin{figure*}[t] \\begin{center} \\includegraphics[height=4.5cm]{f2.sm.eps} \\caption{CMDs of the central parts of Com, CVn II, Segue 1 and Her from the Subaru/INT follow-up data. \\label{fig:extra}} \\end{center} \\end{figure*} \\begin{figure*}[t] \\begin{center} \\includegraphics[height=9cm]{f3.eps} \\caption{Differential Hess diagrams using SDSS (upper panels) and Subaru or INT (lower) data for Com, CVn II, Segue 1 and Her. In each case, the normalized Hess diagram constructed with stars selected within $r_1$ is subtracted from the normalized Hess diagram constructed with stars selected between $r_2$ and $r_3$. \\label{fig:objsb}} \\end{center} \\end{figure*} \\begin{figure*}[t] \\begin{center} \\includegraphics[height=9cm]{f4.eps} \\caption{Contour levels on the differential Hess diagrams using SDSS (upper panels) and Subaru or INT (lower) data for Com, CVn II, Segue 1 and Her. The ridgeline of M92 (M13) is overlaid as a solid (dotted) line, using the data of \\citet{Cl05}. \\label{fig:objsc}} \\end{center} \\end{figure*} \\begin{figure*}[t] \\begin{center} \\includegraphics[height=7.5cm]{f5.eps} \\caption{Isodensity contours for Com, CVn II, Segue 1 and Her I. Membership is determined using a mask constructed from the M92 ridgeline.  The top panels show CMD-selected stars with $18 < i < 22.5$. There are $30 \\times 30$ pixels, smoothed with a Gaussian with FWHM of 3 pixels.  Contour levels are 2, 3, 5, 7, 10, 15 $\\sigma$ above the background. The bottom panels show the central parts of the objects in Subaru/INT data.  There are $30 \\times 30$ pixels, smoothed with a Gaussian with FWHM of 2.2 pixels. Contour levels are 2, 3, 4, 5, 7, 10, 15 $\\sigma$ above the background. \\label{fig:objsd}} \\end{center} \\end{figure*} \\begin{figure*}[t] \\begin{center} \\includegraphics[height=5cm]{f6.eps} \\caption{Left: Differential Hess diagram for Leo IV, together with the color-magnitude box used to select BHB candidate stars. Middle: Contours of the differential Hess diagram, with overplotted M92 ridgeline and mask used to select members. Right: Isodensity contours of Leo IV, together with locations of BHB candidate stars. \\label{fig:leo4}} \\end{center} \\end{figure*}  \\begin{deluxetable*}{lccccc} \\tablecaption{Properties of the New Milky Way Satellites \\label{tbl:pars}} \\tablewidth{0pt} \\tablehead{ \\colhead{Parameter\\tablenotemark{a}} &  \\colhead{Coma Berenices} & \\colhead{Canes Venatici II} & \\colhead{Segue 1} &  \\colhead{Hercules} & \\colhead{Leo IV} {~~~ } } \\startdata  Coordinates (J2000) & 12:26:59 +23:54:15 &  12:57:10 +34:19:15 &  10:07:04 +16:04:55 &  16:31:02 +12:47:30 & 11:32:57 -00:32 00 \\\\ Galactic $(\\ell,b)$ & $241.9^\\circ$, $83.6^\\circ$  & $113.6^\\circ$, $82.7^\\circ$ & $220.5^\\circ$, $50.4^\\circ$ & $28.7^\\circ$, $36.9^\\circ$ & $265.4^\\circ$, $56.5^\\circ$ \\\\ Position Angle & $120^\\circ$ & $0^\\circ$ & $60^\\circ$ & $125^\\circ$  & $355^\\circ$ \\\\ Ellipticity & $0.5$ & $0.3$ & $0.3$ & $0.5$ & $0.25$ \\\\ $r_h$ (Plummer) & $5\\farcm0$ & $3\\farcm0$ & $4\\farcm5$ & $8\\farcm0$ & $3\\farcm3$\\\\   $r_h$ (Exponential) & $5\\farcm9$ & $3\\farcm3$ & $4\\farcm6$ & $8\\farcm4$ & $3\\farcm4$\\\\ V$_{\\rm tot}$ & $14\\fm5\\pm0\\fm5$ & $15\\fm1\\pm0\\fm5$ & $13\\fm8\\pm0\\fm5$ & $14\\fm7\\pm0\\fm5$ & $15\\fm9\\pm0\\fm5$ \\\\ (m$-$M)$_0$ & $18\\fm2\\pm0\\fm2$ & $20\\fm9\\pm0\\fm2$ &$16\\fm8\\pm0\\fm2$ &  $20\\fm7\\pm0\\fm2$ & $21\\fm0\\pm0\\fm2$ \\\\ Heliocentric distance & $44^{+4}_{-4}$ kpc & $150^{+15}_{-13}$ kpc  & $23^{+2}_{-2}$ kpc & $140^{+13}_{-12}$ kpc & $160^{+15}_{-14}$ kpc\\\\ M$_{\\rm tot,V}$ & $-3\\fm7\\pm0\\fm6$ & $-4\\fm8\\pm0\\fm6$ & $-3\\fm0\\pm0\\fm6$ & $-6\\fm0\\pm0\\fm6$ & $-5\\fm1\\pm0\\fm6$ \\enddata \\tablenotetext{a}{Integrated magnitudes are corrected for the Galactic foreground reddening reported by \\citet{Sc98}} \\label{tab:objs} \\end{deluxetable*} \\begin{figure*} \\begin{center} \\includegraphics[height=7cm]{f7.sm.eps} \\caption{The locations of Milky Way satellites in Galactic coordinates.  Filled circles are satellites discovered by SDSS, unfilled circles are previously known Milky Way dSphs. The light grey shows the area of sky covered by the Sloan survey and its extensions to date.  The dashed and dotted lines show the orbital planes of the Sagittarius and Orphan Streams, respectively, taken from \\citet{Fe06a} and \\citet{Fe06b}.} \\label{fig:dwarfslb} \\end{center} \\end{figure*} ",
        "conclusions": "In this paper, we announced the discovery of five new satellites of the Milky Way. One is a probable new globular cluster, which has been named Segue 1 after the survey in which it was found. It is in the Sagittarius stream and was possibly stripped off the Sagittarius progenitor.  The remaining four are probably new dwarf galaxies, which have been named according to their constellations as Coma Berenices, Canes Venatici II, Leo IV and Hercules. We have presented SDSS and deeper Subaru/INT photometry (where available) on these objects. We provide color-magnitude diagrams, distances, absolute magnitudes and half-light radii for all the satellites.  Taken together with the earlier announcements of Ursa Major I, Willman 1, Canes Venatici I, Bootes and Ursa Major II, ten new Milky Way satellites have been discovered in SDSS data in very rapid succession. This abundance of discoveries is occurring because the survey is probing down to hitherto uncharted surface brightnesses. All the SDSS discoveries are effective surface brightness fainter than $\\mu_V =27 \\magsq$. The obvious conclusion is that there are more low surface brightness Milky Way satellites waiting to be discovered.  The SDSS discoveries occupy a distinct region in the plane of absolute magnitude versus half-light radius. They are typically fainter, more metal-poor and more irregular than the previously-known Milky Way dwarf spheroidals (dSphs).  They are larger, and somewhat less luminous, than typical Galactic globular clusters.  Taking the known globular clusters, the previously known Milky Way satellites and the SDSS discoveries, there is still a scarcity of objects with half-light radii between $\\sim 40$ pc and $\\sim 100$ pc. This may represent the division between star clusters and dwarf galaxies.  The SDSS discoveries could have a bearing on the ``missing satellite'' problem.  Preliminary indications from studies of UMa I~\\citep{Kl05} and Boo~\\citep{Mu06} suggest that these objects may be dark matter dominated.  It seems possible that a population of ultra-faint, dark dwarf galaxies really does surround the Milky Way. However, it is not yet clear that these are the ``missing satellites'' predicted by the simulations of \\citet{Kl99} and \\citet{Mo99}. The match of the data to CDM halos should be carried out in the plane of cumulative number versus halo mass or circular velocity.  Kroupa's (1997) study of the evolution of tidal dwarf galaxies reproduces some of the properties of the new satellites. This opens up the possibility that some of these objects may be tidal dwarf galaxies, or shreds from the violent building phase of the Milky Way. In this case, the satellites will not have substantial dark matter. Kinematic data are now urgently needed to confirm whether or not these objects are dark matter dominated."
    },
    "0608/astro-ph0608481_arXiv.txt": {
        "abstract": "Multifrequency, high-resolution radio observations of the quasar 3C254 using MERLIN are presented. The quasar has a highly asymmetric radio structure, with the eastern component of the double-lobed structure being much closer to the nucleus and significantly less polarized than the western one. However, the two lobes are more symmetric in their total flux densities. The observations show the detailed structure of the hotspots which are very different on opposite sides of the radio core, reveal no radio jet and suggest that the oppositely-directed jets may be intrinsically asymmetric. ",
        "introduction": "The majority of high-luminosity extragalactic radio sources selected at a low frequency tend to be symmetric in the brightness and location of the components relative to the nucleus of the parent galaxy. However, a small but significant fraction are highly asymmetric.  The asymmetries of the lobes of emission are important because they could provide useful insights into the environments of sources (e.g. Pedelty et al. 1989; McCarthy, van Breugel \\& Kapahi 1991), the interaction of jets with external clouds or galaxies, and also provide tests of the orientation-based unified scheme for radio galaxies and quasars (Barthel 1989). For example, the structure of the highly-asymmetric double-lobed source B0500+630 appears to be inconsistent with the unified scheme and suggests that there is an intrinsic asymmetry in the oppositely-directed jets from the nucleus (Saikia et al. 1996). The most extreme form of asymmetry is when the source is completely one-sided with radio emission on only one side of the nucleus. Although most of these objects are core-dominated and their apparent asymmetry is likely to be due to bulk relativisitic motions in the extended lobes of emission, there does appear to be a number of weak-cored one-sided sources which are difficult to reconcile with the simple relativistic beaming scenario (Saikia et al. 1990).  As part of a study of highly asymmetric radio sources, the quasar 3C254 has been observed at radio wavelengths with high resolution using MERLIN. The results of these observations are presented in this paper. 3C254 is identified with a quasar at a redshift of 0.73612$\\pm$0.00160 (NASA Extragalactic Database) so that 1 arcsec corresponds to 7.290 kpc (H$_o$=71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m$=0.27, $\\Omega_\\Lambda$=0.73, Spergel et al. 2003).  The radio source has an overall angular extent of 13.5 arcsec (98 kpc) with the eastern component being much closer to the nucleus. The ratio of the separations of the outer hotspots, r$_D$, defined to be $>$1, is 7.1 while the flux density ratio of the corresponding lobes from the MERLIN {\\it L}-band images presented here is $\\sim$0.55. Earlier Very Large Array (VLA) and MERLIN observations with lower resolution have shown the highly asymmetric structure and a bridge of emission connecting the two lobes (e.g. Owen \\& Puschell 1984; Liu \\& Pooley 1991; Liu, Pooley \\& Riley 1992; Reid et al. 1995).  The observations by Liu \\& Pooley have also shown that the eastern lobe is more depolarized, and has a steeper spectrum, consistent with the relationship reported by them that the more depolarized lobe tends to have a steeper radio spectrum. This relationship is significantly stronger for smaller sources, but is similar for both radio galaxies and quasars suggesting that, in addition to Doppler effects, there are intrinsic differences between the lobes on opposite sides (cf. Ishwara-Chandra et al. 2001).  The large separation ratio, the distortion in the eastern lobe and the stronger depolarization all suggest that the jet on the eastern side is colliding with a dense gas cloud which is influencing the properties and appearance of the radio source.  Evidence of the gas clouds has been reported by Forbes et al. (1990)  who found this steep-spectrum quasar to be embedded in a spectacular, extended, emission-line region emitting strongly in [O{\\sc ii}] $\\lambda$3727 and [O{\\sc iii}] $\\lambda$5007. Forbes et al. have also suggested that the emitting gas is at high pressure, consistent with confinement by a hot intracluster medium. Subsequent observations have shown that the optical continuum and line emission are extended along the radio axis (Bremer 1997; Crawford \\& Vanderriest 1997), illustrating the so-called alignment effect which is seen in powerful radio galaxies. Bremer also shows that part of the extended emission line region which is close to the eastern lobe has a blueshift of $\\sim$650 km s$^{-1}$ relative to the nuclear emission along a position angle of 105$^\\circ$. There is also an overdensity of galaxies within approximately 10 arcsec of the quasar, suggesting that 3C254 lies at the centre of a compact group of galaxies. High-resolution X-ray observations with the Chandra Observatory show X-ray emission from the quasar and the western hotspot, and demonstrate that the source does not lie in a hot, massive cluster (Donahue, Daly \\& Horner 2003).  \\begin{figure*} \\begin{center} \\vbox{ \\hspace{0.0cm} \\psfig{file=3C254-1539-I.PS,width=6.5in,angle=-90} } \\end{center} \\caption{The MERLIN images of 3C254 at 1539 MHz with an angular resolution of $\\sim$0.17 arcsec. Peak brightness: 142 mJy/beam; Contours: 0.4$\\times$($-$1, 1, 2, 4, 8, 16 $\\ldots$) mJy/beam. } \\end{figure*}  In this paper, high-resolution MERLIN observations at {\\it L}- and {\\it C}-bands are presented. These observations were made to examine the structure and asymmetry of the hotspots on opposite sides and determine any possible intrinsic asymmetry in the oppositely-directed jets. The observations and analyses of the data are described in Section 2, and the results are presented in Section 3. Possible explanations for the structure are discussed in Section 4, while the conclusions are summarised in Section 5. ",
        "conclusions": "The radio galaxy, 3C254, is one of the most asymmetric radio sources in terms of the location of the outer lobes of radio emission relative to the nucleus of the host galaxy. MERLIN observations of the source have been presented at {\\it L}- and {\\it C}-bands with resolutions ranging from 0.18 (1.3 kpc) to 0.057 arcsec (0.4 kpc) respectively. These sensitive, high-resolution observations do not show any evidence of a jet on either the western or eastern side of the source, making it difficult to identify the approaching side on the basis of jet sidedness. Identifying the eastern side as the approaching one on the basis of the blue-shifted emission lines being pushed outwards by interaction with the radio jet (Bremer 1997), it appears that while the large-scale asymmetries may be caused by interaction with an asymmetric external environment, the small-scale asymmetries in the hotspots are probably the result of an intrinsic asymmetry in the jets on opposite sides. The usual assumption of oppositely-directed symmetric jets does not appear to be universally applicable, but it may be possible to establish this in only the most asymmetric of sources."
    },
    "0608/astro-ph0608162_arXiv.txt": {
        "abstract": "We present new, sensitive, near-infrared images of the Class I protostar, Elias 29, in the Ophiuchus cloud core.   To explore the relationship between the infall envelope and the outflow, narrowband H$_2$ 1$-$0 S(1), Br$\\gamma$, and $K_{\\rm{cont}}$ filters were used to image the source with the Wide-Field Infrared Camera on the Hale 5.0-meter telescope and with Persson's Auxiliary Nasmyth Infrared Camera on the Baade 6.5-meter telescope. The source appears as a bipolar, scattered light nebula, with a wide opening angle in all filters, as is typical for late-stage protostars. However, the pure H$_2$ emission-line images point to the presence of a heretofore undetected precessing jet.  It is argued that high velocity, narrow, precessing jets provide the mechanism for creating the observed wide-angled outflow cavity in this source. ",
        "introduction": "Powerful, bipolar supersonic outflows are a hallmark of the earliest protostellar stages. Molecular outflows generally exhibit wide opening angles and velocities of $\\sim$ 10--30 km s$^{-1}$. Optical jets powered by protostars \\textit{}% exhibit small opening angles and velocities of $\\sim$ 100-400 km s$^{-1}$ \\cite{mun90,fri05}.   Due to the disparate spatial and velocity distributions of the molecular and optical outflow components, a two-wind structure has often been invoked to explain the instances in which both outflow components are observed to be powered by a single source \\cite{sto88,dav02,lis05}. However, an alternative scenario, in which optical jets and molecular outflows result from a single, collimated flow is gaining wide acceptance \\cite{har94, san99, wol00, yun01}.  An outstanding, unsolved problem with jet-driven outflow models is how jets produce the wide-angle outflow cavities observed in both protostellar envelopes and in molecular outflows. A proposed solution is that of jets changing direction \\cite{mas93}, for which there is now growing evidence \\cite{brp96, rei97, bar98, arc05}. Nevertheless, the spatial relationship between the outflow cavity and the currently active wandering jet has not yet been directly observed within the boundaries of a protostellar envelope. We use sensitive, narrowband near-infrared imaging to study this relationship: the scattered light emission from the protostellar envelope is imaged with a narrowband continuum filter, whereas the currently active jet regions within the infall envelope are traced in continuum-subtracted, narrowband H$_2$ images, sensitive to shock-excited emission. Based on such observations of the Class I protostar, Elias 29$=$WL15 ($\\alpha_{2000}=16^h27^m09.43^s$, $\\delta_{2000}=-24^{\\circ}37^{\\prime}18\\farcs7$) in the nearby (d$=$125 pc) $\\rho$ Ophiuchi molecular cloud, we report the first detection of a precessing jet carving out a protostellar envelope's cavity. ",
        "conclusions": "The reduced WIRC and PANIC images of Elias 29 appear essentially identical. Only the PANIC images are presented here, since these were obtained in better seeing and have higher signal-to-noise.  Figure 1 shows our results, all presented at a linear greyscale stretch to emphasize the presence of the pure emission features, 3a and 3b in Figure 1a, and their absence from Figure 1b. Figure 1a shows the appearance of Elias 29 through the narrowband H$_2$  filter. Figure 1b shows Elias 29 through the surrogate ``narrowband continuum'' (Br $\\gamma$) filter. Since the NIR continuum emission overwhelms the strength of the pure H$_2$ line emission for this object, Elias 29 looks similar in both filters. Figure 1c shows the continuum-subtracted, pure H$_2$ emission-line image of Elias 29. Saturation effects from the extreme brightness of Elias 29 prevent correct image subtraction within the central $\\sim$ 5\\farcs5 radius region, resulting in the black and white circular artifacts at the center of Figure 1c. Outside this area, five regions of pure H$_2$ line emission are discovered, having been missed by previous investigators \\cite{gom03, kha04}. The appearance of Elias 29 in Figures 1a \\& 1b is consistent with scattered light models of Class I protostars consisting of an infall envelope containing a bipolar cavity \\cite{whi03}. The boundary of the wide-angled cavity walls within the protostellar infall envelope is indicated by the lowest contour levels overplotted on Figures 1b \\& 1c. The lack of a narrow ``waist'' or distinctive hourglass shape at this angular resolution is consistent with a disk of small spatial extent (B. Whitney, priv. comm.). Elias 29 appears strikingly different in the continuum-subtracted, pure H$_2$ emission line image than in the narrowband filter images. The pure H$_2$ emission objects, labelled 1, 2a, 2b, 3a, and 3b, appear at the same location, and have similar morphologies in both the PANIC and the WIRC (not shown here) continuum-subtracted images.  Taken together, the H$_2$ emission knots form a sinuous structure with S-shaped point-symmetry about the center.  Table 1 lists the observed characteristics of each H$_2$ feature. The first column lists the feature designations from Figure 1c. The second and third columns list the coordinates of each H$_2$ emission region, corresponding to the WIRC pixel with the highest counts in each object. The coordinates are derived from the WIRC data, since there were not enough sources in the smaller PANIC field to provide a plate solution for absolute astrometry. The final, continuum-subtracted PANIC image was used to determine the signal to noise ratio of each emission feature, listed in the last column of Table 1, since the seeing and signal to noise were better in the PANIC than in the WIRC dataset. A rectangular aperture (of dimensions given in the fourth column of Table 1) was used to determine the average counts in each feature. Several background regions were used to calculate the average background noise level and its standard deviation.  The backgound level was subtracted from the average counts in the rectangular aperture containing each H$_2$ emission region, and the result was divided by the standard deviation of the background to give the SNR. Only features 3a and 3b have $\\sigma$ $<$ 3. Nevertheless, these objects are real, since they appear at identical locations and have similar morphologies in both datasets (WIRC  and PANIC).  H$_2$ emission associated with molecular outflows generally arises from shock interactions of the unseen, fast jet/wind component with the ambient medium.  Elias 29 is known to drive a CO outflow \\cite{bon96, sek97}, with velocities as high as $\\sim$ 80 km s$^{-1}$ \\cite{boo02}, although its morphology is unknown at the scales of the images presented here. In some Class I protostellar envelopes, both spatially diffuse H$_2$ emission and compact H$_2$ emission knots have been identified in the same source \\cite{yam92, dav02}.  In such cases, the diffuse component is found to trace the infrared continuum reflection nebula, whereas knotty, compact emission traces the fast jets observed at optical or cm-continuum wavelengths. By analogy, the knotty H$_2$ emission features identified in Figure 1c are likely to be associated with the jets driven by Elias 29.  Since the cooling time of shocked 2.12 $\\mu$m H$_2$ emission is of order $\\sim$ a few years \\cite{shu82,all93}, the newly identified features in Figure 1c are tracing recent shock activity. The three H$_2$ emission features in Figure 1c: 1, 2a, and 2b, describe an S-shaped, point-symmetry, about the center, consistent with a precessing jet. Features 3a and 3b are extended and diffuse, close to the inferred cavity wall within the southern hemisphere of the protostellar envelope of Elias 29. More quantitative results on the precessing jet, such as its precession period and initial opening angle await HST NICMOS polarimetry. Ground-based, AO-assisted  integral field spectroscopy of the H$_2$ line would establish the velocity field.  Hydrodynamic simulations of molecular outflows driven by pulsed, precessing protostellar jets have become available only recently \\cite{lim01,ros04, smi05}. One stated aim of these studies is to model the outflows to infer properties of the (unseen) driving jets, such as jet power and precession.  Currently, only a limited set of models have been calculated with accompanying simulated H$_2$ images.  Assumed jet radii are 1.7 $\\times$ 10$^{15}$ cm, corresponding to $\\sim1\\farcs0$ at the distance to Elias 29 (so, well beyond the disk region where jets may be generated). The pulsation is input as a 30\\% velocity modulation with a 60 yr period into the models. The timespan over which the models are calculated is short (t $<$ 500 yr), but much longer than the dynamical time of a few decades for a jet travelling a conservative $\\sim$ 100 km s$^{-1}$ to reach the H$_2$ Features 1 and 2a in Elias 29. This, compared with the very fast cooling time of H$_2$, validates comparison of the Elias 29 pure emission-line H$_2$ image (Figure 1c) with simulated H$_2$ images.  Models of molecular jets with precession angles of 5$^{\\circ}$, 10$^{\\circ}$, and 20$^{\\circ}$ have been calculated, although the precession angle of the Elias 29 jet may well be even wider, based on the projected half-angles formed by Features 1 and 2a with the central symmetry axis of the observed reflection nebulosity.  Simulated H$_2$ images from two general classes of precessing jet models have been published: Jets with assumed precession periods of 50 yr, termed ``fast'' precessing, and jets with precession periods of 400 yr, termed ``slow'' precessing \\cite{ros04,smi05}.  The terms ``fast'' and ``slow'' here are relative to the flow evolution time of 500 yr. Comparison of Figure 1c with synthetic H$_2$ $v=1-0$ S(1) model images from such simulations leads to the conclusion that this jet exhibits slow, rather than fast, precession, that is to say, qualitatively, Figure 1c resembles the H$_2$ 1$\\rightarrow$0 morphologies presented in Figures 6 \\& 7, rather than those of Figure 5, of Smith \\& Rosen (2005). For slow precession, the simulations show ordered chains of bow shocks and meandering streamers of H$_2$, which contrast with the chaotic H$_2$ structures produced by jets in rapid precession. The production of specific simulated H$_2$ images for direct comparison with Figure 1c is beyond the scope of this paper, but may be attempted in the future.  Images of both scattered light cavities and H$_2$ jet emission within them exist for only a handful of Class I objects \\cite{dav02}.  For all of these, the jets lie well within the outflow cavity walls, close to, if not precisely along, the symmetry axis of each outflow cavity. By contrast, the jet of Elias 29, as traced by H$_2$, is found close to, and, in one case (Feature 2b) along, the cavity walls (Fig. 1c).  Examples of other S-shaped, precessing outflows are known (Gueth et al. 1996; Schwartz \\& Greene 1999; Hodapp et al. 2005; Sahai et al. 2005). However, no precessing protostellar jets have been imaged at their base within their scattered light cavities, either because they are too embedded to be detected in the NIR (L1157, HH211-mm), or because of lack of narrowband continuum data excluding the H$_2$ emission line for proper image subtraction(IRAS 03256$+$3055).Binarity has inevitably been invoked to account for the presence of precessing jets \\cite{ter99}.  However, Elias 29 is known to be single \\cite{sim95, ghe93}. At scales between $\\sim 1-60$AU, lunar occultation observations show that 90\\% of the K-band flux from Elias 29 originates in a component $\\sim$ 7 mas diameter ($\\sim$ 0.8 AU) with the remaining 10\\% of the K-band flux originating in a much larger, 415 mas diameter, diffuse component \\cite{sim87}. It is then inferred that within $\\approx$ 0.8 AU, Elias 29 is a single object. Alternative scenarios leading to jet precession that may be operative in Elias 29 might be disk oscillations associated with a propeller-driven outflow, recently discovered in MHD simulations \\cite{rom06} or warping of the accretion disk caused by the magnetically driven jets \\cite{lai03}. Future high-resolution integral field spectroscopic observations of both the jet and the disk components could test the latter model, which predicts the disk and jets to precess in opposite senses.  We thank our very professional night assistants, Jean Mueller at Palomar Observatory and Hernan Nu\\~nez at Las Campanas Observatory. Portions of this research were funded by NSF AST-0206146 to M.B. and by the NASA LTSA Program (NAG5-8933).  NASA/LTSA Award 399-30-61-00-00 to R.S. is gratefully acknowledged.   AJW acknowledges support from NASA Space Science grant NAG5-3175.  \\clearpage"
    },
    "0608/astro-ph0608472.txt": {
        "abstract": "We measure the relation between galaxy luminosity and disk circular velocity (the Tully-Fisher, or TF, relation), in the $g$, $r$, $i$, and $z$-bands, for a broadly selected sample of galaxies from the Sloan Digital Sky Survey, with the goal of providing well defined observational constraints for theoretical models of galaxy formation.  The input sample of 234 galaxies has a roughly flat distribution of absolute magnitudes in the range $-18.5$ $>$ $\\Mr$ $>$ $-22$, and our only morphological selection is an isophotal axis-ratio cut $b/a$ $<$ 0.6 to allow accurate inclination corrections.  Long-slit spectroscopy from the Calar Alto and MDM observatories yields usable \\halpha\\ rotation curves for 162 galaxies (69\\%), with a representative color and morphology distribution.  We define circular velocities $\\Veight$\\ by evaluating the rotation curve at the radius containing 80\\% of the $i$-band light. Observational errors, including estimated distance errors due to peculiar velocities, are small compared to the intrinsic scatter of the TF relation.  The slope of the forward TF relation steepens from $-5.5$ $\\pm$ $0.2\\,\\hbox{mag}/\\hbox{log}_{10}\\kms$ in the $g$-band to $-6.6$ $\\pm$ $0.2\\,\\hbox{mag}/\\hbox{log}_{10}\\kms$ in the $z$-band.  The intrinsic scatter is $\\sigma$ $\\approx$ 0.4 mag in all bands, and residuals from either the forward or inverse relations have an approximately Gaussian distribution.  We discuss how Malmquist-type biases may affect the observed slope, intercept, and scatter.  The scatter is not dominated by rare outliers or by any particular class of galaxies, though it drops slightly, to $\\sigma$ $\\approx$ 0.36 mag, if we restrict the sample to nearly bulgeless systems.  Correlations of TF residuals with other galaxy properties are weak:  bluer galaxies are significantly brighter than average in the $g$-band TF relation but only marginally brighter in the $i$-band; more concentrated (earlier type) galaxies are slightly fainter than average; and the TF residual is virtually independent of half-light radius, contrary to the trend expected for gravitationally dominant disks.  The observed residual correlations do not account for most of the intrinsic scatter, implying that this scatter is instead driven largely by variations in the ratio of dark to luminous matter within the disk galaxy population. ",
        "introduction": "The observed correlation between luminosity and disk rotation speed \\citep{tully77} is one of the fundamental empirical clues to the physics of galaxy formation, in particular to the relation between dark matter halos and their luminous baryonic components. The Tully-Fisher (hereafter TF) relation has been widely exploited as a distance indicator, in studies of the cosmic distance scale (e.g., \\citealt{tully77,aar86,tul00,fre01}) and the large scale peculiar velocity field (e.g., \\citealt{wil90,mat92,wil97b,cou00}).  The ambitious surveys constructed for such studies have usually focused on a relatively narrow range of galaxy types, typically undisturbed late-type spirals, with the goal of obtaining a tight relation that can yield precise distances. However, the small scatter and the measured parameters of the TF relation are adopted as key constraints on galaxy formation theories (e.g., \\citealt{kau93,col94,mo98,som99,nav00}), even though these may not yet have the detail to predict the precise types of model galaxies.  The goal of this paper is to measure the TF relation for a broadly selected sample of galaxies, with a focus on quantifying rather than minimizing the intrinsic scatter and on measuring the correlation of TF residuals with other galaxy properties.  Our sample of galaxies is drawn from the Sloan Digital Sky Survey (SDSS, \\citealt{yor00}), an imaging and spectroscopic survey of the North Galactic Cap and selected regions of the South Galactic Cap.  The SDSS galaxy spectra are obtained through a $3''$ diameter fiber, so they do not yield reliable estimates of rotation velocities for galaxies that are spatially well resolved.  We have therefore obtained long-slit \\halpha\\ rotation curve data for a sample of 234 SDSS galaxies using the Calar Alto 3.5-m telescope and the MDM 2.4-m telescope.  The Calar Alto observations of 189 galaxies constituted a substantial portion of the Calar Alto Key Project contribution to the SDSS.  Selection from the SDSS redshift survey allows us to define our sample based on absolute magnitude rather than apparent magnitude, while working in a redshift range $5000\\kms \\leq v \\leq 15000\\kms$ where peculiar velocities add little uncertainty to individual galaxy distances.  Equally important, the SDSS provides high quality, 5-band imaging for all of the program galaxies, and the connection to the much larger SDSS database allows statistical application of our results to, e.g., infer the distributions of galaxy potential well depths or angular momenta in the local universe.  As discussed in \\S 2, our selection criteria produce a roughly flat distribution in $r$-band absolute magnitude over the range $-18.5$ $>$ $\\Mr$ $>$ $-22$.\\footnote{Throughout the paper, we adopt $h$ $\\equiv$ $H_0$ / 100 $\\kms$ Mpc$^{-1}$ $=$ 0.7 and quote absolute magnitudes for this value.  One should add 5\\,log($h$/0.7) for other values of $h$.  All logarithms are base 10.}  We select galaxies with isophotal axis ratios $b/a \\leq 0.6$ so that we can make accurate inclination corrections to rotation speeds, but we impose no other morphological selection criteria.  We obtain usable \\halpha\\ rotation curves for 170 of our 234 selected targets (73\\%).  The axis ratio cut slightly reduces the representation of early-type galaxies, and the early-type galaxies that pass this cut are slightly less likely to yield usable \\halpha\\ rotation curves.  Nonetheless, relative to most previous large TF samples, our sample is more representative of the range of disk galaxy types.  In addition, the redshift range and accurate photometry make our typical observational errors smaller than the intrinsic TF scatter, which is essential for an accurate estimate of this scatter and a full understanding of residual correlations.  In brief, this is a TF sample designed for studies of galaxy formation, not for measurements of the Hubble constant or peculiar velocities.  We presented and discussed the scaling relations for a disk-dominated subset of this sample in (\\citealt{piz05}, hereafter P05). The selection of inclined galaxies ensures a small uncertainty on the inclination corrected velocity width.  There are a variety of ways to define a galaxy's rotation speed from optical or 21cm data, and in general they yield similar but not identical results for the correlation between luminosity and rotation speed (see  \\cite{ver01} for a careful investigation of this issue).  As discussed in \\S 4, our measure of disk rotation speed is based on the value of an arc-tangent fit to the rotation curve (\\citealt{cou97}, hereafter C97) evaluated at a position containing 80\\% of the total $i$-band flux.  The work of \\cite{tully77} used 21cm line widths rather than optical rotation curves, so strictly speaking our analysis is not the ``Tully-Fisher relation'', but we will follow common practice in using this term for the more general correlation between luminosity and gas rotation speed.  The observational study most similar to our own, in terms of a broad sample selection, is that of \\cite{kan02}, hereafter K02, who measured TF relations and residual correlations in $U$, $B$, and $R$-bands for a sample of 68 galaxies with morphological types Sa-Sd from the Nearby Field Galaxy Sample \\citep{jan00}.  Our sample is larger by a factor of $\\sim$ 2.5, and the luminosity distribution is rather different:  its absolute magnitude histogram is roughly flat in the range $-18.5$ $>$ $\\Mr$ $>$ $-22$, with somewhat higher representation on the bright end, while K02's histogram peaks at $M_B$ $=$ $-18$ and declines at brighter magnitudes.  The selection of our sample at $v$ $\\geq$ 5000 $\\kms$ and the use of SDSS surface photometry to determine disk inclinations makes our observational errors much smaller than those of K02, giving us a better handle on the TF intrinsic scatter.  However, unlike K02 we do not have literature HI data for many galaxies in our sample, nor do we extend our sample to faint luminosities.  In comparisons to the voluminous observational TF literature, we will mostly focus on K02 and on the work of C97, who analyzed large samples of disk galaxy optical rotation curves using methods similar to those adopted here, and Verheijen (2001, hereafter V01), who carried out a detailed investigation of the TF relation in the Ursa Major cluster using resolved HI rotation curves and optical and infrared surface photometry.  Numerous papers have used TF results to derive conclusions about the physics of galaxy formation and the relative importance of baryons and dark matter in the luminous regions of disk galaxies (e.g., \\citealt{fab82,gun82,per88,col89,whi91,kau93,col94, eis96,per96,dal97,mo98,cou99,fir00,nav00,dut06,gne06}). Our underlying goal is to provide better observational inputs for these kinds of theoretical and modeling efforts.  We have therefore paid particular attention to characterizing our sample selection and observed quantities as straightforwardly as we can, and to estimating the intrinsic scatter in the different SDSS bands.  We have already presented the scaling relations and inferred baryonic mass fractions of a disk-dominated subset of our full sample in P05, and \\cite{gne06} have used these results to constrain the distributions of disk-to-halo mass ratios and spin parameters of spiral galaxies. ",
        "conclusions": "We have measured the TF relation for a sample of 162 galaxies selected from the SDSS, using follow-up spectroscopy to obtain \\halpha\\ rotation curves.  We targeted 234 galaxies from the SDSS spectroscopic galaxy catalog that have a roughly flat absolute magnitude distribution in the range $-18.5 > \\Mr > -22$.  For the purpose of testing galaxy formation models, our sample has several advantages relative to most previous TF studies.  Our target selection is blind to morphology, except for an $i$-band axis ratio cut of 0.6. Our completeness is high, and the galaxies with usable \\halpha\\ rotation curves have distributions of color and concentration similar to a large control sample of galaxies selected from DR4 with the same $\\Mr$ distribution but no inclination cut or \\halpha\\ requirement. Due to the uniform SDSS photometry and the selection of galaxies at redshifts greater than 5000 $\\kms$, all sources of observational error are small compared to the estimated intrinsic scatter.  The uniform multi-wavelength SDSS photometry allows us to investigate the TF relations and the correlations between TF residuals, color residuals, and other structural parameters in the $g$, $r$, $i$, and $z$-bands with minimal photometric uncertainties.  We adopt $\\Veight$\\ as a measure of the rotation speed and use extinction corrected absolute magnitudes. We estimate the slope, intercept, and intrinsic scatter of the TF relations simultaneously, using a maximum likelihood procedure that accounts for individual observational errors in absolute magnitude and rotation speed. We measure forward TF slopes between $-5.5$ and $-6.6\\,{\\rm mag}/{\\rm log_{10}}(\\kms)$, with typical uncertainty of 0.2, and intercepts between $-20.7$ and $-21.5$ mag at $\\eta = 2.22$, with typical uncertainty of 0.04. The slope becomes systematically steeper with wavelength. Inverse fits always yield steeper slopes, but once the effects of Gaussian scatter are included, the forward and inverse fits describe essentially the same two-dimensional distribution of data points over the range covered by our data. Corrections to these slopes due to Malmquist-type biases are discussed in Appendix B, where it is shown that our observed TF relation can be reproduced with a steeper slope and larger intrinsic scatter given certain assumptions about selection effects.  The intrinsic scatter appears to be nearly independent of wavelength or fitting procedure, typically $\\sigma = 0.42-0.46$ mag, with a higher value (0.54 mag) for the inverse fit in $g$-band. The distribution of residuals is approximately Gaussian for both the forward and inverse relations, and there is no indication of rare outliers inflating the TF scatter, with the possible exception of one galaxy whose axis ratio may be incorrectly measured because of spiral arms. Omitting this one system reduces the estimated intrinsic scatter by 0.02 mag. The intrinsic scatter is slightly smaller for a disk-dominated subset of galaxies, decreasing in the $i$-band from 0.42 mag to 0.36 mag for galaxies having disk-to-total flux ratios greater than 0.9. Morphologically asymmetric, barred, or possibly interacting galaxies show no clear evidence for offsets from the mean TF relation or for larger scatter, though our statistics for these subsets are limited.  In an attempt to understand the origin of scatter in the TF relation, we study correlations between TF residuals and other galaxy properties: $g-r$ color, half-light radius, and $i$-band concentration index. The $g$-band TF residual shows a significant correlation with color, in the sense that bluer galaxies are brighter at fixed $\\eta$ or (more clearly visible in the data) that redder galaxies rotate faster at fixed $M_g$.  However, these correlations are much weaker in the $i$-band, and they are much weaker than predicted for pure self-gravitating disks with the stellar mass-to-light ratios predicted from $g-r$ color.  The correlation with half-light radius is very weak, with a slight tendency for larger galaxies to be more luminous at fixed $\\eta$. More concentrated (earlier type) galaxies tend to be slightly fainter at fixed $\\eta$ or to rotate slightly faster at fixed $M_i$, but the trend is weak compared to the TF scatter.  Our TF relation is similar to that found for field spirals by C97, who used similar analysis methods. The TF relation found by V01 for Ursa Major spirals, using the flat portion of HI rotation curves as a velocity measure, is significantly steeper than ours.  This difference may partly reflect the difference in velocity width definitions, and Malmquist-type biases as described in Appendix B.  The residual correlations in our sample are similar to those found by K02 using the Nearby Field Galaxy Survey of \\cite{jan00}, but we find that stellar population effects explain only a small fraction of the scatter in the TF relation.  The nearly constant intrinsic scatter of the TF relation has interesting implications for galaxy formation theories.  The scatter is only slightly larger in bluer bands and only slightly smaller for a morphologically selected disk-dominated sub-sample of galaxies, and the correlation of TF residuals with other properties explains only a small fraction of the scatter.  The lack of correlation with radius implies that disk gravity has small influence on galaxy dynamics at the radii used for TF measurements; this observational result provides a strong constraint on galaxy formation models \\citep{dut06,gne06}.  The weak correlation with color, size, and morphology implies that most of the scatter in the TF relation arises from genuine variation in the ratio of dark matter to baryonic matter from halo to halo.  Our results suggest many fruitful directions for future observational investigations, each of them comparable in magnitude to the one undertaken here. Absorption-line velocity measurements, which require longer exposures than our \\halpha\\ emission line measurements, could yield rotation velocities for the remaining 30\\% of our sample, showing whether galaxies without \\halpha\\ emission follow a systematically different TF relation. HI measurements would allow detailed comparison between TF relations defined by \\halpha\\ and HI velocity widths (see C97) for our broadly selected galaxy sample, and they would provide gas masses and gas fractions as additional parameters for residual correlation studies, which would be valuable for testing theoretical models \\citep{dut06,gne06}. Extending the TF relation to brighter and fainter luminosities is of great interest (e.g., \\citealt{mcg00}), though complications arise from the growing paucity of disks in the first case and the increasing importance of non-circular motions in the second. Finally, it is clear that detailed investigation of TF residuals would benefit from a much larger sample with similarly broad selection criteria and similarly small observational errors, so that one could, for example, investigate residual correlations with size, morphology, or environment within restricted regions of luminosity-color space. Since large samples of galaxies ($\\ge 2000$) have been observed for TF studies in the past (see, e.g., \\citealt{cou03,mas06} and references therein), it may be possible to construct such a sample largely from archival data, if the data quality can be made sufficiently homogeneous and the effects of sample selection can be sufficiently well understood. Our data already provide a well defined and highly constraining target for theoretical models of galaxy formation.  Extended observational studies would improve our ability to deduce the essential physics of disk galaxy formation empirically."
    },
    "0608/astro-ph0608212_arXiv.txt": {
        "abstract": "We present {\\it Galaxy Evolution Explorer} ({\\it GALEX}) far (FUV) and near (NUV) ultraviolet imaging of the nearby early-type galaxy NGC~2974, along with complementary ground-based optical imaging. In the ultraviolet, the galaxy reveals a central spheroid-like component and a newly discovered complete outer ring of radius $6.2$~kpc, with suggestions of another partial ring at an even larger radius. Blue FUV$-$NUV and UV$-$optical colours are observed in the centre of the galaxy and from the outer ring outward, suggesting young stellar populations ($\\la1$~Gyr) and recent star formation in both locations. This is supported by a simple stellar population model which assumes two bursts of star formation, allowing us to constrain the age, mass fraction and surface mass density of the young component pixel by pixel. Overall, the mass fraction of the young component appears to be just under $1$~per cent (lower limit, uncorrected for dust extinction). The additional presence of a nuclear and an inner ring (radii $1.4$ and $2.9$~kpc, respectively), as traced by [O\\,{\\small III}] emission, suggests ring formation through resonances. All three rings are consistent with a single pattern speed of $78\\pm6$ km~s$^{-1}$~kpc$^{-1}$, typical of S0 galaxies and only marginally slower than expected for a fast bar if traced by a small observed surface brightness plateau.  This thus suggests that star formation and morphological evolution in NGC~2974 at the present epoch are primarily driven by a rotating asymmetry (probably a large-scale bar), despite the standard classification of NGC~2974 as an E4 elliptical. ",
        "introduction": "\\label{sec:intro} Early-type galaxies are traditionally viewed as dynamically simple stellar systems with homogeneous stellar populations \\citep[e.g.][]{g77}. However, building on earlier work (e.g.\\ \\citealt{zf91} and references therein), the {\\tt SAURON} survey \\citep{betal01,zetal02} of the two-dimensional ionised-gas/stellar kinematics and stellar populations of nearby early-type galaxies has recently revealed kinematically-decoupled components and substantial metallicity and/or age gradients in many objects \\citep[e.g.][]{eetal04,setal06,kuetal06}. Early-type galaxies are thus likely to have had complex and varied formation histories. Such properties are often taken as evidence against the classical monolithic collapse model \\citep*[e.g.][]{els62,la74}, in which elliptical galaxies are old systems that formed early in highly efficient starbursts and have evolved passively since.  Indeed, while standard optical imaging shows that most early-type galaxies have red colours and old stellar populations, deep imaging surveys have shown that many possess shells and tidal features \\citep*[e.g.][]{ss92}, telltale signs of a turbulent past. Colour gradients also indicate that early-type galaxies have composite stellar populations \\citep[e.g.][]{pe90}, and some nearby early-type galaxies have signatures of ongoing or recent star formation \\citep*[RSF; e.g.\\ NGC~2865;][]{hcb99,retal05}. Spectroscopic studies further suggest that many elliptical galaxies have intermediate-age stars \\citep[e.g.][]{oco80,ba87,tetal00}. The spread in the luminosity-weighted age \\citep[e.g.][]{tetal00} is evidence for RSF. The so-called E~+~A galaxies \\citep{dg83} possess characteristic spectral features defined by strong hydrogen Balmer absorption lines (H$\\delta$, H$\\gamma$, H$\\beta$), indicating the presence of a young stellar population, but have negligible [O\\,{\\small II}] $\\lambda$3727 emission. These galaxies are believed to be `post-starburst' systems, i.e.\\ they have recently undergone a burst of star formation. \\citet{fetal04} also suggested that a small fraction of low redshift elliptical galaxies have unambiguous signatures of active star formation with rates as high as those in spirals.  Far ultraviolet (FUV) radiation was first discovered in early-type galaxies by the {\\it Orbiting Astronomical Observatory-2} in 1969 \\citep*{cwp72}, but it is now generally recognized that UV emission is a ubiquitous phenomenon in these objects. Although star formation (SF) can also be measured indirectly at other wavelengths, particularly in the infrared \\citep[e.g.][]{hetal86,metal97} and through H$\\alpha$ emission \\citep*[e.g.][]{ke83,hbi03}, UV light is very sensitive to stellar populations younger than a few hundred megayears and the UV flux can give important and direct information about the star formation rate (e.g.\\ \\citealt{hbi03}; \\citealt*{betal05}). The FUV is sensitive to the UV upturn flux originating from old hot helium-burning horizontal branch (HB) stars (see, e.g., \\citealt*{ydo97} and \\citealt{oco99} for reviews of this issue and the UV upturn phenomenon), hence SF studies usually focus on the near ultraviolet (NUV) passband. \\citet{fs00} studied young stars in the cluster Abell~851 ($z=0.41$) using NUV$-$optical colours and suggested that, for some of the early-type galaxies, the stellar mass fraction in young stars is higher than $10$~per cent. More recently, \\citet{yetal05} investigated the NUV colour-magnitude relation of early-type galaxies classified by the Sloan Digital Sky Survey. They concluded that roughly $15$~per cent of nearby bright ($M_r<-22$) early-type galaxies show signs of recent ($\\la1$~Gyr) star formation (at the $1-2$~per cent level by stellar mass).  Here, we focus on the nearby early-type galaxy NGC~2974, classified as E4 in the RC3 \\citep{vetal91} but possessing many features (mainly dynamical) more reminiscent of lenticular galaxies \\citep[see, e.g.,][]{cm94}. NGC~2974 has in fact been widely studied because of its unusual physical properties, including a dust lane containing ionised gas \\citep[e.g.][]{k89} and substantial neutral hydrogen (H\\,{\\small I}) consistent with a rotating disc aligned with the optical isophotes \\citep[e.g.][]{kjgkg88}. \\citet{betal93} also suggested that it has a complex lens-like feature. Recently, using data from the {\\tt TIGER} integral-field spectrograph and {\\it Hubble Space Telescope}'s ({\\it HST}) Wide Field and Planetary Camera~2 (WFPC2), \\citet*{egf03} discovered a gaseous two-arm spiral structure in the central $500$~pc, suggesting the presence of a nuclear bar. The larger scale {\\tt SAURON} integral-field observations of \\citet{kcemz05} reveal an apparent nuclear bar, nuclear ring, larger scale spiral arms and a possible inner ring in the ionised gas map (see Section~\\ref{sec:rings}), which suggest the presence of a large-scale bar as well.  In this paper, we present and discuss new UV imaging observations of NGC~2974 obtained with {\\it Galaxy Evolution Explorer} ({\\it GALEX}). In Section~\\ref{sec:obs}, we describe the {\\it GALEX} observations and data reduction, supporting optical imaging, and the main results from UV$-$optical surface photometry. Recent star formation in NGC~2974 is quantified and discussed in Section~\\ref{sec:spop_sf}, and evidence for a rotating pattern, mainly based on the discovery of a star-forming outer ring, is discussed in Section~\\ref{sec:fig_rot}. We summarize our results and discuss their implications briefly in Section~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary} We have presented {\\it GALEX} FUV and NUV and ground-based F555W imaging of the nearby early-type galaxy NGC~2974. The optical images are characteristic of elliptical galaxies but the UV images reveal both a central spheroid-like component and a complete (outer) ring of radius $6.2$~kpc. Another partial ring is also suggested at an even larger radius. Both rings are detected here for the first time and are only visible in the UV. The FUV luminosity of the outer ring is $\\approx1.2~\\times10^{8}$~$L_{\\odot}$.  Blue FUV$-$NUV colours are observed in the centre of the galaxy and around the outer ring, suggesting the presence of young stellar populations ($\\la1$~Gyr) and recent star formation. UV$-$optical colours are also blue at both locations, indicating that (very) young stars contribute significantly to the galaxy's UV flux. Simple two-component stellar population modeling allows us to constrain the age, stellar mass fraction and surface mass density of the young component pixel by pixel and confirms this behaviour. The latter are formally lower limits since we have not corrected for dust extinction, but the effect is expected to be small.  Considering the mass model and observations of \\citet{kcemz05}, the {\\tt SAURON} [O$\\,${\\small III}] nuclear and inner rings and the {\\it GALEX} UV outer ring are all consistent with a single pattern speed of $78\\pm6$ km~s$^{-1}$~kpc$^{-1}$, suggesting that NGC~2974 harbours a large-scale $m=2$ asymmetry such as a bar. Assuming that the surface brightness plateau observed at a radius of $20$--$25\\arcsec$ marks its limits, this large-scale bar would end at roughly $60$~per cent of its own corotation radius, slightly smaller but not inconsistent with expectations.  Star formation, rings and bars are not topics normally associated with elliptical galaxies, yet in recent years we have seen a growing number of studies reporting evidence for complex kinematics \\citep[e.g.][]{eetal04} and recent star formation \\citep[e.g.][]{yetal05}. The origin of those features are of key importance. Here, we have witnessed evidence for recent star formation in a particular early-type galaxy, in the form of a ring, fueled by a newly discovered co-spatial ring of neutral hydrogen \\citep{wkomz06}. A rotating bar (i.e.\\ a disc phenomenon) seems to be regulating both phenomena, adding to earlier evidence for lenticular-like properties in NGC~2974. Star formation may thus not be all that unusual in early-type galaxies, or perhaps elliptical galaxies in the classical sense are simply far rarer than usually assumed."
    },
    "0608/astro-ph0608024_arXiv.txt": {
        "abstract": "The temporal and spectral properties of ultraluminous X-ray sources (ULXs, $L_{\\rm X}>2\\times10^{39}$ ergs s$^{-1}$) and bright X-ray sources ($L_{\\rm X}>3\\times10^{38}$ ergs s$^{-1}$) are examined and compared in two extremely different host environments: the old elliptical galaxy NGC 1399 and the young, starforming Antennae galaxies (NGC 4038/4039).  ULXs in NGC 1399 show little variability on either long or short time scales.  Only 1 of 8 ULXs and 10 of 63 bright sources in NGC 1399 are variable at a confidence level of 90\\%. On long timescales, the NGC 1399 sources are steadier than most Galactic black hole X-ray binaries, but similar to GRS 1915$+$105.  The outburst duration of the NGC 1399 sources is about 20 yrs, again, similar to that of GRS 1915$+$105.  The bright X-ray sources in NGC 1399 may be black hole X-ray binaries with giant star companions similar to GRS 1915$+$105.  The brightest ULX ({PSX-1}) in NGC 1399 is coincident with a globular cluster, shows a hard spectrum with a photon index around 1.5, and has a nearly constant luminosity around $5\\times10^{39}$~erg~s$^{-1}$.  It may be an intermediate-mass black hole (IMBH) in a hard spectral state.  In contrast to NGC 1399, the ULXs in the Antennae are all variable and a large fraction of the bright sources (9 of 15) are also variable.  The variability and luminosity of ULXs in the Antennae suggest they are black hole high mass X-ray  binaries accreting via Roche-lobe overflow.  A flare with a duration of about 5~ks is found from Antennae X-42.  The most luminous ULX, X-16, with a very hard  spectrum ($\\Gamma$ = 1.0$\\sim$1.3) and a luminosity which varies by a factor of 10, could be an IMBH candidate. ",
        "introduction": "Ultraluminous X-ray sources are point-like X-ray sources in external galaxies with positions off the nucleus and apparent isotropic luminosities in excess of the Eddington luminosity of a 10$M_\\sun$ black hole \\citep[$\\sim 2\\times10^{39}$ ergs s$^{-1}$; for review see][]{fab06,mil04a}.  ULXs are often found in actively starforming galaxies, e.g.\\ M82 \\citep{kaa01} and the Antennae galaxies \\citep{zez02-2}.  There are statistically more ULXs in starburst galaxies compared to spiral galaxies \\citep{kil05}.  \\citet{kaa04} investigated three star burst galaxies and found X-ray sources were preferentially near star clusters but displaced with higher luminosity sources nearer to clusters.  All these suggest that ULXs trace young stellar objects.  Some surveys have reported higher numbers of ULXs in elliptical galaxies than in spirals \\citep{col02}, but the elliptical galaxies suffer more contamination from background sources \\citep{lop06}. \\citet{irw03} found that bright X-ray sources in elliptical galaxies have spectra consistent with Galactic X-ray binaries in the high state, but that the number of sources in the ULX range is consistent with the number of expected background sources.  However, some elliptical galaxies contain ULXs associated with globular clusters, e.g.\\ NGC 1399 \\citep{ang01} and NGC 720 \\citep{jel03}, which strongly suggests an association with the elliptical galaxy.  \\begin{deluxetable*}{lllcll}[t] \\tablecaption{Chandra Observations of NGC 1399 and the Antennae. \\label{tab:obs}} \\tablehead{ \\colhead{Galaxy} & \\colhead{index} & \\colhead{obs. ID} & \\colhead{exposure (ks)} & \\colhead{Instrument} & \\colhead{obs. date} } \\startdata & (a) & 320 & 3.5 & ACIS-I/NONE & 1999 Oct 18 \\\\ & (b) & 319 & 55.9 & ACIS-S/NONE & 2000 Jan 18 \\\\ & (c) & 239 & 3.6 & ACIS-I/NONE & 2000 Jan 19 \\\\ NGC 1399 & (d) & 49898 & 13.1 & ACIS-S/HETG & 2000 Jun 15 \\\\ & (e) & 240 & 43.5 & ACIS-S/HETG & 2000 Jun 16 \\\\ & (f) & 2389 & 14.7 & ACIS-S/HETG & 2001 May 08 \\\\ & (g) & 4172 & 44.5 & ACIS-I/NONE & 2003 May 26 \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} & (i) & 315 & 72.2 & ACIS-S/NONE & 1999 Dec 01 \\\\ & (ii) & 3040 & 69.0 & ACIS-S/NONE & 2001 Dec 29 \\\\ & (iii) & 3043 & 67.1 & ACIS-S/NONE & 2002 Apr 18 \\\\ The Antennae & (iv) & 3042 & 67.3 & ACIS-S/NONE & 2002 May 31 \\\\ & (v) & 3044 & 36.5 & ACIS-S/NONE & 2002 Jul 10 \\\\ & (vi) & 3718 & 34.7 & ACIS-S/NONE & 2002 Jul 13 \\\\ & (vii) & 3041 & 72.9 & ACIS-S/NONE & 2002 Nov 22 \\\\ \\enddata \\end{deluxetable*}  NGC 1399 and the Antennae are cases, respectively, of an elliptical galaxy and a starforming galaxy that host a large number of ULXs. \\citet{ang01} compared {\\it Chandra} and {\\it Hubble Space Telescope} ({\\it HST}) images and found 70\\% of low mass X-ray binaries (LMXBs) in NGC 1399 were associated with globular clusters (GCs); the GC associated sources are brighter than field sources on average.  The X-ray source population in NGC 1399 may be unusual because a recent survey of LMXBs in six elliptical galaxies shows that, in general, there is no significant difference in luminosity between GC associated and field LMXBs \\citep{kim05}.  Several ULXs in NGC 1399 are detected but their natures are still unclear; they could be IMBHs or multiple LMXBs \\citep{ang01}.  X-ray sources in the Antennae have been well studied along with a series of observations from {\\it Chandra} \\citep{fab01,zez02-2}.  Comparison with optical images shows that the Antennae X-ray sources are predominately associated with young star clusters \\citep{zez02-3}.  The Antennae ULXs have spectra consistent with ULXs in nearby galaxies and are not supernova remnants because they are variable and lack radio emission \\citep{zez02-3}. \\citet{zez02-4} argued that IMBHs could account for a few of the ULXs in the Antennae but not most of them, which were more likely to be supercritical accreting stellar mass X-ray binaries due to their displacement from young star clusters \\citep{zez02-2}. The ULX X-37 in the Antennae has been identified as a background quasar \\citep{cla05}.  The most interesting ULXs are potential IMBHs. NGC 1399 and the Antennae are two extremely different host environments for ULXs. The comprehensive connection between ULXs and young stellar objects makes ULXs in NGC 1399 mysterious.  A comparison of ULXs between these two system may shed light on their natures.  We analyzed {\\it Chandra} archival data for both galaxies (\\S~\\ref{sec:obs}) in order to measure their long-term and short-term time variability (\\S~\\ref{sec:var}) and spectral variation (\\S~\\ref{sec:spec}).  We discuss the results and the implications in \\S~\\ref{sec:diss}.  We adopt the distance to NGC 1399 as 20.5 Mpc \\citep{mer01} and the distance to the Antennae as 13.8 Mpc \\citep{sav04}.  This new distance to the Antennae is obtained by measurement of the red giant branch tip and appears accurate to within 2~Mpc, while the older distance was estimated from the recession velocity.  The new distance is more likely to be accurate.  We note that this revised distance to the Antennae has reduced the Antennae source luminosities by a factor of 4 relative to previous X-ray studies. ",
        "conclusions": "\\label{sec:diss}  The most dramatic difference between NGC 1399 and the Antennae is that the ULXs as well as the bright X-ray sources in the elliptical galaxy are less variable.  In Fig.~\\ref{fig:lln}, only a couple sources show apparent variability in the L-L diagram of NGC 1399 between (b)-(e) or (b)-(g) observations.  In Fig.~\\ref{fig:lla}, the number of sources in the Antennae with significant variability is obviously larger than that in NGC 1399.  This is verified by lightcurves of ULXs and the variability functions of bright sources.  The variability functions of Galactic LMXBs show that most black holes are highly variable, with the notable exception of GRS 1915$+$105, whereas bright neutron stars are less variable. The normalized variability function (bottom panel of Fig.~\\ref{fig:vfun}) presents the bright sources in NGC 1399 as similar to 4U 1820$-$30, a little less variable than the group of neutron star LMXBs and GRS 1915$+$105, and far less  variable than the group of black hole LMXBs.  ULXs in NGC 1399 are even less variable than other bright sources, with 7 of 8 showing constant flux and the other having little variability.  None of the bright X-ray sources or ULXs in NGC 1399 show short-term variability with a significance level of 3$\\sigma$ in any individual observation.  The lack of variability should provide a strong constraint on the nature of these sources.  \\begin{figure}[!h] \\centering \\plotone{f8.eps} \\caption{0.3-10 keV luminosity vs. power-law photon index ($\\Gamma$) for the 8 ULXs in NGC 1399. All points in the plot come from observations that have enough photons for a spectral fitting. PSX-1, PSX-2, PSX-4 and PSX-7 have two points from observations (b) and (g), and rest sources have one from the observation (b). The absorption in the fitting is set as a free parameter but with the lower bound to the Galactic value ($1.34\\times10^{20}$ cm$^{-2}$). The best fitted $n_{\\rm H}$ are all close or equal to the Galactic value (less than 3$\\sigma$ except PSX-2 in observation (b), which is nearly 6$\\sigma$). \\label{fig:nlg}} \\end{figure}  \\begin{figure}[!h] \\centering \\plotone{f9a.eps}\\\\ \\plotone{f9b.eps}\\\\ \\plotone{f9c.eps}\\\\ \\plotone{f9d.eps} \\caption{0.3-10 keV luminosity vs. power-law photon index ($\\Gamma$) for the 4 ULXs in the Antennae in seven observations. The absorption in the fitting is set as a free parameter but with the lower bound to the Galactic value ($3.83\\times10^{20}$ cm$^{-2}$). The best fitted $n_{\\rm H}$ values are consistent in all observations and the weighted average number is shown for each source. \\label{fig:alg}} \\end{figure}  The long term variability also allow us to estimate the outburst duration of X-ray sources in NGC 1399.  We found that the outburst duration is around 20 yrs for both bright X-ray sources and ULXs.  A natural way to produce low variability would be if the bright X-ray sources in NGC 1399 consist of multiple unresolved LMXBs.  However, a large number of sources is required.  From the ASM data, we estimated that neutron star LMXBs are  above 10\\% of their maximum flux 80\\% of the time, while black holes are above their 10\\% of their maximum flux only 15\\% of time (the estimate for black holes is an upper bound because several of our selected black hole systems have had only one outburst during the lifetime of RXTE).  Using the average luminosity of 4U 1820$-$30 of $\\sim 5\\times 10^{37} \\rm \\, erg \\, s^{-1}$, about 8 neutron star binaries would be needed to produce the least luminous of our bright sources and 50 would be needed to produce the least luminous ULX.  Large numbers of neutron stars are known to be present in some globular clusters in our galaxy, e.g.\\ there are around 50 in the Galactic GC NGC 6266 \\citep{poo03}, but only very rarely does a cluster contain more than one bright X-ray source.  It is very puzzling to understand how individual clusters in NGC 1399 could contain so many bright X-ray sources.  Though the luminosity of black holes in the accreting phase is, generally, several times larger than that of neutron stars, the smaller duty cycle partially offsets the increased luminosity and comparable numbers of black hole binaries are needed to reach the same luminosity.  Perhaps, the most promising interpretation of the NGC 1399 bright sources is that they are are individual black holes of a type similar to GRS 1915$+$105.  As noted above, GRS 1915$+$105 has been a persistent X-ray source since its discovery 14 years ago \\citep{cas94} and has reached luminosities in the ULX range.  The variability function of GRS 1915$+$105 is similar to that of the NGC 1399 bright sources.  The companion star in GRS 1915$+$105 is a K or M giant \\citep{gre01}.  The time scale for the evolution of such a binary system to the mass-transferring phase is roughly 6~Gyr \\citep{bel02}. Allowing for different companion star masses, similar mass transferring binaries could currently be present in either the old ($\\sim 11$~Gyr) or young ($\\sim 2$~Gyr) globular clusters in NGC 1399 \\citep{for01}. \\citet{kin02} has previously suggested that the ULXs in ellipticals may be similar to GRS 1915$+$105.  \\citet{pir02} suggested that the bright X-ray sources in ellipticals are X-ray binaries with Roche-lobe filling giant companions with mass transfer driven by nuclear evolution of the companion.  They predicted outburst durations similar to what we find for NGC 1399.  For the sources associated with globular clusters, another possibility is that the companion stars are initially isolated stars which are captured by the black hole after evolving to the giant phase.  The capture cross-section is enhanced for stars in the giant phase.  However, GRS 1915$+$105 shows strong variability on short time scales, which we do not observe from the bright X-ray sources in NGC 1399.  The sensitivity of Chandra to short timescale variability is limited due to the low counting rate, but our detection of short timescale variability from sources in the Antennae (with comparable counting rates) indicates that similar flares would have been detected from the NGC 1399 sources.  The short timescale variability of sources in NGC 1399 should be investigated in more detail with new observations.  \\citet{bil04} has suggested that the bright X-ray sources in NGC 1399 may be neutron stars accreting from white dwarf companions in ultracompact binaries, similar to 4U 1820$-$30.  The variability of 4U 1820$-$30 is similar to the NGC 1399 bright sources. Hence, neutron star ultracompact binaries may account for some or all of the bright sources if mildly super-Eddington (apparent) luminosities can be achieved.  However, the short lifetimes expected for the high accretion rates required given the very low-mass companions ($0.04 - 0.08 M_{\\odot}$) make it unlikely that the ULXs in NGC 1399 are neutron star ultracompact binaries.  Another possible interpretation of the bright sources is that they are neutron-star binaries similar to the neutron-star X-ray binaries known as Z sources.  Z sources are persistent with variability similar to that seen from the NGC 1399 sources.  Again, mildly super-Eddington (apparent) luminosities would be required to produce the dimmest of the bright sources and such objects are unlikely to explain the ULXs. Interestingly, \\citet{web83} suggested that the Z sources have giant star companions with mass loss driven by nuclear evolution.  This would make the systems the neutron star analogs of black hole/giant star binaries such as GRS 1915$+$105.  The three most luminous X-ray sources in NGC 1399 are interesting. The brightest one, PSX-1, has a hard spectrum ($\\Gamma=1.2\\sim1.5$), little variability ($L_{\\rm max}/L_{\\rm min}=2.2\\pm0.8$), and a high luminosity around $5\\times10^{39}$ ergs s$^{-1}$.  The spectrum indicates the source is not in its high/soft state ($\\Gamma=2.1-4.8$) but at the hard extreme of the low/hard state ($\\Gamma=1.5-2.1$) \\citep{mcc06}.  The low variability indicates that the source is probably not relativistically beamed \\citep{kor02}, because relativistically beamed sources are expected to be highly variable. This source is associated with a bright GC with $C-R=2.12$, $m_R=21.02$ \\citep{dir03}, which means it is unlikely to be a background AGN.  A huge number of sources would be required to explain the emission as a collection of neutron star or black hole X-ray binaries.  If the source is an IMBH in the low state, its practically constant luminosity is reasonable if compared to the long-term behavior of M82 X-1, which varies in luminosity by a factor of 2-3 \\citep{kaa06a,kaa06b} and shows a hard spectrum with a power-law index $\\Gamma \\approx 1.6$ \\citep{kaa06b}.  NGC 1399 PSX-2 and PSX-4 display typical behaviors of black holes in the high/soft state.  PSX-2 is associated with a bright GC  ($C-R = 2.24$, $m_R = 21.36$) but PSX-4 is not. The plausible cool disk ($kT \\approx 0.4$ keV) emission is consistent with the soft spectral index ($\\Gamma \\approx 2.5$).  Their low variability ($L_{\\rm max}/L_{\\rm min} = 2.5 \\pm 0.7$ for PSX-2 and $1.5 \\pm 0.6$ for PSX-4) is suggestive of some Galactic black holes in the high state, but no Galactic source has been found to remain in the high state for four years \\citep{mcc06}.  The cool disk has been a useful probe in finding IMBH candidates \\citep{col99,kaa03,mil04,fen05}.  However, their disk temperatures are somewhat higher than most IMBH candidates, and their luminosities ($2\\sim3\\times10^{39}$ ergs s$^{-1}$) are also around the lower bound of those candidates \\citep{mil04,fen05}.  These sources do not have high mass companion stars to produce geometrically beamed emission as suggested by \\citet{kin01}.  Therefore, it could be possible that these two sources are 10$\\sim$100$M_\\sun$ black holes emitting at the high state if the disk components are true.  One possible way to determine the reality of the disk emission is an accurate measurement of the column density, since these two sources, in particular PSX-2, show significantly higher intrinsic absorption if modeled by a single power-law spectrum.  The bright X-ray sources in the Antennae are variable.  This suggests that they are black hole binaries and accreting from massive stars to produce unstable disks \\citep{kal04}.  They are not as variable as black hole LMXBs (Fig.~\\ref{fig:vfun}).  Considering the variability together with their high luminosity, we suggest they are HMXBs accreting via Roche-lobe overflow.  There is no disk emission found in any of the 4 ULXs.  The most luminous source, X-16, manifests a very hard spectra ($\\Gamma=1.0-1.3$) and has strong variability with a change of luminosity of $\\sim$10 (Fig.~\\ref{fig:alg}).  The source could be an IMBH with a mass of a few hundred to thousand $M_\\sun$ emitting at the hard state.  The X-ray behavior is also consistent with a relativistically beamed jet pointing at us, but \\citet{zez02-3} pointed out that there is no radio emission found from any ULX in the Antennae.  It is interesting that none of the ULXs shows the very soft spectrum normally seen in ULXs in nearby galaxies \\citep{fen05}.  We note that X-37 was reported to have a soft component in its spectrum, but has been subsequently identified as a background AGN.  X-11 and X-42 are likely to be in the hard state ($\\Gamma \\le 2$), X-44 looks like it is between the hard and soft states ($\\Gamma \\sim 2$).  A plausible correlation ($r=0.61$) is found in X-11 between luminosity and $\\Gamma$, which shows similar behaviors as stellar mass black holes, but we should be cautious because of the weak correlation coefficient.  The X-ray properties of these three sources are consistent with HMXBs, but whether they are IMBHs or $\\sim$20$M_\\sun$ black holes emitting in near or super Eddington luminosities is still unclear."
    },
    "0608/nucl-ex0608027_arXiv.txt": {
        "abstract": "Targets consisting of $^{3,4}$He implanted into thin aluminum foils (approximately 100, 200 or 600 $\\mu$g/cm$^2$) were prepared using intense (a few $\\mu$A ) helium beams at low energy (approximately 20, 40 or 100 keV). Uniformity of the implantation was achieved by a beam raster across a 12 mm diameter tantalum collimator at the rates of 0.1 Hz in the vertical direction and 1 Hz in the horizontal direction. Helium implantation into the very thin (approximately 80-100 $\\mu$g/cm$^2$) aluminum foils failed to produce useful targets (with only approximately 10\\% of the helium retained) due to an under estimation of the range by the code SRIM. The range of low energy helium in aluminum predicted by Northcliffe and Shilling and the NIST online tabulation are observed on the other hand to over estimate the range of low energy helium ions in aluminum. An attempt to increase the amount of helium by implanting a second deeper layer was also carried out, but it did not significantly increase the helium content beyond the blistering limit ($~6 \\ \\times \\ 10^{17}$ atoms/cm$^2$). The implanted targets were bombarded with moderately intense \\he4 and $^{16}$O beams of 50-100 particle nA . Rutherford Back Scattering of 1.0 and 2.5 MeV proton beams and  recoil helium from 15.0 MeV oxygen beams were used to study the helium content and profile before, during and after bombardments. We observed the helium content and profile to be very stable even after a prolonged bombardment (up to two days) with moderately intense beams of $^{16}$O or \\he4. Helium implanted into thin (aluminum) foils is a good choice for thin helium targets needed, for example, for a measurement of the \\xhe3ag reaction and the associated $S_{34}$ astrophysical cross section factor (S-factor). ",
        "introduction": "Several key reactions of interest in nuclear physics and nuclear astrophysics involve light particles such as protons, deuterons, and alpha particles. The use of radioactive beams and inverse kinematics makes the availability of targets containing these elements necessary. Previous experiments in the field have used a variety of cryogenic, solid, and gas targets. One solution that has been used for hydrogen and deuterium is a thin polyethylene target. This solution cannot be used for helium, however, and some published experiments have used gas cells for helium targets.  A helium gas jet target that was used in a measurement of the \\xhe3ag reaction \\cite{Kre82} was found later to have significant non-uniformity and yield a cross section which is approximately 40\\% smaller than measured by the same group in a later experiment \\cite{Hil88}. Another method that has shown promise is the implantation of helium into a thin metal foil ({\\em e.g.}, Al, Ni, Cu, Mo, Au). This has been used in a number of experiments, such as measurements of the \\xhe3ag reaction cross section \\cite{Al84}, as well as the $^3He(d, p)^4He$ reaction cross section \\cite{Ge96}. Noble gases, such as Ar and Xe \\cite{Sp93,Gr95} implanted into metals were also used as targets. Indeed, noble gasses are known to form stable structures composed of bubbles with radii of approximately 100 nm when implanted into metals.  An implanted target  suffers from a low number density, approximately a factor of 10 below targets commonly used in reaction studies in nuclear astrophysics, but has many advantages over a gas cell, with the most significant advantage being a well-localized target with high energy resolution. For example, a \\ahe3 implanted target with $5 \\times 10^{17} \\  atoms/cm^2$ (just below the blistering limit) and a 1.0 MeV \\he4 beam of modest intensity of 500 nA (280 pnA) yield a count rate of 40 counts per hour in a HPGe detector with a 2\\% efficiency for detecting the resulting 2.016 MeV direct capture gamma rays at 90$^\\circ$. The large angular acceptance of such a detector placed close to the target results an energy spread with a FWHM of approximately 30 keV. Hence an implanted target in combination with moderately intense beams is useful for studying the \\xhe3ag reaction and its associated astrophysical cross section factor ($S_{34}$) at modestly low energies of approximately $E_{cm}$ = 400 keV.  One area of concern with implanted targets is the stability of the implanted gas over time, especially when used in beams. Beam heating has been cited as a problem \\cite{Al84}. For example it was noted \\cite{Al84} that the helium content in the center of the target was depleted after bombardment with an intense (2-3 $\\mu$A) \\he4 beam even though the implanted target was water cooled.  In this study we used helium implanted into thin aluminum foils (approximately 100, 200 or 600 $\\mu g/cm^2$) in order to minimize the beam heating. We tested these thin implanted targets before and after bombardment with moderately intense beams (approximately 100 pnA ) of $^{16}$O and \\he4. We did not observe a depletion in the helium content even after a prolonged (as much as two day) bombardment. The helium content and profile measured using back scattering of 1.0 MeV proton beams were observed to be very stable. ",
        "conclusions": "We prepared $^{3,4}$He targets by implanting low energy helium beams into thin aluminum foils and tested them using (mostly Rutherford) Back-Scattering of low energy proton and oxygen beams before, during and after bombardment with modestly intense beams of \\he4 and $^{16}$O.  We measured the helium profile and content Using (mostly Rutherford) Back-Scattering and demonstrated that these targets are robust against bombardment by modestly intense beams. No  degradation in the helium profile or content was observed after a prolonged bombardment with moderately intense beams. Such targets can be very useful for certain studies with modestly intense beams, including a measurement of the \\xhe3ag reaction at energies as low as $E_{cm}$ = 400 keV."
    },
    "0608/astro-ph0608628_arXiv.txt": {
        "abstract": "Hardness-intensity diagrams (HIDs) have been used with great success to study the accretion states and their connection to radio jets in X-ray binaries (XRBs).  The analogy between XRBs and active galactic nuclei (AGN) suggests that similar diagrams may help to understand and identify accretion states in AGN and their connection to radio loudness. We construct ``disc-fraction luminosity diagrams'' (DFLDs) as a generalization of HIDs, which plot the intensity against the fraction of the disc contribution in the overall spectral energy distribution (SED).  Using a sample of 4963 Sloan Digital Sky Survey (SDSS) quasars with ROSAT matches, we show empirically that an AGN is more likely to have a high radio:optical flux ratio when it has a high total luminosity or a large non-thermal contribution to the SED.  We find that one has to consider at least two-dimensional diagrams to understand the radio loudness of AGN.  To extend our DFLD to lower luminosities we also include a sample of low-luminosity AGN.  Using a simulated population of XRBs we show that stellar and supermassive BHs populate similar regions in the DFLD and show similar radio/jet properties.  This supports the idea the AGN and XRBs have the same accretion states and associated jet properties. ",
        "introduction": "In spite of a vast difference of scales, both black-hole X-ray binaries (XRBs) and active galactic nuclei (AGN) are thought to be powered by an essentially scale-invariant central engine consisting of a black hole, an accretion disc surrounded by a corona and a relativistic jet \\citep[e.g.,][]{ShakuraSunyaev1973,Antonucci1993,MirabelRodriguez1999}. In the study of AGN, it is a long-standing puzzle that some objects show strong radio emission while in others only little radio emission is found, even though their optical properties are similar \\citep{KellermannSramekSchmidt1989,MillerRawlingsSaunders1993}.  XRBs, i.e., stellar-mass accreting black holes, show states associated with strong radio emission as well as states during which the radio emission is strongly suppressed or ``quenched'' \\citep{TananbaumGurskyKellogg1972,FenderCorbelTzioumis1999,CorbelFenderTzioumis2000,Fender2001}. The different levels of radio emission have been ascribed to different jet properties of the different accretion states \\citep[e.g.,][]{FenderBelloniGallo2004}.  Due to the shorter timescales of accretion onto XRBs, it is possible to study the different accretion states and their association with jets in detail.  The quantitative comparison of XRBs and AGN has made a major step forwards in recent years with the discovery of a \\emph{fundamental plane of black-hole activity} \\citep{MerloniHeinzdiMatteo2003,FalckeKoerdingMarkoff2004}. This is a correlation between radio and X-ray luminosities and black hole mass which fits both classes of objects (both individually and as an ensemble).  Also the variability properties of AGN and XRBs show many similarities \\citep{UttleyMcHardyPapadakis2002,MarkowitzEdelsonVaughan2003,KoerdingFalcke2004,AbramowiczKluzniakMcClintock2004}. Therefore, it is a promising route to use the concepts and methods developed to study the evolution of radio jets in XRBs for exploring the puzzle of the radio loudness in AGN.  XRB outbursts seem to occur in cycles of several distinct states that are similar in every large outburst and from object to object (e.g., \\citealt{Nowak1995,BelloniHomanCasella2005},see Fig.~\\ref{fiHIDSketch}).  At the beginning of the cycle the source is in the \\emph{hard state} characterised by a hard power law in the X-ray spectrum. At radio frequencies, one ususally observes a steady jet \\citep{Fender2001}. In this state the standard accretion disc \\citep{ShakuraSunyaev1973} is probably truncated or unobservable for some other reason and the inner part of the accretion disc replaced by an inefficient accretion flow (e.g., an advection dominated accretion flows, ADAF: \\citealt{NarayanYi1994,ChenAbramowiczLasota1995,EsinMcClintockNarayan1997} but see also \\citealt{MillerHomanSteeghs2006}). The X-rays are ususally described with Comptonization models (e.g. \\citealt{ThornePrice1975,SunyaevTruemper1979}). However, it is also possible that the compact jet may contribute to the X-ray emission or even dominate it \\citep{MarkoffFalckeFender2001,MarkoffNowakWilms2005}. Once the source leaves the hard state, it is often found in an \\emph{intermediate state} (IMS). One finds two IMSs: first, the source enters a \\emph{hard IMS} characterised by a hard spectral component and band-limited noise in the power spectrum.  Also in this state a jet is observed in the radio. Later, a \\emph{soft IMS} is found, dominated by a soft spectral component with power-law noise in the power spectrum (e.g., \\citealt{BelloniHomanCasella2005}).  Approximately at the transition from the soft to the hard IMS one usually finds a radio flare and the ejecta can often be resolved in the radio band. In the soft IMS no radio emission is found. After leaving the IMS, the source may go to the \\emph{soft state}, where the X-ray spectrum is dominated by a soft thermal component that is thought to originate from a standard accretion disc. The radio jet is still quenched in this state \\citep{FenderCorbelTzioumis1999,CorbelFenderTzioumis2000}.  It has not been established definitely how the phenomenology of AGN can be mapped onto these XRB states. Quasars are AGN with bright optical continuum emission, which are thought to have high accretion rates and strong disc emission \\citep[e.g.,][]{UrryPadovani1995}. Here, we use the term AGN in a general sense to mean all supermassive accreting black holes, and we use the term ``quasar'' irrespective of the radio properties, i.e., synonymous with ``quasi-stellar object'' (QSO), and without imposing any luminosity cut.  For AGN, one commonly adopted definition of a ``radio-loud'' source is for the $R$ parameter, the ratio of the radio flux (at 1.4 GHz) to the optical $B$-band flux, to be larger than 10.  We follow this convention here. It has been known for over two decades that radio-loud quasars have higher X-ray fluxes that radio-quiet objects (e.g., \\citealt{ZamoraniHenryMaccacaro1981,ElvisWilkesMcDowell1994,ZdziarskiJohnsonDone1995}), perhaps hinting (with hindsight) that they represent different accretion states.  \\citet{PoundsDoneOsborne1995} suggested that narrow-line Seyfert 1 objects may be analogues of the soft state. The fundamental plane of accreting black holes has been interpreted as evidence that low-luminosity AGN (LLAGN) are hard-state objects \\citep{FalckeKoerdingMarkoff2004, KoerdingFalckeCorbel2005}. This has also been suggested by \\citet{Ho2005}, based on the lack of an observable ``big blue bump'' in the spectral energy distributions (SEDs) of LLAGN. Again using the fundamental plane, \\citet{MaccaroneGalloFender2003} suggested that AGN with an intermediate accretion rate (a few percent of the Eddington rate) are the analogue of soft-state XRBs.  On the other hand, \\citet{Boroson2002} has suggested that radio-loud and radio-quiet quasars occupy distinct regions of a two-dimensional diagram of physical accretion rate $\\dot{M}$ against Eddington-scaled luminosity $L/L_\\mathrm{Edd}$.  \\citeauthor{Boroson2002} considers these two quantities as the likely physical origin of the ``eigenvectors'' 1 and 2 \\citep{BorosonGreen1992}, and eigenvector 1 appears to be related to the radio properties.  The division between radio-loud and radio-quiet objects proposed by \\citet{Boroson2002} is one in black-hole mass, similar to those reported elsewhere \\citep{FranceschiniVercelloneFabian1998,LacyLaurent-MuehleisenRidgway2001,JarvisMcLure2002}. However, neither \\citet{Ho2002} nor \\citet{WooUrry2002b} found any evidence for a mass threshold between radio-loud and radio-quiet objects.  Thus, the physical mechanism making some sources much more radio-loud than others remains a puzzle.  For X-ray binaries, hardness-intensity diagrams (HIDs) have been used with great success to study the evolution of outbursts and separate the different accretion states (e.g, \\citealt{HomanWijnandsBelloni2001,FenderBelloniGallo2004,BelloniHomanCasella2005}). The radio-loud and radio-quiet phases of a source are well separated on this HID. Additionally, the timing properties of the X-ray lightcurve change dramatically with the position in the HID (e.g., \\citealt{HomanWijnandsBelloni2001,BelloniHomanCasella2005,Remillard2005}). Hence, the construction of equivalent diagrams promises to be a useful tool for the study of AGN, too.  For XRBs, X-ray HIDs are useful diagnostic tools because spectral features of both the accretion disc and the hard power-law component are visible in the X-rays. The blackbody temperature of the accretion disc scales with black hole mass as $M^{-1/4}$ so that the SEDs of accretion discs around supermassive black holes peak in the optical/ultra-violet and an X-ray HID for AGN would not contain any information about the disc emission. Therefore, we have to construct a new type of diagram which is applicable to the study of XRBs as well as of AGN. In this paper, we present the definition of an analogous diagram for AGN (\\S\\ref{s:method}) and use our method to explore the origin of radio loudness in a sample of quasars from the Sloan Digital Sky Survey (SDSS) and LLAGN (\\S\\ref{s:results}), also considering the impact of selection effects and other biases (\\S\\ref{s:critical}).  We compare the behaviour of AGN and X-ray binaries in \\S\\ref{s:disc} and conclude in \\S\\ref{s:conc}. ",
        "conclusions": "\\label{s:conc}  We have shown that the hardness-intensity diagrams used to diagnose the states of XRBs can be generalised to disc-fraction/luminosity diagrams (DFLDs), which can be used both for AGN and XRBs. The shape of the distribution of low-luminosity AGN and SDSS quasars in this diagram is similar to that of XRBs. Additionally, supermassive and stellar-mass black holes show a similar dependence of radio loudness on position in the DFLD. This suggests that radio loudness can only be understood in context of a two-dimensional diagram, i.e., as function of non-thermal fraction and luminosity, and is not directly related to a single variable like the accretion rate or black-hole mass. The similarity of the diagrams for AGN and XRBs supports the idea that AGN have the same accretion states as XRBs.  At low luminosities, AGN and XRBs are always in their hard state, which is characterized by strong X-ray and radio emission compared to their disc emission. At higher luminosities, AGN can either be in their radio-loud hard IMS (radio-loud quasars) or their soft state or soft IMS (radio-quiet quasar). The position of an AGN in the DFLD determines the probability with which it is in each of these states. Thus, we find the same coupling between accretion flow/disc and jet properties in AGN as in X-ray binaries.  To avoid all evolutionary and selection effects, it will be desirable to repeat our investigation with a sample of AGN from a single narrow redshift range, but with high dynamic range in luminosity and well-determined masses.  Ideally, the sample would extend down to LLAGN, which can to date only be found in the nearby universe, but not in the same volume as large numbers of quasars.  Further strong tests of the AGN-XRB unification and our proposed equivalence between XRB states and members of the AGN zoo will be performed by looking for correlations between DFLD position and other properties, such as QPO frequencies."
    },
    "0608/astro-ph0608302_arXiv.txt": {
        "abstract": "We present a measurement of the D/H abundance ratio in a metal-poor damped \\lya\\ (DLA) system along the sightline of QSO \\object.  The DLA system is at redshift $z = 2.70262$, has a neutral column density of \\lnhobj\\ \\cmm, and a gas-phase metallicity [O/H]$= -1.49$ which indicates that deuterium astration is negligible. Deuterium absorption is observed in multiple Lyman series with a column density of \\lndobj\\ \\cmm, best constrained by the deuterium Lyman-11 line. We measure \\lndhobj, which when combined with previous measurements along QSO sightlines gives a best estimate of \\lndhall\\, where the 1$\\sigma$ error estimate comes from a jackknife analysis of the weighted means. Using the framework of standard big bang nucleosynthesis, this value of D/H translates into a baryon density of \\dhobh, where the error terms represent the $1\\sigma$ errors from D/H and the uncertainties in the nuclear reaction rates respectively. Combining our new measurement with previous measurements of D/H, we no longer find compelling evidence for a trend of D/H with \\nhi. ",
        "introduction": "For the last decade, measurements of the primordial D/H ratio in QSO sightlines have provided increasingly more precise constraints on the cosmological baryon density.  Although the measurement of D/H is simple in principle, compared to the other light elements produced during big bang nucleosynthesis, finding those QSO absorption lines systems which are suitable for measuring D/H has proven observationally challenging \\citep{tytler_review}.  For a QSO absorption system to show D/H, a number of criteria must be met (see \\citet{kirkman03}, hereafter K03,  for a more detailed discussion).  First, the hydrogen column density must be large enough (since D/H is of the order of one part in $10^5$) such that deuterium can be observed using modern high-resolution spectrographs.  Second, the velocity structure of the hydrogen absorption must be simple enough, ideally a single component of gas, so that the deuterium absorption is well resolved given the small 82 \\kms\\ offset from the hydrogen Lyman lines. Third, there can be little to no interloping \\lya\\ forest or metal lines at the position of the deuterium absorption, since such absorption strongly complicates attempts to constrain the deuterium column density.  Unfortunately, \\lya\\ forest absorption is both ubiquitous and stochastic in high redshift QSO spectra.  Finally, the background QSO must be bright enough to obtain high signal-to-noise, high-resolution spectroscopy at $\\lambda < 4000$\\AA\\ with a reasonable allotment of telescope time.  Each one of these criteria act to decrease the probability that a D/H measurement can be made towards any given QSO, and since \\textit{all} the criteria must be met, the resultant probability of a QSO sightline being suitable for measuring D/H is very low, with only approximately 1\\% of QSOs at $z \\simeq 3$ able to provide a measurement of D/H.  To date, there are few measurements of D/H in QSO spectra (\\citet{bt98a},\\citet{bt98b}, \\citet{omeara01}, \\citet{pettinibowen}, K03, \\citet{levshakovDH}, \\citet{crighton04}).  These measurements constrain the baryon density \\obh\\ through the framework of standard big bang nucleosynthesis (SBBN), which predicts the abundances of the light elements as a function of the baryon--to--photon ratio $\\eta$ and the expansion rate of the universe \\citep{kolb}, and through the cosmic microwave background radiation, which provides the photon density.  A measurement of the ratio of any of the light nuclei produced in SBBN gives the baryon density, and  measurement of additional abundance ratios test the theory (see \\citet{steigman05},  \\citet{pettinidh}, and references therein for a current census of D/H and the other light element abundances).   Recent measurements of the temperature angular power spectrum of the CMB (Spergel \\etal\\ 2006) also provide a measurement of \\obh\\, depending on the assumptions made, with a level of accuracy roughly equal to or greater than that provided by D/H. Nevertheless, measurements of D/H are still important for a number of reasons. First, primordial D/H probes the universe at one of the earliest times in the universe accessible with current observational and theoretical techniques. Second, the light element abundances predicted from SBBN do not all agree with each other; most notably the observationally inferred abundance of \\lisv\\ is significantly lower than that expected from SBBN and D/H \\citep{fields06}. Third, D/H can help constrain deviations from SBBN, such as inhomogeneous BBN \\citep[e.g.][]{lara06}, relic primordial particle decays \\citep[e.g.][]{jedamzik04}, or non-standard neutrino physics \\citep[e.g.]{abazajian05}. Fourth, the dispersion in the measurements of D/H is larger than would be expected from the individual measurement errors (K03), i.e., the data demand both a better understanding of the errors on the current measurements and new constraints. Fifth, the value of \\ob\\ derived from D/H and SBBN requires many fewer priors than the CMB derived value.  Moreover, the \\ob\\ from D/H can be used in principle as a prior in the CMB analysis, and the ratio of the values for \\ob\\ from D/H and the CMB offer a precision test of the hot Big Bang model. Finally, the existing measurements of D/H show evidence of a trend of decreasing D/H with increasing \\nhi. K03 suggested that this trend (and the dispersion in D/H values) is due to error under-estimation.  Fortunately, the Sloan Digital Sky Survey \\citep[SDSS,][]{sdssdr4}, by virtue of its large sample of high redshift QSO spectra (the Data Release 4 alone contains 5036 QSOs with $z > 2.7$) gives us a new data set to find those special sightlines which can show D/H. In this \\textit{Letter}, we present a new measurement of D/H in a QSO sightline from the SDSS, \\object, which was chosen as part of our high-resolution survey for Lyman limit absorption \\citep{omeara06}. ",
        "conclusions": "We now discuss the value of D/H we obtain for the absorber and place it within the context of the combined D/H ratio for all QSO absorption systems. The best estimates of the \\ion{H}{1} and \\ion{D}{1} column densities in the DLA towards \\object\\ imply a value of \\lndhobj. The errors on D/H stem primarily from the effect of continuum placement uncertainty on the H~I column density, and the signal-to-noise ratio of the data at Lyman-11 where the \\ion{D}{1} column density is best constrained.  Turning now to the combined D/H value from QSO sightlines, Figure \\ref{fig_alldh} shows the new value of D/H from \\object\\ along with the previous values of D/H taken from the sample discussed in K03.  We do not include the result of Crighton \\etal\\ (2004), since we feel that the errors on D/H in this system have been under-estimated, particularly for the reported \\nhi\\ value. We do not include the results of \\citet{levshakovDH} for the reasons given in K03. The horizontal solid line shows the value for the weighted mean of the data \\lndhallfull, which we round to $\\log \\rm{D/H} = -4.55 \\pm 0.04$ to keep consistent with the literature, and the dashed lines show the $\\pm 1 \\sigma$ uncertainties estimated from a jackknife analysis of the weighted means. In the case of asymmetric errors on individual D/H measurements, we have adopted the larger of the errors for the calculation of the weighting.  Prior to the addition of \\object\\ to the sample of D/H measurements of K03, a $\\chi^2$--minimizing linear fit to the data of the form $\\log \\rm{D/H} = -2.914 \\pm 0.467 - (0.087 \\pm 0.025)\\times$\\lnhi\\ provided an acceptable fit to the data ($P_{(\\chi^2 > \\chi^2_{\\rm{fit}})} = 0.74$).  With the inclusion of \\object, however, we see a significant decrease in the likelihood that there is a D/H trend with \\nhi. Although the data are best fit with a non--zero slope, $\\log \\rm{D/H} = -3.707 \\pm 0.385 - (0.044 \\pm 0.021)\\times$\\lnhi, the slope differs from zero at only the $2\\sigma$ level (even before accounting for the presence of \\nhi\\ in both axes). Furthermore, this is not a good model of the data, with $P_{(\\chi^2 > \\chi^2_{\\rm{fit}})} = 0.04$.  As such, the data give little confidence to the existence of a trend of D/H with \\nhi. Finally, if we were to include the \\cite{crighton04} and \\cite{levshakovDH} results, the likelihood that D/H depends on \\nhi\\ is further diminished.  The presumption of a single value for D/H, however, is still not supported by the observations; the observed scatter exceeds that expected assuming the error estimates reported in the literature.  Our adopted error on the best estimate for D/H from the jackknife estimation exceeds the error on the weighted mean by a factor of two. Likewise, the $\\chi^2$ of the six measurements of D/H about the weighted mean value of \\lndhall\\ is high, with $P_{(\\chi^2 > \\chi^2_{\\rm{D/H}})} = 0.01$.  Although there is the possibility that the scatter in the individual D/H measurements is real, we prefer the hypothesis of K03 that the errors in some of the individual measurements, if not all of them, are underestimated.  It is likely that a combination of new methods of error analysis and new QSO  sightlines is required to fully address the excess scatter in the D/H measurements.   Using the framework from SBBN \\citep{bnt2001}, the value for D/H of \\lndhall\\ translates to a value for the cosmological baryon density of \\dhobh, where the first error term comes from the errors on D/H explained above, and the second term from the uncertainties in the nuclear reaction rates.  By comparison, the WMAP three year result provides an estimate of \\obh $=0.0223^{+0.0007}_{-0.0009}$, which lies  within the $1\\sigma$ error estimate on \\obh\\ from D/H (Spergel \\etal\\ 2006).   Finally, we note that the absorber showing D/H in the spectrum of \\object\\ was discovered serendipitously as part of our survey for Lyman limit absorption, and is the first D/H measurement from a QSO first discovered by the SDSS. The SDSS Data Release 3 alone has 405 DLA with redshifts optimal for D/H \\citep[$2.51 \\le z \\le 4.0$][]{dla_dr3}, Assuming a small fraction of these DLA provide measurements of D/H, the SDSS will give many tens of measurements.  The situation improves further if one considers the Lyman limit systems toward the SDSS QSO sample. This contrasts with the SLLS and DLA which give the \\nhi\\ from the \\lya\\ and \\lyb\\ lines, and the \\ndi\\ from the unsaturated D~I lines in the Lyman series.  Because of this effect, the DLA and SLLS offer more pathlength per QSO to potentially find D/H.  All of these effects combine to give a likely distribution of D/H measurements which is roughly independent of \\nhi, a hypothesis which is already being hinted at in the current sample, since two measurements come from LLS, two from SLLS, and two from DLA. Altogether, the SDSS offers the best opportunity for investigating the larger than expected scatter in D/H and correlations with \\nhi, metallicity, etc."
    },
    "0608/astro-ph0608134_arXiv.txt": {
        "abstract": "{We report on the physical properties of the disk-like structure of B8\\,IIIp star 36\\,Lyncis from line syntheses of phase-resolved, high resolution spectra obtained from the {\\it International Ultraviolet Explorer} archives and from newly obtained ground-based spectra of the H$\\alpha$ absorption profile. This disk is highly inclined to the rotational axis and betrays its existence every half rotation cycle as one of  two opposing sectors pass in front of the star. Although the disk absorption spectrum is at least ten times too weak to be visible in optical iron lines during these occultations, its properties can be readily examined in a large  number of UV ``iron curtain\" lines because of their higher opacities. The UV Fe\\,II and Fe\\,III lines in particular permit a determination of the disk temperature: 7,500$\\pm{500}$\\,K and a column density of 3$_{-1.5}$$\\hspace*{-.25in}^{+3}$$\\hspace*{.1in}\\times$10$^{20}$\\,cm$^{-2}$. The analysis of the variations of the UV resonance lines brings out some interesting details about the radiative properties of the disks: (1) they are optically thick in the C\\,IV and Si\\,IV doublets, (2) the range of excitation of the UV resonance lines is larger at the primary occultation ($\\phi$ = 0.00) than at the secondary one, and (3) the {\\bf relative strengths of the absorption peaks} for the two occultations varies substantially from line to line. We have modeled the absorptions of the UV C\\,IV resonance and H$\\alpha$ absorptions by means of a simulated disk with opaque and translucent components. Our simulations suggest that a gap separates the star and the inner edge of the disk. The disk extends radially out to $\\geq$10 R$_{*}$. The disk scale height perpendicular to the plane is $\\approx$1R$_{*}$. However, the sector causing the primary occultation is about four times thicker than the opposite sector.  The C\\,IV scattering region extends to a larger height than the H$\\alpha$ region does, probably because it results from shock heating far from the cooler disk plane. \\keywords {stars: individual: 36\\,Lyncis -- stars: chemically peculiar -- stars: magnetic fields -- stars: circumstellar matter -- ultraviolet: stars -- X-rays: stars}} ",
        "introduction": "In a companion to this paper, Wade et al. (2006; hereafter Paper 1) have redetermined the magnetic and rotational periods and surface abundances of the He-weak B8p star 36\\,Lyncis (HD\\,79158). Using data from the {\\it International Ultraviolet Explorer (IUE)} satellite, Shore \\& Brown (1990; hereafter SB90) discovered the presence of circumstellar plasma through the periodic modulations of the strengths of its C\\,IV and Si\\,IV resonance lines. Following the model of Shore (1987), SB90 found that these modulations are due to a disk-like structure surrounding the star.\\footnote {The term ``disk\" will be used in this paper to describe circumstellar matter corotating in or near the magnetic equatorial plane of the star. Evidence is accumulating that the gas density of these structures is unevenly distributed in azimuth in this plane. The structure may have a discrete inner edge, may be decentered, may be nonisothermal around the disk {\\bf at a given radius from the star's center}, and may be warped out of coplanarity (see e.g., Townsend \\& Owocki 2005; hereafter TO05).} According to this picture, the wind from the magnetic poles flows out along open field lines to infinity, much as for unmagnetized stars. However, at lower magnetic latitudes the wind behavior and geometry is complex. Since the wind particles from this zone are charged, they cannot easily cross the magnetic loops, and the high velocity flow is effectively quenched. At intermediate latitudes, the flow is constrained to follow closed magnetic loops. As they cross the magnetic equator, the particles collide and shock with streams originating from the other hemisphere, generating EUV or X-ray emission in a large, spheroidal post-shock region (e.g., Babel 1997, Babel \\& Montmerle 1997 (hereafter BM97), Babel \\& North 1997). Post-shock particles cool and settle toward the equatorial plane. Steady state is achieved as particles return to the star's surface or escape through the disk's outer edge. In this picture a cool disk co-rotates with the star and resembles a wobbling inner tube. As fixed external observers, we see ``shell\" absorption components in the UV resonance and H$\\alpha$ lines near the rest frame at those times during the rotation cycle when a section of the disk {\\bf transits} the star. In spectra of hot Bp stars a slightly redshifted emission is observed in the UV lines when the disk is viewed ``face on.\"  Since the pioneering studies of Babel and collaborators, several important advances have been made in the hydrodynamic modeling of Bp disks. The first of these was the inclusion of the feedback effects of the wind particle weight on the magnetic field loops (ud-Doula \\& Owocki 2002, uO02). This weight causes the loops to sag, thereby permitting disk particles to fall back to the star at low magnetic latitudes. This process leaves an effectively gas free region close to the star. In addition, for sufficiently high ratios of magnetic to wind energy density $\\eta$, circulation of wind efflux proceeds on a timescale of a few days and, in certain $\\eta$ regimes, chaotically. These models suggest that the effective coverage factor of the star by the disk can be large, even occulting the entire star for near edge-on viewing angles.  Models that include radiative cooling of the wind/disk shock interface (Tonneson et al. 2002) indicate that the hot zone is mainly confined to the outer region of the disk yet also extends out of the equatorial plane. Recent theoretical work (e.g., Preuss et al. 2004, TO05) suggests that matter can accumulate on low potential energy surfaces, well outside the disk plane.     A large number of studies have shown that many and perhaps all Bp stars with measurable dipolar fields harbor co-rotating magnetospheres. Periodic ultraviolet variations have also been found in a few $\\beta$\\,Cep stars (Henrichs et al. 1998, Floquet et al. 2002, Neiner et al. 2003), in the O7\\,V star $\\theta^1$\\,Ori\\,C  (Walborn \\& Nichols 1994, Stahl et al. 1996), and in He-weak stars as late as type B8 (Shore \\& Brown 1990b, Shore et al. 2004). To generalize, the periodic variability of the C\\,IV and/or N\\,V lines is a hallmark of Bp stars with oblique dipolar magnetic fields (Smith \\& Groote, hereafter SG01).  36\\,Lyn is one such star. It lies near the cool edge of the He-weak Bp domain and shows prominent C\\,IV and N\\,IV line variations. In order to take advantage of the available H$\\alpha$ observations and {\\it IUE} archival data, the improved determination of its rotational and magnetic geometry, and of our improved understanding of the theoretical properties of disks, we have modeled the variations of several UV and optical lines in this star's disk. From our results we are able to provide quantitative estimates of its thermodynamic and geometrical properties.  In their analysis of disk properties of four Bp and $\\beta$\\,Cep stars, SG01 utilized archival spectra from the {\\it International Ultraviolet Explorer} satellite to analyze the time variations of UV ultraviolet resonance lines and the ``iron curtain\" background of weak lines populating the far-UV spectrum. For this study we were able to include new phased H$\\alpha$ absorptions, as taken from Paper 1; see the journal of observations in this paper for details. The finer {\\bf sampling} in phase and the higher quality of these observations over the {\\it IUE} data provide additional constraints on the properties of the circumstellar material.  The UV data for these programs are extant high-dispersion {\\it IUE} echellograms obtained through the large-aperture and of the Short Wavelength Prime (SWP) camera. These data were obtained from the MAST archives.\\footnote{Multi-Mission Archive at Space Telescope Science Institute, in contract to NASA.} and constitute a total of 26 spectra (a 27th spectrum, SWP\\,32982, was lost to the {\\it IUE} project; Levay 2003).  \\begin{table}[ht!] \\begin{center}{}  \\begin{tabular}{r|l}  \\hline  \\noalign{\\smallskip}  Mass  $M$& $4.0\\pm 0.2 M_\\odot$\\\\  Radius $R$& $3.4\\pm 0.7 R_\\odot$\\\\  Effective temperature $T_{\\rm eff}$          & $13,300\\pm 300$ K\\\\  Surface gravity $\\log g$                  & $3.7-4.2$\\\\  Period $P_{\\rm rot}$      & 3.83475 days\\\\  Projected rotation speed $v_{\\rm e}\\sin i$     & $48\\pm 3$ km\\,s$^{-1}$ \\\\  Inclination $i$           & $56-90^{o}$ \\\\  Dipole strength $B_{\\rm d}$        & 3210-3930 G\\\\  Phase of mag. equator passage & 0.0 \\\\  Photospheric abundance [Fe]   & +0.8-1.1 dex \\\\  \\hline\\hline\\noalign{\\smallskip}  \\end{tabular} \\caption{\\label{tab1}{Some characteristics of 36 Lyn from Paper 1.} }   \\end{center}   \\end{table} ",
        "conclusions": "We can recap the conclusions of this paper as follows:  \\begin{enumerate}  \\item Like the circumstellar disks of early-type Bp stars, the static ``disk\" of 36\\,Lyn consists of both a cool and a hot component. In general terms, this confirms the hydrodynamic result of Babel \\& Montmerle (1997) for IQ\\,Aur. The cool component of 36\\,Lyn is observed in the Fe-group lines (mainly Fe$^{2+}$) arising from low excitation states; the hot component is visible in the UV resonance lines. The volume in which H$\\alpha$ absorption is formed probably overlaps both these regions, but we do not know the extent of this overlap.  \\item We have determined the temperature of the cool disk component to be 7,500$\\pm{500}$\\,K from an analysis of Fe\\,II and Fe\\,III lines variations during primary disk occultation. The temperature ratio $T_{disk}$/$T_{eff}$ $\\approx$0.6 for 36\\,Lyn. This is {\\bf also} the ratio SG01 found for the hot Bp stars $\\sigma$\\,Ori\\,E and $\\beta$\\,Cep.  \\item The physical parameters of this disk, including its column density, insure that the strongest low-excitation lines are readily visible in {\\it IUE} spectra, but they are not quite sufficient to allow the disk component to be visible in even strong optical Fe lines.  \\item The ionization temperature(s) determined from the presence of disk N\\,V and C\\,IV absorptions reveal that part of the disk is heated, perhaps to 30,000\\,K or more. For C\\,IV there is only a hint of emission at non-occultation phases, unlike the disks of hotter Bp stars.  \\item The relative heights of the absorption maxima corresponding to the occultations by two disk segments differ among the highly excited ``super ions\" (N\\,V, C\\,IV, and Si\\,IV) themselves and with respect to resonance lines of less excited ions. The two opposing disk segments do not exhibit the same ranges of ionization.  The disk sector responsible for absorptions at $\\phi$ = 0.00 has a more restricted temperature range than the sector visible at $\\phi$ $\\approx$ 0.55. The existence of restricted temperatures implies that the production of the ``super ions\" does not result from the irradiation of gas by local soft X-rays.   \\item Geometric modeling of the occultation from the UV resonance and H$\\alpha$ lines suggests that the inner edge of the disk (ring) is detached from the star's surface. The inner edge is is located at $\\approx$ ${\\frac 13}$$r_o$.  This conclusion is secure for H$\\alpha$ but is less secure for the UV data.  \\item Analysis of the equivalent width maxima of the $H\\alpha$ shows clearly that the associated opposing disk sectors have nearly the same inner and outer disk values. They differ by a factor of 3 in their heights above the magnetic plane. The secondary absorption appears to be weaker because the disk sector responsible for it is thinner (measured out of the disk plane) than the opposing segment by about the same factor.  \\item Like H$\\alpha$, the UV resonance lines are opaque near the central plane. Their absorptions are formed in volumes having nearly the same range of radii as the H$\\alpha$ absorption (i.e., from 3R$_*$ to 10R$_*$). The UV resonance lines indicate significant additional translucent extensions beyond this {\\bf height,} with  h$_{S}$/$h$ $\\approx$ 3.  \\end{enumerate}"
    },
    "0608/astro-ph0608244_arXiv.txt": {
        "abstract": "We propose a methodology to perform a self-consistent analysis of the physical properties of the emitting gas of HII galaxies adequate to the data that can be obtained with the XXI century technology. This methodology requires the production and calibration of empirical relations between the different line temperatures that should superseed currently used ones based on very simple, and poorly tested,  photo-ionization model sequences.  As a first step to reach these goals we have obtained simultaneous blue to far red longslit spectra with the William Herschel Telescope (WHT) of three compact \\HII\\ galaxies selected from the  Sloan Digital Sky Survey (SDSS) Data Release 2 (DR2) spectral catalog using the INAOE Virtual Observatory superserver. Our spectra cover the range from 3200 to 10500\\,\\AA, including the Balmer jump, the [O{\\sc ii}]\\,$\\lambda\\lambda$\\,3727,29\\,\\AA\\ lines, the [S{\\sc iii}]\\,$\\lambda\\lambda$\\,9069,9532\\,\\AA\\ doublet as well as various weak auroral lines such as [O{\\sc iii}]\\,$\\lambda$\\,4363\\,\\AA\\ and [S{\\sc iii}]\\,$\\lambda$\\,6312\\,\\AA.  For the three objects we have measured at least four line temperatures: T([O{\\sc iii}]), T([S{\\sc iii}]), T([O{\\sc ii}]) and T([S{\\sc ii}]) and the Balmer continuum temperature T(Bac). These measurements and a careful and realistic treatment of the observational errors yield total oxygen abundances with accuracies between 5 and 9\\%. These accuracies are expected to improve as better calibrations based on more precise measurements, both on electron temperatures and  densities, are produced.  We have compared our obtained spectra with those downloaded from the SDSS DR3 finding a satisfactory agreement. The analysis of these spectra yields values of line temperatures and elemental ionic and total abundances which are in general agreement with those derived from the WHT spectra, although for most quantities they can only be taken as estimates since, due to the lack of direct measurements of the required lines, theoretical models had to be used whose uncertainties are impossible to quantify.  The ionization structure found for the observed objects from the O$^{+}$/O$^{2+}$ and S$^{+}$/S$^{2+}$ ratios points to high values of the ionizing radiation, as traced by the values of the ``softness parameter\" $\\eta$ which is less than one for the three objects. The use of line temperatures derived from T([O{\\sc iii}]) based on current photoionization models yield for the two highest excitation objects, much higher values of $\\eta$ which would imply lower ionizing temperatures. This is however inconsistent with the ionization structure as probed by the measured emission line intensities.  Finally, we have measured the T(Bac) for the three observed objects and derived temperature fluctuations. Only for one of the objects, the temperature fluctuation is significant and could lead to higher oxygen abundances by about 0.20\\,dex. ",
        "introduction": "\\HII\\ galaxies are low mass irregular galaxies with, at least, a recent episode of  violent star formation \\citep*{1985MNRAS.216..255M,1985RMxAA..11...91M} concentrated in a few parsecs close to their cores. The ionizing fluxes originated by these young massive stars dominate the light of this subclass of Blue Compact Dwarf galaxies (BCDs) which show  emission line spectra very similar to those of giant extragalactic  \\HII\\ regions \\citep[GEHRs; ][]{1970ApJ...162L.155S,1980ApJ...240...41F}. Therefore, by applying the same measurement techniques as for \\HII\\ regions, we can derive the temperatures, densities and chemical composition of  the interstellar gas in this type of generally  metal-deficient galaxies. In some cases, it is possible to detect in these objects, intermediate-to-old stellar populations which have a more uniform spatial distribution than the bright and young stellar  populations associated with the ionizing clusters \\citep{1998ApJ...493L..23S}. This older population produces a characteristic spectrum with absorption features which mainly affect the hydrogen recombination emission lines \\citep{1988MNRAS.231...57D}, that is the Balmer and Paschen series in the spectral range of interest. In some cases, the underlying stellar absorptions can severely affect the ratios of H{\\sc i} line pairs and hence the determination of the reddening constant (C(H$\\beta$)). They must therefore be measured with special care (see discussion in \\S \\ref{measure}).  A considerable number of the {blue} objects observed at intermediate redshifts seem to have properties {(mass, R$_e$, velocity width of  the emission lines) similar to local \\HII\\ galaxies} \\citep{1994ApJ...427L...9K,1995ApJ...440L..49K,1996ApJ...460L...5G,1998ApJ...495L..13G}. {In particular, those with $\\sigma$\\,$<$\\,65\\,km\\,s$^{-1}$ follow the same $\\sigma - L_B\\ {\\rm and}\\ L_{H\\beta}$ relation as seen in \\HII\\ galaxies \\citep{2000MNRAS.311..629M,2003ApJ...598..858M,2003RMxAC..16..213T,2005MNRAS.356.1117S}.} {Similar conclusion is drawn from recent studies on Lyman Break galaxies that also suggest that strong narrow emission line galaxies  might have been very common in the past \\citep[e.g.][]{2000ApJ...528...96P,2001ApJ...554..981P,2001A&A...379..393E}. To detect possible evolutionary effects like systematic differences in their chemical composition, accurate and reliable methods for abundance determination are needed.}  This is usually done by combining photoionization model results and observed emission line intensity ratios.  {There are  several major problems with this approach that limit the confidence of  present results.} Among them: the effect of temperature structure in multiple-zone models \\citep{2003MNRAS.346..105P}; the presence of temperature fluctuations across the nebula \\citep{2003ApJ...584..735P};  collisional and density effects on ion temperatures (Luridiana et al.\\ 1999, P\\'erez-Montero \\& D\\'{\\i}az 2003); the presence of neutral zones affecting the calculation of ionization correction factors (ICFs; \\nocite{2002ApJ...565..668P} Peimbert et al.\\ 2002); the ionization structure not adequately reproduced by current models \\citep{2003MNRAS.346..105P}; ionization vs.~matter bounded zones, affecting the low ionization lines formed in the outer parts of the ionized regions \\citep*{2002MNRAS.329..315C}. \\nocite{1999ApJ...527..110L} \\nocite{2005MNRAS.361.1063P} On the other hand, the understanding of the age and evolutionary state of \\HII\\ galaxies require the use of self-consistent models for the ionizing stars and the ionized gas. However, model computed evolutionary sequences show important differences with observations \\citep{2003A&A...397...71S}, including: (a) He{\\sc ii} is too strong in a substantial number of objects as compared to model predictions; (b) [O{\\sc iii}]/H$\\beta$ vs.\\ [O{\\sc ii}]/H$\\beta$ and [O{\\sc iii}]/H$\\beta$  vs.~[O{\\sc i}]/H$\\beta$ are not  well reproduced by evolutionary model sequences in the sense that predicted collisionally excited lines are too weak compared to observations; (c) there is a large spread in the [N{\\sc ii}]/[O{\\sc ii}] values (more than an order of magnitude) for galaxies with the same value of ([O{\\sc ii}]+[O{\\sc iii}])/H$\\beta$ in the metallicity range from 8 to 8.4 (see e.g. P\\'erez-Montero \\& D\\'iaz 2005). \\nocite{2002AJ....123..485S}  \\begin{table*} \\centering \\caption[]{Journal of observations.  Right ascension, declination and redshift were obtained from the SDSS.} \\label{jour} \\begin{tabular} {l c r r c c} \\hline \\multicolumn{1}{c}{Object  ID}   &  spSpec SDSS   & \\multicolumn{1}{c}{RA} & \\multicolumn{1}{c}{Dec} &  z  &    Exposure (s) \\\\ \\hline SDSS J002101.03+005248.1 & spSpec-51900-0390-445  &  5.254297 & 0.880038 &  0.098 & 1 $\\times$ 1200 + 2 $\\times$ 2400\\\\ SDSS J003218.60+150014.2 & spSpec-51817-0418-302  &  8.077479 & 15.003949 &  0.018 & 1 $\\times$ 1200 + 2 $\\times$ 2400\\\\ SDSS J162410.11-002202.5  & spSpec-52000-0364-187  &  246.042122 & -0.367378 &  0.031 & 1 $\\times$ 1200 + 3 $\\times$ 1800\\\\ \\hline \\end{tabular} \\end{table*}  \\begin{table*} \\centering \\caption[]{SDSS photometric magnitudes obtained using the DR4 explore tools.} \\label{obj} \\begin{tabular} {l c r r c c} \\hline \\multicolumn{1}{c}{Object  ID}   & \\multicolumn{1}{c}{u} & \\multicolumn{1}{c}{g} & \\multicolumn{1}{c}{r} & \\multicolumn{1}{c}{i}   &    \\multicolumn{1}{c}{z} \\\\ \\hline SDSS J002101.03+005248.1 & 17.56 & 17.35 & 17.51 & 16.98 & 17.45  \\\\ SDSS J003218.60+150014.2 & 17.04 & 16.49 & 16.53 & 16.74 & 16.65  \\\\ SDSS J162410.11-002202.5  & 17.07 & 16.46 & 16.91 & 16.80 & 16.74  \\\\ \\hline \\end{tabular} \\end{table*}  Substantial progress toward solving the problems listed above has to come from the accurate measurement of weak emission lines  which will allow to derive [O{\\sc ii}], [S{\\sc ii}] and [S{\\sc iii}] temperatures and densities allowing to constrain the ionization structure as well as Balmer and Paschen discontinuities which will provide crucial information about the actual values of temperature fluctuations. It is possible that these fluctuations produce the observed differences between the abundances relative to hydrogen derived from recombination lines (RLs) and collisionally excited lines (CELs) when a constant electron temperature is assumed \\citep{1967ApJ...150..825P,1969BOTT....5....3P,1971BOTT....6...29P}. { These discrepancies have been observed in a good sample of objects, such as galactic \\HII\\ regions \\citep[e.g.][and references therein]{2004MNRAS.355..229E,2005MNRAS.362..301G,2006MNRAS.368..253G}, \\HII\\ regions in the Magellanic Clouds \\citep[e.g.][and references therein]{2000ApJ...541..688P,2003ApJ...584..735P,2003MNRAS.338..687T}, extragalactic \\HII\\ regions \\citep[e.g.][and references therein]{2003RMxAC..16..113P,2005ApJ...634.1056P,2006astro.ph..3134G} and planetary nebulae \\citep[e.g.][and references therein]{2002MNRAS.334..777R,2005MNRAS.362..424W,2006MNRAS.tmp..404L,2006astro.ph..5082L,2006astro.ph..5595P}. Likewise, there are relatively recent theoretical works that study the possible causes of these discrepancies in abundance determinations using photoionization models of different complexity \\citep[e.g.][and references therein]{2005A&A...434..507S,2005A&A...444..723J,2005MNRAS.364..687T}.}  Unfortunately most of the available starburst and \\HII\\ galaxy spectra have only a restricted wavelength  range (usually from about 3600 to 7000 \\AA) and do not have the adequate S/N to accurately measure the intensities of the weak diagnostic emission lines. Even the Sloan Digital Sky Survey (SDSS) spectra (Stoughton et al.\\ 2002) do not cover simultaneously the [O{\\sc ii}]\\,$\\lambda\\lambda$\\,3727,29 and the [S{\\sc iii}]\\,$\\lambda$\\,9069\\,\\AA\\ lines. We have therefore undertaken a project with the aim of obtaining a database of top quality line ratios { for a sample that includes the best objects for the task. The data is collected using exclusively two arm  spectrographs in order to guarantee both high quality spectrophotometry in the whole spectral range from 3500 to 10500\\,\\AA\\ and good spectral and spatial resolution.} In this way we will be in a  position to vastly improve constraints on the photoionization models including the  mapping of the ionization structure and the measurement of temperature fluctuations about which very little is known.  In this first work we present observations of three \\HII\\ galaxies selected from the SDSS. Details regarding the selection of the objects as well as the observations are given in Section 2. Section 3 presents the results including line measuring techniques. The methodology for the derivation of gaseous physical conditions and elemental abundances is presented in Sections 4 and 5 respectively. The discussion of our results, including a detailed comparison with the SDSS data is presented in Section 6. Finally, Section 7 summarizes the main conclusions of our work. ",
        "conclusions": "We have performed a detailed analysis of newly obtained spectra of three \\HII\\ galaxies selected from the Sloan Digital Sky Survey Data Release 2, covering from 3200 to 10550 \\AA\\ in wavelength. For the three objects we have measured four line temperatures: T([O{\\sc iii}]), T([S{\\sc iii}]), T([O{\\sc ii}]) and T([S{\\sc ii}]) and the Balmer continuum temperature T(Bac). For two of the objects we have also measured  T([N{\\sc ii}]). These measurements and a careful and realistic treatment of the observational errors yield total oxygen abundances with accuracies between 5 and 9\\%. The fractional error is as low as 2\\% for the ionic O$ ^{2+} $/H$ ^{+} $ ratio due to the small errors associated with the measurement of  the strong nebular lines of [O{\\sc iii}] and the derived T([O{\\sc iii}]), but increases to 15\\% for the O$^{+}$/H$^{+}$ ratio. The accuracies are lower in the case of the abundances of sulphur (of the order of 20\\% for S$^+$ and 12\\% for S$^{2+}$) due to the presence of larger observational errors both in the measured line fluxes and the derived electron temperatures. The error increases further (up to 30\\%) for the total abundance of sulphur due to the uncertainties in the ICF.  This is in contrast with the small errors quoted for line temperatures other than T([O{\\sc iii}]) in the literature, in part due to the commonly assumed methodology of deriving them from the measured T([O{\\sc iii}]) through theoretical relations. These relations are found from photoionization models and no uncertainty is attached to them; therefore, the obtained line temperatures carries only the observational error found for the T([O{\\sc iii}]) measurement. If this methodology were to be adopted, the theoretical relations should be adequately tested and empirical relations between the relevant line temperatures should be obtained in order to quantify the corresponding model uncertainties.  We have compared our obtained spectra with those downloaded from the SDSS DR3 finding a satisfactory agreement. The analysis of these spectra yields values of line temperatures and elemental ionic and total abundances which are in general agreement with those derived from the WHT spectra, although for most quantities, they can only be taken as estimates since, due to the lack of direct measurements of the required lines, theoretical models had to be used whose uncertainties are impossible to quantify. Unfortunately, the spectral coverage of SDSS precludes the simultaneous observation of the [O{\\sc ii}]\\,$\\lambda\\lambda$\\,3727,29\\,\\AA\\ and [S{\\sc iii}]\\,$\\lambda\\lambda$\\,9069, 9532\\,\\AA\\ lines, and therefore the analysis can never be complete.  The ionization structure found for the observed objects from the O$^{+} $/O$ ^{2+} $ and S$^{+} $/S$ ^{2+} $ ratios points to high values of the ionizing radiation as traced by the values of the ``softness parameter\" $\\eta$ which is less than one for the three objects. Line temperatures derived from T([O{\\sc iii}]) with the use of current photoionization models yield for the two highest excitation objects,  much higher values of $\\eta$ which would imply lower ionizing temperatures. This is however inconsistent with the ionization structure as probed by the measured emission line intensities.  Finally, we have measured the Balmer continuum temperature for the three observed objects and derived the temperature fluctuations as defined by \\cite{1967ApJ...150..825P}. Only for one of the objects, the temperature fluctuation is significant and could lead to higher oxygen abundances by about 0.20\\,dex."
    },
    "0608/astro-ph0608072_arXiv.txt": {
        "abstract": "We searched for redshifted O emission lines from the possible warm-hot intergalactic medium (WHIM) surrounding the cluster of galaxies A2218 at $z=0.1756$ using the XIS instrument on {\\it Suzaku}. This cluster is thought to have an elongated structure along the line of sight based on previous studies. We studied systematic uncertainties in the spectrum of the Galactic emission and in the soft X-ray response of the detectors due to the contamination building up on the XIS filters. We detected no significant redshifted O lines, and set a tight constraint on the intensity with upper limits for the surface brightness of O\\emissiontype{VII} and O\\emissiontype{VIII} lines of $1.1\\times10^{-7}$ and $3.0\\times10^{-7}$~$\\mathrm{photons~cm^{-2}~s^{-1}~arcmin^{-2}}$, respectively.  These upper limits are significantly lower than the previously reported fluxes from the WHIM around other clusters of galaxies.  We also discuss the prospect for the detection of the WHIM lines with {\\it Suzaku} XIS in the future. ",
        "introduction": "Based on several N-body simulations of cosmological large-scale structure formation \\citep[e.g.,][]{cen99:_where_are_baryon, dave01:_baryon_warm_hot_inter_medium,2003ApJ...594...42C}, a significant (30--50\\%) fraction of baryons is thought to reside in the form of gas in a `warm-hot' phase ($T=10^{5-7}$~K), which is hard to detect with the instruments currently in operation.  This warm-hot gas, whose density is $10^{-6}$--$10^{-4}\\mathrm{~cm^{-3}}$, is called the warm-hot intergalactic medium (WHIM)\\@.  The firm detection of the WHIM is important because it is the most promising candidate for the ``missing baryons''; i.e., it is thought to explain the discrepancy between the baryon density observed in the local universe \\citep{fukugita98:_cosmic_baryon_budget} and that in the distant universe \\citep{rauch97:_opacit_ly_fores_implic_omega} or that calculated from the observed fluctuations of the cosmic microwave background \\citep{spergel03:_first_year_wilkin_microw_anisot}.  The WHIM may be detected via emission or absorption lines from highly ionized elements in UV or X-ray spectra. Given the relative elemental abundances, --- on the order of 0.1 to 0.3 of Solar --- oxygen is the most promising element to provide detectable transitions at the expected temperatures. Over 40 systems show O\\emissiontype{VI} absorption features, based on observations from {\\it FUSE} and {\\it HST} \\citep[e.g.][]{2005ApJ...624..555D}. However, their analysis shows that the baryon density in the temperature range $T=10^{5-6}$~K is about an order of magnitude lower than the expected WHIM level. Therefore, most of the missing baryons seem to reside in a hotter phase at $T=10^{6-7}$~K, which can be probed with X-rays. There is only one sightline with a likely detection of the WHIM via O\\emissiontype{VII} or O\\emissiontype{VIII} absorption lines so far \\citep{2005ApJ...629..700N}. Their results suggest that most of the missing baryons can be explained by the absorbing clouds inferred from their data. However, recent work by \\cite{2006astro.ph..4515R} and \\cite{2006astro.ph..4519K} showed that the detection with the Chandra LETGS by \\cite{2005ApJ...629..700N} was not consistent with RGS observations of the same sight line.  Thus, the existence of the WHIM in the hotter phase is not yet confirmed.   Several possible detections of WHIM emission have been reported. \\citet{2003A&A...397..445K} and \\citet{2004JKAS...37..375K} report the detection of redshifted O\\emissiontype{VII} and O\\emissiontype{VIII} emission lines as well as soft X-ray ($E<0.5~\\mathrm{keV}$) excesses in the spectra of 7 clusters of galaxies out of 21 measured systems. \\citet{finoguenov03:_xmm_newton_x_coma} also reported detection of O\\emissiontype{VII} and O\\emissiontype{VIII} emission lines in the outskirts of the Coma cluster. \\citet{fujimoto04:_probin_warm_hot_inter_medium} and Takei et al.\\ (2006) report not only emission lines in the spectrum of cluster outskirts but also absorption lines in the spectrum of a quasar behind the clusters, although the significance of the absorption measurement is not high ($<3\\sigma$). \\citet{fujimoto04:_probin_warm_hot_inter_medium} observed O\\emissiontype{VIII} features in the Virgo cluster, while Takei et al.\\ (2006) detected Ne\\emissiontype{IX} in the Coma cluster.  They estimated the path length of absorbing/emitting medium assuming that the absorber and the emitter are the same cloud, and concluded that the length scale of the medium exceeds the virial scale of the clusters. There are counterarguments against the discovery of O\\emissiontype{VII} or O\\emissiontype{VIII} emission lines associated with clusters. For example, \\citet{2005A&A...443...29B} and \\citet{2006ApJ...644..167B} concluded that some of the emission lines observed by \\citet{2003A&A...397..445K} are likely due to field-to-field variation of the soft X-ray background or Galactic emission. Such controversies occurred because the CCD instruments cannot clearly distinguish O lines emitted by  the gas associated with clusters of moderate redshift ($z\\lesssim 0.05$) from those of Galactic origin. These arguments suggest the importance of careful target selection and accurate measurement of background level. Note that the observations of redshifted emission lines also give important clues to solve another mystery ``cluster soft excess'', the excess emission in $E<0.3~\\mathrm{keV}$, reported since the 1990s in the extreme ultraviolet and the X-ray bands from {\\it EUVE} and {\\it ROSAT} \\citep[e.g.,][]{lieu96:_diffus_coma, 1996ApJ...458L...5L, 1998ApJ...498L..17M,2004ApJ...605..168B}. Thus, the study of X-ray emission in cluster outskirts brings us important information about the cluster system itself and its environment.  Since there are only a few reports of possible emission lines from the WHIM, further sensitive study of other systems is desirable.  In this paper, we report our observation of A2218 with the XIS (Koyama et al.\\ 2006) onboard {\\it Suzaku} (Mitsuda et al. 2006). A2218 is a well-known cluster because of its strong gravitational lensing arcs.  Its redshift is 0.1756 \\citep{struble99:_compil_redsh_veloc}, where one arcminute corresponds to 179 kpc with the cosmology we have adopted. The cluster is thought to be still in a dynamically young state --- X-ray analysis suggests that the cluster core is not in hydrostatic equilibrium due to an ongoing or recent merger \\citep{2005A&A...433..777P, 2002ApJ...567..188M} in the line-of-sight direction. The distribution of galaxy velocities based on the optical study of \\citet{1997ApJ...490...56G} indicates substructure in the line of sight, and \\citet{2005MNRAS.359..417S} studied the mass structure in detail and have concluded that A2218 is unrelaxed.  The dynamics of the cluster is considered to cause the discrepancy between the mass estimated from gravitational lensing and that determined from X-ray observations; gravitational lensing indicates $\\sim 3$ times higher mass than X-ray estimation for the central ($\\sim20''$) region. These observations indicate an elongated structure in the line-of-sight direction and suggests existence of a large-scale filament in this direction.  If this is the case, we would expect a high surface brightness of the WHIM emission. The fairly large redshift of A2218 would also help us observe a significant energy shift of emission lines with the XIS instrument. Thus, A2218 seems to be one of the most suitable targets to search for the WHIM emission from Suzaku.  We assume a Hubble constant of 70 km s$^{-1}$ Mpc$^{-1}$ or $h_{70}=1$ and $\\Omega_\\mathrm{m} = 0.3$, $\\Omega_\\mathrm{\\Lambda} = 0.7$. The solar metal abundances used in this paper are given by \\cite{anders89:_abund}. Unless otherwise stated, errors in the figures are at the 68~\\% confidence level and at the 90~\\% confidence level in the text and tables. ",
        "conclusions": "We have observed A2218 with the XIS instrument onboard {\\it Suzaku} to search for the redshifted O emission lines from the WHIM, which is possibly forming a large-scale filament around the cluster.  The object was selected because its redshift ($z=0.1756$) would allow the XIS to separate the WHIM lines from the Galactic ($z=0$) ones, and because an elongated structure in the line-of-sight direction was suggested from the previous studies of this cluster.  We detected no redshifted O lines, and set a constraint on their intensities.  In this section, we compare our constraints on the line intensity with those reported in other works and also discuss the future prospect of studying the WHIM emission with the {\\it Suzaku} XIS.  \\subsection{Comparison with other results}  \\citet{2003A&A...397..445K} and \\citet{finoguenov03:_xmm_newton_x_coma} reported positive detections of O lines around clusters of galaxies based on {\\it XMM-Newton} observations. \\citet{2003A&A...397..445K} detected significant O lines in three clusters, S\\'{e}rsic~159--03, MKW 3s and A2052, and possible, but not uniquely proven, O lines in the Coma and A1795 clusters. \\citet{finoguenov03:_xmm_newton_x_coma} detected the O lines in the outskirts of the Coma cluster, in particular, the Coma-11 field. Fig.~\\ref{fig:compareoxygen} compares the O line surface brightness $I$ for their observations and our results.  The left and right panels are for the O\\emissiontype{VII} and O\\emissiontype{VIII} cases, respectively.  The intensity of the O lines of Galactic emission measured by a microcalorimeter experiment of \\cite{2002ApJ...576..188M}, the one compiled by \\citet{2002A&A...389...93L} and the one estimated from our Offset-A observation are also shown.  The surface brightness quoted for \\citet{2003A&A...397..445K} was calculated using the temperature and emission measure in their Table 7 and the metal abundance from Tables 4 and 5 of \\citet{2004A&A...420..135T}, assuming the solar abundance ratio Fe to O \\citep[Fe/O number density ratio of 0.55;][]{anders89:_abund}. The surface brightness levels of the O\\emissiontype{VII} or O\\emissiontype{VIII} emission lines that were reported so far from the {\\it XMM-Newton} observations are similar to or higher than that of the Galactic emission.  In contrast, the upper limit of O\\emissiontype{VII} line intensity in the A2218 vicinity obtained in this work is about six times lower than the Galactic level.  The upper limit for O\\emissiontype{VIII} is also lower than the level reported as a positive detection in other works. The tight upper limit we obtained here demonstrates the good spectral capability of the XIS, in particular below 0.5 keV\\@.  \\begin{figure}[hbt] \\begin{center} \\FigureFile(0.47\\columnwidth,clip){figure7a.eps} \\hspace{0.03\\columnwidth} \\FigureFile(0.47\\columnwidth,clip){figure7b.eps} \\end{center} \\caption{Comparison of O\\emissiontype{VII} (left) and O\\emissiontype{VIII} (right) surface brightness. From left to right, those in the Coma-11 field \\citep{finoguenov03:_xmm_newton_x_coma}, S\\'{e}rsic~159--03, MKW 3s and A2052 \\citep{2003A&A...397..445K}, Galactic emission of \\cite{2002ApJ...576..188M}, \\citet{2002A&A...389...93L}, and our Offset-A observation, and the upper limits in A2218 outskirts (this work) are plotted. } \\label{fig:compareoxygen} \\end{figure}  \\subsection{Constraint on the WHIM density}  If we assume that the O\\emissiontype{VII} emission line is produced in a cloud with uniform density and temperature under collisional ionization equilibrium, the surface brightness $I$ is determined by the electron density $n_\\mathrm{e}$, the hydrogen density $n_\\mathrm{H}$, path length $L$, and metal abundance $Z$ of the cloud as \\begin{equation} I = C(T)~ (1+z)^{-3}~ n_\\mathrm{e}n_\\mathrm{H}~ ZL, \\end{equation} where $C(T)$ is a coefficient that depends on the temperature of the cloud.  Here, we will constrain the density of the cloud assuming the temperature to be $T\\sim 2\\times 10^{6}$~K, since the O\\emissiontype{VII} line is strongest at this temperature. Substituting $C(T=2\\times 10^{6}~\\mathrm{K})$ with the value calculated with the SPEX code \\citep{1996uxsa.coll..411K}, our constraint, $I < 1.1 \\times 10^{-7} $~$ \\mathrm{photons~cm^{-2}~s^{-1}~arcmin^{-2}}$ at $z=0.1756$, and electron-to-hydrogen number density ratio $n_\\mathrm{e}/n_\\mathrm{H}=1.2$ for ionized gas, gives the following condition \\begin{equation} n_\\mathrm{H} < 7.8\\times10^{-5}~\\mathrm{cm^{-3}}~ \\left(\\frac{Z}{0.1~Z_\\mathrm{\\odot}}\\right)^{-1/2}~ \\left(\\frac{L}{2~\\mathrm{Mpc}}\\right)^{-1/2}. \\end{equation} Note that $L=2$~Mpc corresponds to $11'$, the average diameter of the annular region where we extracted the spectra. The overdensity $\\delta\\equiv n_{\\rm H}/\\bar{n}_{\\rm H}$ of this cloud is  \\begin{equation} \\delta < 270~ \\left(\\frac{Z}{0.1~Z_\\mathrm{\\odot}}\\right)^{-1/2}~ \\left(\\frac{L}{2~\\mathrm{Mpc}}\\right)^{-1/2}, \\end{equation} where $\\bar{n}_{\\rm H} = X\\Omega_\\mathrm{b}\\rho_\\mathrm{crit} (1+z)^3 /m_\\mathrm{p} = 1.77\\times10^{-7}~(1+z)^3~\\mathrm{cm^{-3}}$ is the mean hydrogen density in the universe, in which $X=0.71$ is the hydrogen-to-baryon mass ratio, $\\Omega_\\mathrm{b}=0.0457$ is the baryon density of the universe, $\\rho_\\mathrm{crit}=9.21\\times10^{-30}~\\mathrm{g~cm^{-3}}$ is the critical density of the universe, and $m_\\mathrm{p}$ is the proton mass.  Even though this level of overdensity is much higher than the typical WHIM density ($\\delta \\sim 10$), it shows that Suzaku can certainly detect the high-density end of the WHIM distribution that is predicted to exist near clusters.  \\subsection{Prospect for the WHIM observation with the XIS}  Although we have set a tight constraint on the intensity of redshifted O lines, the inherent sensitivity of the XIS was presumably not achieved, since the redshifted O\\emissiontype{VIII} line fell almost on the Galactic O\\emissiontype{VII} line, while the redshifted O\\emissiontype{VII} line overlapped with the instrumental neutral O edge.  Therefore, better sensitivity is expected for clusters with a redshift such that the redshifted lines are free from the Galactic and instrumental features.  For example, a redshift of $z\\sim0.07$ or $z>0.2$ is suitable for this purpose.  The former case is more promising, because the surface brightness of high-$z$ clusters is lower and because the contamination of bright ICM emission in the center of the cluster makes the detection of warm-hot emission in the outskirts difficult for distant clusters.  Note that the XIS is capable of separating a redshift difference of $\\Delta z\\sim0.07$ at the O line energy.  Because of the contamination on the XIS OBF, the effective area at the redshifted O\\emissiontype{VII} line had dropped by $\\sim 25\\%$ at the time of the A2218 observation.  If there had been no contamination on the OBF, we could have obtained more photons by this fraction.  This contamination also caused a large systematic uncertainty corresponding to about 5\\% of the OVII line flux.  Since the thickness and chemical composition of the contaminant will be better understood as more data are accumulated, the uncertainty will be less influential in future observations.  Furthermore, we hope that the original XIS sensitivity in the soft X-ray energy range can be recovered in the future by hardware changes including warming the contaminated OBFs.  It should be noted that in most cases concerning the search for faint WHIM emission with CCD sensors, the variation of Galactic emission causes a larger uncertainty than the statistical one.  Offset observations, such as those carried out in our work, are always desirable to determine the background emission level reliably. Looking at a somewhat longer time frame, observations with microcalorimeters are promising, because their outstanding energy resolution ($\\Delta E \\lesssim 7$ eV) for extended sources can easily separate the WHIM emission from the Galactic spectral features. For example, two orders of magnitude higher sensitivity than the present work is expected with a small mission such as {\\it DIOS} \\citep[Diffuse Ionized Oxygen Surveyor;][]{2007_ohashi_dios_spie}, which is proposed to be launched in early 2010s \\citep{2004PASJ...56..939Y}."
    },
    "0608/astro-ph0608522_arXiv.txt": {
        "abstract": "We present wide-field H$\\alpha$ images of the Galactic supernova remnant G156.2+5.7 which reveal the presence of considerable faint H$\\alpha$ line emission coincident with the remnant's X-ray emission. We also present low-resolution optical spectra for a few representative emission regions.  The outermost H$\\alpha$ emission consists largely of long and thin (unresolved), smoothly curved filaments of Balmer-dominated emission presumably associated with the remnant's forward shock front. Patches of brighter H$\\alpha$ emission along the western, south-central, and northeastern regions appear to be radiative shocked ISM filaments like those commonly seen in supernova remnants, with relatively strong [O~I] $\\lambda\\lambda$6300,6364 and [S~II] $\\lambda\\lambda$6716,6731 line emissions. Measured [S~II] $\\lambda\\lambda$6716,6731/H$\\alpha$ ratios range from 0.84 to 1.57.  Comparison of the observed H$\\alpha$ emission with the {\\it ROSAT} PSPC X-ray image of G156.2+5.7 shows that the thin Balmer-dominated filaments lie along the outermost edge of the remnant's detected X-ray emission. Brighter radiative emission features are not coincident with the remnant's brightest X-ray or radio regions. Areas of sharply weaker X-ray flux seen in the {\\it ROSAT} image of G156.2+5.7 appear spatially coincident with dense interstellar clouds visible on optical and {\\it IRAS} 60 and 100 ~$\\mu$m emission images, as well as maps of increased optical extinction. This suggests significant X-ray absorption in these regions due to foreground interstellar dust, especially along the western and southern limbs. The close projected proximity and alignment of the remnant's brighter, radiative filaments with several of these interstellar clouds and dust lanes hint at a possible physically interaction between the G156.2+5.7 remnant and these interstellar clouds and may indicate a smaller distance to the remnant than previously estimated. ",
        "introduction": "A majority of the $\\sim$230 currently confirmed Galactic supernova remnants (SNRs) listed in the latest SNR catalogue \\citep{Green04} were first identified through radio surveys. However, with the advent of sensitive X-ray satellites, about half a dozen remnants have now been discovered via their X-ray emission (cf. \\citealt{Busser96} and references therein).  The first Galactic SNR discovered through its X-ray emission was reported by \\citet*{Pfeff91}. They found a previously unknown Galactic SNR based on the {\\it ROSAT} detection of a relatively bright and nearly circular 108' diameter emission structure centred on $\\alpha$(J2000) = 4$^{\\rm h}$ 59$^{\\rm m}$ $7^{s}$, $\\delta$(J2000) = 51$^{\\rm o}$ $46'$ $35''$. With an estimated 0.1 -- 2.4~keV X-ray flux  of $1.9 \\times 10^{-10}$ erg cm$^{-2}$ s$^{-1}$, this new remnant ranked among the ten brightest X-ray Galactic SNRs known at that time.  \\citet{Pfeff91} concluded that the remnant was located in a region of very low interstellar density (0.01 atoms cm$^{-3}$), which might help account, in part, for the remnant's lack of strong radio emission.  The new remnant, designated G156.2+5.7 (also RXJ04591+5147), was subsequently detected and confirmed in the radio by \\citet*{Reich92}.  In contrast to the object's bright X-ray flux, its 1 GHz radio surface brightness of $5.8 \\times 10^{-23}$ W m$^{-2}$ Hz$^{-1}$ sr$^{-1}$ ranks it as perhaps the faintest Galactic SNR currently known \\citep{Green04}, being a factor of two in 1 GHz surface brightness below than of the exceptionally radio faint but optically [O~III] bright remnant, G65.3+5.7 \\citep{Gull77,Reich79}. Follow-up searches for a central radio pulsar or central compact object have proven unsuccessful \\citep{Lorimer98,Kaplan06}.  \\begin{figure*} \\includegraphics[width=170mm]{fig1.eps} \\caption{ H$\\alpha$\\ mosaic image of the G156.2+5.7 supernova remnant.  Two distinct emission morphologies are seen: a faint rim of thin filamentary emission which runs most of the way around the remnant; and patches of much brighter emission with a clumpy, complicated structure. Note the white dust absorption patch along the remnant's western limb. To reduce the appearance of background noise the data have been smoothed slightly using a Gaussian filter. \\label{ha}} \\end{figure*}  Since shortly after its discovery, G156.2+5.7 has been studied only in the X-rays. Using the {\\it ASCA} satellite, \\citet{Yama99,Yama00} found both hard and soft X-ray components, an estimated Sedov age of $\\sim 1.5 \\times 10^{4}$ yr, and an ambient gas density of $\\sim$0.2 cm$^{-3}$. More recently, \\citet{Pann04} used both {\\it ASCA} and {\\it RXTE} to study the the northwestern rim of G156.2+5.7 and investigate the presence of a non-thermal X-ray component and cosmic-ray acceleration.  \\citet{Pfeff91} reported finding no associated optical emission on the Palomar Sky Survey plates and there are no discernible emission features in this region visible on the relatively deep, narrow passband filter images of \\citet*{Parker79} in their emission-line survey of the Milky Way.  However, here we present wide field H$\\alpha$ images of the remnant which reveal a remarkably extensive and complex optical emission structure covering a large fraction of the remnant's X-ray shell. ",
        "conclusions": "Our mosaic H$\\alpha$ image of G156 (Fig.\\ 1) shows a rather surprising amount of optical shock emission for a remnant previously reported to exhibit no detectable optical emission.  G156's detected optical emission structure now ranks it among the most extensive seen among galactic remnants, albeit at the faint end of surface brightness.  The lack of visibility of G156 on the Palomar Sky Survey is not particularly surprising given the broad passband of its images. Nor is it surprising that the SNR emission was not detected by \\citet{Stecklum04} who imaged a region around the Bok globule CB 26 close to the northwest filament complex in both H$\\alpha$ and [S~II], but were limited by a small ($2'$) field of view.  But the lack of detection by the narrow passband filter galactic plane survey of \\citet{Parker79} is puzzling. That survey's l=155.5 $\\degr$, b=+5 $\\degr$ plates are well centred on G156 and yet shows no discernible H$\\alpha$, [S~II] $\\lambda\\lambda$6716,6731, or [O~III] $\\lambda$5007 emission.  G156's detected optical emission consists of a considerable amount of both radiative and non-radiative filaments, something that is not commonly seen in SNRs.  Only a handful of galactic remnants show such extensive non-radiative, Balmer-dominated emission filaments alongside considerable radiative filamentary nebulae. The few comparable objects include the Cygnus Loop, with its bright radiative west and east nebulae (NGC 6960 \\& 6995) and extensive Balmer-dominated filaments outlining much of the remnant's limb \\citep{Levenson98}, and RCW 86 with its lone but bright radiative emission SW limb region and non-radiative filaments outlining most of the rest of the remnant \\citep{Smith97}.  The mixture of both radiative and non-radiative optical emission suggests that the remnant is situated in a highly non-homogenous ISM, with the radiative emission associated with relatively dense cloud--SNR interactions while the thin non-radiative filaments mark regions of much lower interstellar density. However, \\citet{Pfeff91} reported no spectral variations in four equally spaced sectors of the remnant, concluding that there was no significant variation in interstellar absorption across G156.  They did find a range of N(H) column density of $9 \\times 10^{20}$ to $ 2 \\times 10^{22}$ cm$^{-2}$ depending on the choice of model assumed, and chose the lower value in their final analysis. \\citet{Yama00} estimated N(H) = $3-4 \\times 10^{21}$ cm$^{-2}$ using NEI plus thermal or power law models.  Neither study discussed the possibility of N(H) variations, and yet the presence of these large interstellar clouds along the line of sight clearly indicates that the foreground column density varies significantly across the face of the remnant. This complicates the assessment of the X-ray results, as the foreground density is position dependent on scales smaller than the effective apertures used in these studies.  A re-analysis of the {\\it ROSAT} data is certainly warranted, using emission primarily from the relatively unobscured east side of the remnant.  The projected proximity and apparent alignment of the remnant's brighter radiative filaments with several foreground interstellar clouds and dust lanes hint at a possible physically interaction between the G156.2+5.7 remnant and these interstellar clouds. The estimated distance to the T-Tauri star V347 Aur associated with one of the southwestern foreground clouds is 250 -- 300 pc \\citep{Cohen78}, while the estimated distance to the Bok globule CB $26$ near the remnant's bright northwest filament complex ranges from 140 to 300 pc \\citep{LH97,LS01}.  If G156 is indeed interacting with part of the Taurus--Auriga cloud complex in this direction, then the G156 remnant may be only $\\approx$ 0.3 kpc distant, significantly closer than the previous X-ray derived estimates of 1.3 -- 3 kpc. It would also be much younger than the 15,000 -- 26,000 yr ages earlier analyses have suggested.  In fact, assuming G156 is in the Sedov-Taylor (adiabatic) expansion phase of its evolution, lies at a distance of just 0.3 kpc, is expanding in an intercloud density of 0.2 cm$^{-3}$ \\citep{Yama99,Yama00}, and was generated by a SN with energy around the canonical value of of $1 \\times 10^{51}$ erg, then it has a radius of $\\approx$ 5 pc, an age of $\\approx$ 400 yr, and a shock velocity around 5000 km s$^{-1}$. Such a young age seems highly unlikely, as it should then have been seen as a historical supernova.  This extreme age problem is reduced somewhat if either the distance to the clouds is somewhat farther, or the ambient ISM density is higher.  For example, at a distance of 0.6 kpc, a Sedov analysis of G156 would give an age of $\\approx$ 2200 yr, and a shock velocity around 1500 km s$^{-1}$.  Such an age would make G156 more akin to the SN~1006 SNR than the Cygnus Loop.  Its nearly filled X-ray morphology is certainly unlike that seen in moderately old SNRs and an age of only a few thousand years would be more consistent with the observed presence of extensive non-radiative shock filaments which are most often seen associated with high-velocity shocks in young remnants. A relatively youthful remnant might also help explain the apparent enhanced abundances of Si and S seen in the X-ray spectra of the centre of G156 \\citep{Yama99}.  On the other hand, the close apparent placement of the foreground clouds and the radiative emission could be simply a coincidence.  Obviously, further investigations into the distance and general properties of this SNR are needed. For example, if its shock velocity is $\\sim$ 1000 -- 2000 km s$^{-1}$, then high resolution spectra of its Balmer dominated filaments should show evidence for a broad H$\\alpha$ emission component like that seen in the filaments in Tycho and SN~1006.  Moreover, these Balmer dominated filaments should exhibit large proper motions of order 1 arsec yr$^{-1}$. Even null results would provide useful lower limits to the distance to G156, and hence lower limits on the age of the SNR.   Whatever the outcome of future studies, G156 is clearly an interesting remnant, with a now recognized extensive and diverse optical emission structure."
    },
    "0608/astro-ph0608378_arXiv.txt": {
        "abstract": "The All-wavelength Extended Groth Strip International Survey (AEGIS) team presents broad-band spectral energy distributions (SEDs), from X-ray to radio wavelengths, for 71~galaxies spanning the redshift range 0.55--1.16 ($\\langle${$z$}$\\rangle${$\\sim$}0.7).  Galaxies with secure redshifts are selected from a small (22\\,arcmin$^2$) sub-section of the Keck\\footnote{Data presented herein were obtained at the W.\\ M.\\ Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California, and NASA. The Observatory was made possible by the generous financial support of the W.\\ M.\\ Keck Foundation.}/DEIMOS galaxy redshift survey in the Extended Groth Strip field that has also been targeted for deep panchromatic imaging by {\\it Chandra\\/}\\footnote{NASA's {\\it Chandra\\/} X-Ray Observatory was launched in July 1999.  The {\\it Chandra\\/} Data Archive is part of the {\\it Chandra\\/} X-Ray Center which is operated for NASA by the Smithsonian Astrophysical Observatory.} (X-ray: 0.5--10\\,keV), GALEX\\footnote{The Galaxy Evolution Explorer is a NASA Small Explorer, launched in April~2003.  We gratefully acknowledge NASA's support for construction, operation, and science analysis of the GALEX mission, developed in cooperation with the Centre National d'Etudes Spatiales of France and the Korean Ministry of Science and Technology.} (ultraviolet: 120--250\\,nm), Canada-France-Hawaii Telescope\\footnote{Based on observations obtained with the Canada-France-Hawaii Telescope which is operated by the National Research Council of Canada, the Institut National des Science de l'Univers of the Centre National de la Recherche Scientifique of France, and the University of Hawaii.} (optical: 360--900\\,nm), {\\it Hubble Space Telescope\\/}\\footnote{Based on observations with the NASA/ESA {\\it Hubble Space Telescope\\/}, obtained at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS~5-26555.} (optical/near infrared: 440--1600\\,nm), Palomar Observatory\\footnote{Based on observations obtained at the Hale Telescope, Palomar Observatory, as part of a collaborative agreement between the California Institute of Technology, its divisions Caltech Optical Observatories and the Jet Propulsion Laboratory (operated for NASA), and Cornell University.} (near infrared: 1200--2200\\,nm), {\\it Spitzer\\/}\\footnote{Based on observations made with the {\\it Spitzer\\/} Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA.  Support for this work was provided by NASA through contract numbers 1256790, 960785, and 1255094 issued by JPL/Caltech.} (mid/far infrared: 3.6--70~$\\mu$m), and the Very Large Array\\footnote{The Very Large Array of the National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} (radio: 6--20\\,cm). The absolute magnitude of the typical galaxy in our sample is $M_B$=$-$19.82. The ultraviolet to mid-infrared portion of their spectral energy distributions (SEDs) are found to be bracketed by two stellar-only model SEDs: (1)~an early burst followed by passive evolution and (2)~a constant star-formation rate since early times; this suggests that few of these galaxies are undergoing major starbursts.  Approximately half the galaxies show a mid- to far-infrared excess relative to the model SEDs, consistent with thermal emission from interstellar dust.  Two objects have power-law SEDs, indicating that they are dominated by active galactic nuclei; both are detected in X-rays.  The galaxies are grouped by rest-frame color, quantitative optical morphology, and [O\\,{\\sc ii}] emission line strength (possible indicator of star formation).  On average, the panchromatic SEDs of the galaxies, from the ultraviolet to the infrared, follow expected trends: redder SEDs are associated with red $U${$-$}$B$, early-type morphology, and low [O\\,{\\sc ii}] emission, and vice versa for blue SEDs.\\footnote{These data are available at {\\tt http://aegis.ucolick.org/products} .} ",
        "introduction": "\\label{sec:intro}  %  The spectral energy distribution (SED) of a galaxy contains key clues about the dominant physical processes responsible for energy generation.  Space- and ground-based telescopes of the new millennium provide the sensitivity and angular resolution to make measurements across the breadth of the electromagnetic spectrum from X-ray to radio wavelengths.  Panchromatic SEDs can tell us about the stellar populations within the galaxy from their visible light [and ultraviolet (UV) and near-infrared (IR)] properties \\citep[e.g.,][and references therein---hereafter BC03]{bc03}, thermal re-radiation of starlight in the far IR by interstellar dust \\citep{dal01}, and X-ray \\citep[e.g.,][]{mus80} and radio emission from active galactic nuclei (AGNs).  \\citet{dal01} have studied panchromatic SEDs of nearby galaxies and have developed a semi-empirical model for the IR SED of normal star-forming galaxies.  Distant galaxy SEDs are important for look-back studies of galaxy formation and evolution.  The first such survey was GOODS \\citep{gia04}.  Our project, the All-wavelength Extended Groth Strip International Survey \\citep[AEGIS;][hereafter Paper\\,I]{dav06}, extends the work of GOODS over a larger survey area [in a part of the sky known as the Extended Groth Strip (EGS)] thereby driving down cosmic variance.  The cornerstone of AEGIS is the Keck/DEIMOS DEEP2 galaxy redshift survey \\citep{dav03}.  Precise redshift measurements from this survey allow us to determine the environment of each galaxy \\citep{coo05}, a key factor in determining galaxy evolution. Moreover, these DEIMOS spectra allow us to measure linewidths and estimate dynamical masses for the galaxies.  This paper presents SEDs of 71~optically-selected galaxies in AEGIS and is intended to showcase the panchromatic data set.  Given the relatively coarse angular resolution of even the state-of-the-art UV and mid-/far-IR images in our survey, we limit our study to the {\\it global\\/} emission from these distant galaxies.  The goal of this paper is to provide clean SEDs for a handful of objects. Throughout this letter we use a cosmology of ($H_0$,\\,$\\Omega_{\\rm m}$,\\,$\\Omega_\\Lambda$)=(70,\\,0.3,\\,0.7).  Magnitudes are in the AB system. ",
        "conclusions": ""
    },
    "0608/astro-ph0608464_arXiv.txt": {
        "abstract": "{% } {% This paper presents a new homogeneous catalogue of blue straggler stars (BSS) in Galactic open clusters.} {% Photometric data for 216 clusters were collected from the literature and 2782 BSS candidates were extracted from 76 of them.} {% We found that the anticorrelation of BSS frequency vs. total magnitude identified in similar studies conducted on Galactic globular clusters extends to the open cluster regime: clusters with smaller total magnitude tend to have higher BSS frequencies. Moreover, a clear correlation between the BSS frequency (obtained normalising the total number of BSS either to the total cluster mass or, for the older clusters, to the total number of clump stars) and the age of the clusters was found. A simple model is developed here to try to explain this last and new result. The model allows us to ascertain the important effect played by mass loss in the evolution of open clusters.} {% } ",
        "introduction": "Blue straggler stars (BSS) occupy a region above the turn-off (TO) point in the colour-magnitude diagram (CMD) of star clusters, where no stars are expected on the basis of standard stellar evolution, if we assume that all cluster stars are coeval.  A number of exotic explanations have been invoked to interpret the nature of these stars. Briefly, it has been suggested that these stars are the outcome of stellar collisions \\citep{Ben87, Lom96} or of mass exchange in close binary systems \\citep{McC64, Egge89, Mat90}.  Other possible scenarios include second-generation stars \\citep{Egge88}, accretion of gas from the interstellar surrounding medium \\citep{Wil64}, and capture of field stars by a star cluster.  Interestingly, BSS are ubiquitous, as they have been found in globular clusters (GCs), dwarf galaxies, open clusters (OCs), and in the Galactic field. Therefore, important constraints on their nature and possible formation mechanisms can be derived from a comparative study of the BSS properties in different environments. In the past two decades, a huge amount of new, high-quality CMDs from CCD photometry from ground-based and space facilities have become available. Recently, \\citet[hereafter PDK04]{Pio04} published a new photometrically homogeneous catalogue of BSS extracted from the CMDs from HST observations of more than 1/3 of the known Galactic GCs \\citep{Pio02}. The present paper aims to complement the PDK04 catalogue by extending the study of BSS properties to the different environment that is typical of OCs.  The catalogue presented in this paper can be considered as an important extension of the effort made by \\citet[hereafter AL95]{Ahu95} to compile a catalogue of 969 BSS in 390 OCs, based on the photometry obtained with photographic plates and photoelectric detectors in the 60s, 70s, and 80s. The main conclusions drawn from their catalogue are that (1) BSS are present in clusters of all ages, (2) the fraction of BSS increases with the richness and age of a cluster, and (3) BSS show a remarkable degree of central concentration.  Moreover, the number of BSS over the number of main sequence stars seems constant up to an age of 400 Myr.  The main reasons that led us to perform a new search of BSS in OCs are the following: (1) the AL95 catalogue was based on a highly inhomogeneous set of data, (2) a large number of new, high-quality photometric studies of OC stars have appeared since AL95, and (3) the BSS selection criteria needed in our opinion to be revised, especially for young OCs. Most important, in order to estimate the relative number of BSS in different clusters, the total number of BSS must be normalised to some reference stellar population in the same cluster. The choice made by AL95 to normalise the BSS counts to the main sequence stars creates some problems, especially in young clusters, where it is difficult to decide whether a star is a main sequence, a blue straggler, or a field star.  The layout of this paper is as follows.  In Sect.~2 we present the new database.  In Sect.~3 we describe a few interesting properties highlighted by the new catalogue. In Sect.~4 we develop a simple model in order to interpret the observed BSS frequency in OCs.  Section~5 presents our conclusions. ",
        "conclusions": ""
    },
    "0608/astro-ph0608193_arXiv.txt": {
        "abstract": "We measure the volume luminosity density and surface luminosity density generated by the Galactic disc, using accurate data on the local luminosity function and the disc's vertical structure. From the well measured volume mass density and surface mass density, we derive local volume and surface mass-to-light ratios for the Galactic disc, in the bands $B$, $V$ and $I$.  We obtain mass-to-light ratios for the local column of stellar matter of $(M/L)_B = 1.4 \\pm 0.2$, $(M/L)_V = 1.5 \\pm 0.2$ and $(M/L)_I = 1.2 \\pm 0.2$. The dominant contributors to the surface luminosity in these bands are main sequence turn-off stars and giants. Our results on the colours and mass-to-light ratios for the ``Solar cylinder'' well agree with population synthesis predictions using Initial Mass Functions typical of the Solar Neighbourhood. Finally we infer the global luminosity of the Milky Way, which appears to be under-luminous by about 1-$\\sigma$ with respect to the main locus of the Tully--Fisher relation, as observed for external galaxies. ",
        "introduction": "The mass--to--light ($M/L$) ratio\\footnote{Unless otherwise explicitly stated, this paper always refers to the $M/L$ of stellar matter} is an important constraint for studies of stellar populations and for chemo--photometric models of the Milky Way; these often serve as a calibration point for the modeling of disc galaxies in general (e.g.\\ Boissier \\& Prantzos 1999).  The surface luminosity and mean colours of the Solar column are also useful comparison points in extragalactic studies, and for placing the Milky Way on the Tully-Fisher relation.  Studies of the Galactic disc over the last two decades have resulted in good determinations of its local mass density $\\rho(0)$, and the surface mass density, $\\Sigma_0$ (e.g.\\ Holmberg \\& Flynn 2000, 2004 and references therein). In this paper we study the related issue of the luminosity generated by the local disc, both in the local volume (i.e.\\ the volume luminosity density) and integrated perpendicularly to the disc in a column (i.e.\\ the surface luminosity density). These quantities allow us to measure the mass-to-light ratio, $(M/L)$, for the local Galactic disc.  Estimating the luminosity surface density (i.e.\\ surface brightness) of the local Galactic disc requires good knowledge of its vertical structure.  In this respect, Galactic models have much improved since the 1980s, as a particular result of the Hipparcos satellite and star count programs made with the Hubble Space Telescope.  These data allow us to overcome the main deficiency in earlier studies of the local disc surface brightness and disc $M/L$; due to the lack of accurate distances to individual stars, the column luminosity (and mass) density had to be recovered via assumptions about the scale-length $h_R$ and scale-height $h_z$ of the stellar disc, and the results were degenerate with respect to the assumed $h_z/h_R$ ratio. Distances to stars from Hipparcos now allow us to determine directly the distribution of stellar scale-heights and measure column densities, merely by ``counting'' stars and light at the Galactic poles; $h_R$ is no longer required in the modeling and the degeneracy is broken.  The infrared structure of the Milky Way has been extensively studied in the 1990s, mostly taking advantage of the COBE/DIRBE experiment, and is presently well understood (Kent et~al.\\ 1991; Dwek et~al.\\ 1995; Binney et~al.\\ 1997; Freudenreich 1998; Bissantz \\& Gerhard 2002). On the other hand, in the optical, literature estimates of the disc's luminosity and $M/L$ ratio can mostly be traced back to work done prior to the launch of both the Hipparcos satellite and Space Telescope (see Table~\\ref{tab:literature} for a list of the main references). The time is thus ripe for a redetermination of the disc's mass-to-light ratio.  \\begin{table*} \\caption{Previous determinations of the surface brightness and colours of the local Galactic disc} \\begin{tabular}{c c c c c|c c c| p{4truecm}} \\hline \\multicolumn{2}{c}{Solar cylinder} & \\multicolumn{3}{c}{Disc} & \\multicolumn{3}{c}{Galaxy} \\\\ $\\mu_B (R_{\\odot})$ & $\\mu_V (R_{\\odot})$ & $M_{B,d}$ & $M_{V,d}$ & $(B-V)_{0,d}$ & $M_{B,g}$ & $M_{V,g}$ & $(B-V)_{0,g}$ & references \\\\ \\hline 24.15$\\pm$0.07 & & --19.61$\\pm$0.06 & & 0.40 & --20.08$\\pm$0.04 & & 0.53 $\\pm$0.02 & de Vaucouleurs \\& Pence (1978)\\\\ & $\\sim$23 & --19.9 & --20.4 & 0.45 & --20.5 & --20.1 & 0.45 & Bahcall \\& Soneira (1980), \\\\ & & & & & & & & Bahcall (1984) \\\\ 23.3$\\pm$0.3  & 22.7$\\pm$0.2 & & & 0.62$\\pm$0.04 & & & & Ishida \\& Mikami (1982) \\\\ 23.8$\\pm$0.1 & & --20.2$\\pm$0.2 & & 0.84$\\pm$0.15 & --20.3$\\pm$0.2 & & 0.83 $\\pm$0.15 & van der Kruit (1986) \\\\ \\hline \\end{tabular} \\label{tab:literature} \\end{table*}  In sections 2 and 3 we determine the disc luminosity and mass-to-light ratio, both locally and in a column at the Sun, for the optical bands $B$, $V$ and $I$.  In section~4, we compare our results to theoretical predictions from population synthesis modeling. In section 5, we compute the total disc luminosity, for a range of plausible scale-lengths, and make some comparisons between the Milky Way and external galaxies.  In section 6 we summarise and discuss our results. The present study focuses on optical bands, and work is underway to extend these measurements to near infra-red bands. ",
        "conclusions": "We have presented new estimates of the $B, V, I$ luminosity density, colours and mass-to-light ratios of the ``Solar volume'' and ``Solar cylinder'' , i.e.\\ the local Milky Way disc, based on Hipparcos/Tycho data and the use of a calibrated model for the vertical structure of the disc to correct from volume to column quantities (the Tuorla study). Excellent agreement is found with the parallel Heidelberg study determining the local luminosity density on completely empirical grounds, based on the complete CNS4 catalogue of nearby stars, at least over the (more limited) magnitude range probed by this latter study.  Our determination of the $B, V, I$ colours and stellar mass-to-light ratios for the Solar cylinder is in excellent agreement with theoretical expectations, from stellar population synthesis, for a Kroupa/Chabrier IMF representative of the Solar Neighbourhood. The uncertainty in the dynamical estimate of the local surface mass density indicate the Salpeter IMF as the extreme upper limit allowed for the local stellar $M_*/L$; but the Kroupa/Chabrier scaling remains most plausible.  Surface, or column, luminosities and colours provide important constraints for chemo--photometric models of the Solar Neighbourhood and the Milky Way, and allow to compare the Milky Way to external spirals. We have reconstructed from the local disc brightness the total $I$ band luminosity of the Milky Way, which appears to be low with respect to the TF relation; although the discrepancy is not dramatic (about 1~$\\sigma$), we discussed the implications and possible ways out.  On the other hand, were the Milky Way simply a 1~$\\sigma$ outlier from the TF relation (nothing to upset the Copernican principle), it is still worth to underline such offset as our Galaxy is usually taken as the paradigm for disc galaxies in general.  We are presently refining our $I$--band luminosity estimate with the aid of DENIS and other star counts, and defining better our errorbars; we are also extending our surface brightness determination to other photometric bands, especially in the infrared utilizing DENIS and 2MASS."
    },
    "0608/astro-ph0608470_arXiv.txt": {
        "abstract": "We propose a new model within the ``Quark-nova'' scenario to interpret the recent observations of early afterglows of long Gamma-Ray Bursts (GRB) with the Swift satellite. This is a three-stage model within the context of a core-collapse supernova. Stage 1 is an accreting (proto-) neutron star leading to a possible delay between the core collapse and the GRB. Stage 2 is an accreting quark-star, generating the prompt GRB. Stage 3, which occurs only if the quark-star collapses to form a black-hole, consists of an accreting black-hole. The jet launched in this accretion process interacts with the ejecta from stage 2, and could generate the flaring activity frequently seen in X-ray afterglows. This model may be able to account for both the energies and the timescales of GRBs, in addition to the newly discovered early X-ray afterglow features. ",
        "introduction": "With the launch of the Swift satellite \\citep{gehrels04}, observations of early X-ray afterglows from gamma ray bursts (GRBs) have become possible (sometimes as early as 80 seconds after the GRB trigger). This led to some surprising observations, most notably, the existence of one or more (sometimes giant) flares \\citep[first discussed by][]{burrows05} in the early X-ray afterglow (after a few hundred seconds to several thousand seconds) whose rapid variability has been interpreted as the inner engine being active much longer than the duration of the GRB itself \\citep[see for instance][]{zhang06}. Figure~\\ref{genericfigure} shows a canonical X-ray afterglow light curve. The early X-ray afterglow often has a very steep power-law decay, lasting up to about a thousand seconds, followed by a flattening of the light curve, which lasts for $10^4-10^5$ seconds. This is often overlayed with flares and bumps in the light curve. Thereafter, a ``classical'' decaying afterglow is seen, as it was known from the pre-Swift era. It should be noted that not every bursts exhibits all of these features. The flare(s) and the flattening of the light curve are likely due to extended inner engine activity, and, this extended inner engine activity is what we focus on in this paper.  The inner engine for the GRB in our model is an accreting quark-star formed shortly (hours) after the core collapse in a massive star. This process operates in three steps: i) the quark nova \\citep{ouyed02}, where the core is converted to quark matter resulting in  mass ejection, ii) fall back material from the supernova together with some of the matter ejected during the quark nova can form an accretion disk around the quark-star. Accretion onto the quark-star launches a jet that will overtake the shell ejected by the quark nova and is a possible location for the prompt X-ray  emission. Internal shocks within the jet produce the GRB, iii) if enough matter is accreted onto the quark-star, it will collapse and form a black-hole. Continued accretion onto the black-hole can lead to an ultra-relativistic jet. Interactions between this jet and the jet from the quark-star create shocks which lead to the flaring activity frequently seen in X-ray afterglows.  In this paper we assume that the gamma-ray emission is produced by internal shocks. The afterglow is produced when the merged shell creating the gamma-ray emission interacts with the external medium in an external shock.  It should be emphasized that quark stars are hypothetical objects, that have not been observed. They have been discussed theoretically for more than thirty years \\citep[see for instance][]{itoh70}. Two requirements are needed for quark stars to form. First the strange quark matter (SQM) hypothesis must hold true. This hypothesis states that SQM has a lower energy state than $^{56}Fe$, thus being the ground state of hadronic matter. Quark stars are usually thought to form from neutron stars whose central density increases above a critical density. Only the more massive neutron stars are thought to be able to reach this density \\citep{staff06}. Second, the SQM EOS must support a rather massive quark star (the mass will probably not change much in converting from a neutron star to a quark star). Optionally, a quark star can form directly in the collapse of the iron core in massive stars, in which case the second requirement might be relaxed.  In \\S~\\ref{modelsection} we present our model for an inner engine for GRBs in three stages, followed by a discussion of the energies and timescales of the inner engine in \\S~\\ref{ietimeenergiessection}. Then in \\S~\\ref{bhqsintsection} we discuss what happens when the jet from the black-hole and the quark-star interacts followed by a discussion of the formation of the different features in the X-ray afterglow in \\S~\\ref{lcfeaturesection}. We devote \\S~\\ref{flaresbumpsection} to correlations of features seen in the early afterglow light curve, and then in \\S~\\ref{casestudysection} we apply our model to different observed GRBs.  Then in \\S~\\ref{discussionsection} we discuss our model and finally we summarize our model in \\S~\\ref{summarysection}. ",
        "conclusions": "\\label{summarysection} To summarize, the properties of the three stages of the inner engine in our model are: \\begin{itemize} \\item The process is initiated by the collapse of the iron core of a massive, rapidly rotating star. \\item The collapsed core will either leave a neutron star or a quark-star behind. \\item If a neutron star is left behind, it can collapse to a quark-star at a later stage, creating a delay between the supernova and the GRB. \\item Fallback material from the supernova and the quark nova forms a disk around the quark-star. Accretion onto the quark-star generates the GRB by powering an ultrarelativistic jet. Internal shocks in this jet create the GRB. \\item Accretion continues, but at some point it cannot heat up the star sufficiently. This halts the emission of shells, ending the GRB. \\item When the quark-star has accreted too much material, it collapses into a black-hole. Further accretion onto the black-hole launches an ultrarelativistic jet. \\end{itemize} The emission features in our model can be summarized as follows: \\begin{itemize} \\item Early, steep decay in the X-ray afterglow is due to the curvature effect (this is not specific to our model, but rather a generic feature in many models). \\item Flares are created by interaction between the jet from the accretion onto a black-hole and slower parts from the jet from the quark-star. \\item Re-energization of the external shock (seen as flattening of the X-ray afterglow light curve) is due only to the jet from the quark-star. Slower parts of this jet re-energizes the external shock. \\item When the jet from the black-hole collides with the external shock from the GRB, a bump is seen in the afterglow light curve. \\end{itemize}  In conclusion, we have presented a three stage (effectively two stage) model for the inner engine for GRBs involving a neutron star phase, followed by a quark star phase then by a black hole phase. This model seems to account for the observed prompt gamma ray emission, as well as the features of the early X-ray afterglow and as such warrants further study."
    },
    "0608/astro-ph0608646_arXiv.txt": {
        "abstract": "We report \\h2o masers associated with the massive-star forming region \\g192 observed with the new Japan VLBI network at three epochs spanned for two months, which have revealed the three-dimensional kinematical structure of the whole \\h2o  maser region in \\g192, containing two young stellar objects separated by $\\sim$1200~AU. The maser spatio-kinematical structure has well persisted since previous observations, in which the masers are expected to be associated with a highly-collimated bipolar jet and an infalling-rotating disk in the northern and southern clusters of \\h2o maser features, respectively. We estimated a jet expansion speed of $\\sim$100\\kms\\ and re-estimated a dynamical age of the whole jet to be 5.6$\\times$10$^{4}$~yrs. We have investigated the spatial distribution of Doppler velocities during the previous and present observations and relative proper motions of \\h2o\\ maser features in the southern cluster, and a relative bulk motion between the two maser clusters. They are well explained by a model of an infalling--rotating disk with a radius of $\\sim$1000~AU and a central stellar mass of 5--10~$M_{\\odot}$, rather than by a model of a bipolar jet perpendicular to the observed CO outflow. Based on the derived \\h2o maser spatio-kinematical parameters, we discuss the formation mechanism of the massive young stellar objects and the outflow development in \\g192. ",
        "introduction": "\\label{sec:introduction} Formation mechanism of massive stars is one of the most important issues in astrophysics for elucidating their whole evolution as well as their effects on dynamical and chemical composition evolution of galaxies and the universe. It is difficult, however, to observationally find the true mechanism because all of massive young stellar objects (MYSOs) are associated with star clusters that are deeply embedded in molecular clouds located at large distances from the Sun ($d\\gtrsim$500~pc). Therefore, in whichever scenario of massive star formation, observations with very high angular resolution (better than 10~milliarcseconds, mas) is essential to resolve a massive-star forming region into individual MYSOs. Because a massive star also evolves extremely rapidly, it is difficult to find MYSOs at the earliest stage of the star formation process. An MYSO with $M\\gtrsim 8\\:M_{\\odot}$ has a Kelvin-Hermholtz time scale that is shorter than a free-fall time scale, therefore it reaches a zero-age main sequence star while mass accretion still occurs. There exists a basic question how such an MYSO builds its mass against mass ejection by strong stellar wind and expansion of an H{\\rm II} region.  There are two major possible scenarios of massive star formation; the first in term of a massive and thick accretion gas disk/torus and the second in term of lower-mass stars merged into a higher-mass star. In the first scenario, a parent molecular cloud core is still collapsing and providing material into an MYSO with a very high mass-accretion rate (e.g. $\\dot{M}\\gtrsim 10^{-3}\\: M_{\\odot}$~yr$^{-3}$) by way of a thick gas disk even if the central part of the core is photoionized (e.g., \\cite{ket03}). The presence of a disk near a massive star can significantly reduce the effect of radiation on the infalling material by reducing the exposure of matter to the radiation field and by allowing photons to escape along the stellar polar axis (e.g. \\cite{nag89,kru05}). A direct evidence for this scenario should be to directly detect mass accretion onto the surface or in very vicinity of an MYSO within $\\sim$1~AU. Previous observations have reported the existence of such massive disks/tori (e.g. \\cite{sak05,jia05,ket06}), but their physical scales directly elucidated by the observations are still much larger ($>$100~AU). Nevertheless these observations suggested that a single star with a mass at least equal to or lower than 7~$M_{\\odot}$ can be created in this scenario \\citep{jia05}. The formation mechanism of a higher-mass star is still obscure. \\citet{chi04} reported an MYSO in M17 with a stellar mass of 20~$M_{\\odot}$ and an accompanying massive gas/dust disk, but the estimated mass is reevaluated by \\citet{sak05} to be only 3--8~$M_{\\odot}$. In the second scenario, the \"runaway collapse\" is expected, in which two stars are merged while the third companion is kicked away so that an angular momentum is released by the companion from this triple star system (e.g., \\cite{bon04,bon06}). A massive star is formed as a result of a series of events in which two lower-mass stars are merged. This scenario has been also indirectly supported by observational results, in which an MYSO and other young stellar objects (YSOs) are moving from almost the same position on the sky (in Orion KL/BN region, \\cite{rod05,gom05}). An explosion may have happened when these stellar objects were located in the common point with strong dynamical interactions \\citep{bal05}. However, the stellar volume density requested for this scenario ($>$10$^{7}$~pc$^{-3}$) should predict a high frequency of such runaway collapse and explosion events. More observational examples are necessary to widely support this scenario as a major scenario of massive stars.  \\h2o maser emission is one of the most important probes for study on star formation, often based on data obtained by using very long baseline interferometry (VLBI) with high angular and velocity resolution. Analyses of spatial positions, Doppler velocities, and proper motions of individual maser features with a typical size of 1~AU \\citep{rei81} have revealed the 3-D gas kinematics often in very vicinity ($<<$100~AU) of YSOs (e.g.\\ \\cite{cla98,fur00,set02, tor01, tor03, mos05, fur05, god04, god06}). In practice, details of the gas kinematics are complicated but seem to depend on mainly evolutionary status of YSOs. \\h2o maser sources sometimes enable us to measure internal motions of giant molecular clouds by measuring relative bulk motions between clusters of \\h2o  masers that are separated by up to 1~pc but still located within a single antenna beam (e.g., \\cite{ima00,tor01}). Such bulk motions may be owing to, e.g.,  propagation of shock layers from newly-formed H{\\rm II} regions, cloud contraction by the self-gravitation due to a huge mass of a giant molecular cloud, or runaway motions as mentioned above.  The massive-star forming region \\g192 (hereafter G192) is located at a distance of $\\sim$2~kpc, and has massive YSO candidates, one of which is a B2-type star with a giant Herbig-Haro bipolar outflow \\citep{dev99} and a possible circumstellar disk (\\cite{she99}, hereafter SK99; \\cite{she01}, hereafter SCK01; \\cite{ind03}; \\cite{she04}, hereafter SBCSK04). There are two clusters of \\h2o masers separated by $\\sim$0.\\arcsec6 ($\\sim$1200~AU) in the north--south direction (SK99; SBCSK04). The northern and southern clusters exhibit alignments of maser features parallel and perpendicular to two outflows, which are traced by $^{12}$CO emission that are projected roughly in the east--west direction parallel to the Herbig-Haro objects mentioned above (\\cite{she98}, hereafter SWSC98). On the basis of the spatial distribution of \\h2o\\ masers and C$^{18}$O emission tracing high density gas, SK99 and SBCSK04 have proposed that the southern \\h2o\\ maser features are associated with a flattened rotating gas torus around a B2 star, in which the maser velocities are roughly consistent with those in a Keplerian disk.  The Very Long Baseline Array (VLBA) observations of the G192 \\h2o masers were also reported by SBCSK04, but a maser proper motion measurement, especially for the southern cluster of \\h2o masers, was impossible because quenching the maser activity.  Here, we report short-term monitoring observations towards the G192 \\h2o masers with new Japan VLBI network (JVN, Fujisawa \\etal, in preparation), which is newly organized and now consisting of four 20~m telescopes of the VLBI Exploration of Radio Astrometry (VERA), a 45~m telescope of Nobeyama Radio Observatory\\footnote{NRO and VERA observatory are branches of the National Astronomical Observatory, an interuniversity research institute operated by the Ministry of Education, Culture, Sports, Science and Technology.}, a 34~m telescope of Kashima Space Research Center (KSRC), and other 11--64~m dish telescopes in Japan. Fortunately we could detect maser proper motions for both of the maser clusters, leading us to obtaining more direct evidence for the circumstellar disk. Sect.\\ 2 describes the VLBI observations and data reduction. Sect.\\ 3 summarizes analyses of the revealed 3-D kinematics of the \\h2o\\ masers. Sect.\\ 4 discusses the physical conditions and evolutionary statuses of the jet and disk objects in G192. ",
        "conclusions": "We have obtained data of the three-dimensional kinematics of \\h2o  masers in G192 with short-term ($\\sim$2~months) VLBI monitoring observations, which were made with the JVN during 2005 March--June, in different epochs from those previously made with the VLA and the VLBA (SBCSK04). The \\h2o masers are associated with two MYSOs, which have a highly collimated jet and an infalling-rotating gas disk and located at the northern and southern parts in G192, respectively. We have obtained stronger evidence than obtained by previous observations for the existence of an infalling-rotating disk with a radius of $\\sim$1000~AU and a central stellar mass of 5--10~$M_{\\odot}$. The absence of any molecular outflow along the maser alignment found in the previous VLA observation (SK99) may reject the possibility that there exists a jet or an equatorial flow perpendicular to the observed molecular outflow with a total length of 5.7~pc. On the basis of the spatio-kinematical structure of the jet traced by \\h2o maser and $^{12}$CO emission as well as assumption of the same dynamical age of the jet and the disk, we estimated a dynamical age of the jet/disk system in G192 to be about $5.6\\times10^{4}$~yrs. With an assumed gas clump density appropriate for maser excitation, we have speculated that a massive YSO with a mass of $\\sim$8~$M_{\\odot}$ has an accretion disk that has provided the material onto the YSO with a rate of 10$^{-7}-10^{-5}\\: M_{\\odot}$yr$^{-1}$. These partially supports a massive accretion scenario for massive star formation. However, there still exist possibility that this massive YSO will become a higher-mass star by merging with a lower-mass YSO. Observations so as to more clearly elucidate the spatio-kinematics of the \\h2o\\ masers and to directly estimate the pre-shock gas density are crucial for estimating a total mass of the existing massive YSO and destiny of the object. They also should provide clues for elucidating a history of mass accretion and outflow.  \\bigskip We acknowledge all staff members and students who have helped in array operation and in data correlation of the JVN/VERA. H.~I.\\ and T.~H.\\ were supported by Grant-in-Aid for Scientific Research from Japan Society for Promotion Science (16540224).  \\clearpage"
    },
    "0608/astro-ph0608536_arXiv.txt": {
        "abstract": "{ We explore the formation and evolution of debris ejected around quark stars in the Quark Nova scenario, and the application to Soft Gamma-ray Repeaters (SGRs) and Anomolous X-ray Pulsars (AXPs). If an isolated neutron star explodes as a Quark Nova, an iron-rich shell of degenerate matter forms from its crust. This model can account for many of the observed features of SGRs and AXPs such as: (i) the two types of bursts (giant and regular); (ii) the spin-up and spin-down episodes during and following the bursts with associated increases in $\\dot{P}$; (iii) the energetics of the boxing day burst, SGR1806$+$20; (iv) the presence of an iron line as observed in SGR1900$+$14; (v) the correlation between the far-infrared and the X-ray fluxes during the bursting episode and the quiescent phase; (vi) the hard X-ray component observed in SGRs during the giant bursts, and (vii) the discrepancy between the ages of SGRs/AXPs and their supernova remnants. We also find a natural evolutionary relationship between SGRs and AXPs in our model which predicts that the youngest SGRs/AXPs are the most likely to exhibit strong bursting. Many features of X-ray Dim Isolated Neutron stars (XDINs) are also accounted for in our model such as, (i) the two-component blackbody spectra; (ii) the absorption lines around 300 eV; and (iii) the excess optical emission. ",
        "introduction": "Soft $\\gamma$-ray Repeaters (SGRs) are sources of recurrent, short ($t \\sim 0.1\\,\\mathrm{s}$), intense ($L \\sim 10^{44}~\\rm{ergs}$) bursts of $\\gamma$-ray emission with a soft energy spectrum. The normal patterns of SGRs are intense activity periods which can last weeks or months, separated by quiescent phases lasting years or decades. The three most intense SGR bursts ever recorded were the 5 March 1979 giant flare of SGR 0526-66 (Mazets et al.\\ 1979), the similar 28 August 1998 giant flare of SGR 1900+14 and the 27 December 2004 burst (SGR 1806-20). AXPs are similar in nature but with a somewhat weaker intensities and no recurrent bursting. Several SGRs/AXPs have been found to be X-ray pulsars with unusually high spin-down rates of $\\dot{P} / P \\sim 10^{-10}$~s$^{-1}$, usually attributed to magnetic braking caused by a super-strong magnetic field.  Occasionally SGRs enter into active episodes producing many short X-ray bursts; extremely rarely (about once per 50 years per source), SGRs emit  giant flares, events with total energies at least 1000 times higher than their typical bursts. Current theory explains this energy release as the result of a catastrophic reconfiguration of a magnetar's magnetic field. Magnetars are neutron stars whose X-ray emission are powered by ultrastrong magnetic fields, $B \\sim 10^{15}$ G.  Although the magnetar model has had successes, in this paper we present an alternative model which addresses outstanding questions.  We explore these issues within the quark-nova (QN) scenario (Ouyed et al. 2002;  Ouyed et al. 2004;  Keranen, Ouyed, \\& Jaikumar 2005). In our previous studies we have suggested that CFL (color-flavor locked) quark stars could be  responsible for the activity of SGRs and AXPs (Ouyed et al. 2005; Niebergal et al. 2006), and in this paper we extend the QN model by studying in more details the evolution of the QN ejecta (as first discussed in Ker\\\"anen \\& Ouyed 2003).  Following the QN explosion we show that a high metalicity shell can form from the neutron star crust\\footnote{The words ``suspended crust\" has been used in a different context already with respect to quark stars; i.e. the electro-magnetically suspended crust at hundreds of Fermi's above the quark star surface (e.g. Alcock et al. 1986) instead of a few quark star radii as in our case.}. If this matter gains sufficient angular momentum from the central quark star it can form a torus-like structure via the propeller mechanism, which is discussed in the second paper in this series (Ouyed et al. 2006; hereafter refered to as Paper II). However if it does not, as we discuss in this paper, there is instead a thin corotating shell suspended by the quark star's magnetic field pressure. Because the quark star is in a superconducting state it's magnetic field decay is coupled to its period in a manner prescribed by Niebergal et al. (2006), such that as the field decays the shell will drift closer to the star. This movement will shift sections of the shell above the line of neutrality, loosely defined as the polar angle (measured from the equator) above which the magnetic field vector is sufficiently parallel to the gravity vector to allow material to break off from the shell, and fall into the star.  Upon collision with the star the pieces of the shell are instantly converted to CFL quark matter (described by Lugones \\& Horvath (2002)), where excess energy is released as high-energy radiation (Ouyed et al. 2004). We propose this as the mechanism responsible for SGR and AXP bursting activity, and also show that SGRs and AXPs differ primarily by age.  Another class of objects are also explored, namely X-ray Dim Isolated Neutron stars (XDINs; e.g.~Haberl 2004), and we find that these may have evolved from SGRs/AXPs.  This paper is presented as follows. In \\S~\\ref{sec:QN} we review the concepts of a quark-nova and the expansion/evolution of the ejected material into a shell. \\S~\\ref{sec:caseA} contains discussions on the geometry of the resulting shell and its self-similar behaviour in time, leading to pieces of the shell breaking off. The subsequent high-energy bursting from these pieces falling into the quark star is then described in \\S~\\ref{sec:bursting}, along with methods for estimating the age of SGRs/AXPs. In \\S~\\ref{sec:quiescent} the quiescent phase of SGRs/AXPs is discussed in the context of our model, where a simple relation between luminosity and spin-down rate is derived. The changes in period and period derivative observed during bursts are then discussed within the framework of our model in \\S~\\ref{sec:periodchanges}. \\S~\\ref{sec:applications} contains case studies for specific SGRs (1806$-$20 and 1900$+$14), along with explanations for the presence of an iron line during bursts, correlated X-ray to infrared flux ratio, and the hard spectrum seen in giant flares. We then summarize in \\S~\\ref{sec:summary} some outstanding issues in the current understanding of SGRs/AXPs, and show how our model can provide explanations. Finally in \\S~\\ref{sec:xdins}, we discuss XDINs and show how many features are readily explained using our model.  We then conclude in \\S~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  In this paper we present a new model for SGRs and AXPs with possible applications to XDINs. This novel idea relies on the formation of bare CFL quark stars (within the quark-nova scenario) as the underlying engine, with the parent neutron star's crust material surrounding it. Despite the simplifications we have employed to study the dynamics of this ejected shell, specifically how it reaches the equilibrium radius and how it evolves once it is there, our model has a number of attractive features that can account for many observed SGR, AXP, and perhaps XDIN properties.  Missing from our study is the proper treatment of three-dimensional instabilities acting on the shell, such as Raleigh-Taylor or oscillations.  Due to the high conductivity of the shell we expect the Raliegh-Taylor to not operate on very fast timescales, and any radial oscillations to be severly damped as the conducting shell will maintain the magnetic flux enclosed within.  However, azimuthal oscillations may have a significant effect. For a more accuarate description, the dynamics of the shell will require numerical simulations, and is left as future work."
    },
    "0608/astro-ph0608700_arXiv.txt": {
        "abstract": "{ We have developed a regularisation procedure for the direct deprojection and PSF-deconvolution of X-ray surface brightness profiles of clusters of galaxies. This procedure allows us to obtain accurate density profiles in a straightforward manner from X-ray observations (in particular data from {\\it XMM-Newton}, where the PSF correction is important), while retaining information about substructure in the gas distribution, in contrast to analytic modelling of the profiles. In addition to describing our procedure, we present here a detailed investigation of the accuracy of the method and its error calculations over a wide range of input profile characteristics and data quality using Monte Carlo simulations. We also make comparisons with gas density profiles obtained from {\\it Chandra} observations, where the PSF correction is small, and with profiles obtained using analytic modelling, which demonstrate that our procedure is a useful improvement over standard techniques. This type of method will be especially valuable in the ongoing analysis of unbiased and complete samples of X-ray clusters, both local and distant, helping to improve the quality of their results.}  \\authorrunning{J.H. Croston et al.} \\titlerunning{X-ray density profile deprojection} ",
        "introduction": "The density distribution of X-ray-emitting gas in galaxy clusters is an important probe of the underlying mass distribution. By assuming hydrostatic equilibrium and spherical symmetry, the total mass distribution can be obtained from the gas density profile $\\rho_{\\rm gas}(r)$ and temperature profile $T(r)$. The use of galaxy clusters to test structure formation models and as cosmological probes is reliant on these techniques. However, observations also suggest that non-gravitational processes can have important effects on the gas properties of galaxy clusters (e.g. Arnaud, 2005 for a review), which raises questions about our understanding of the links between gas and dark matter properties of clusters. Key diagnostics are provided by the gas entropy distribution, $S=T/\\rho_{\\rm gas}^ {2/3}$, which reflect the specific thermo-dynamical history of the gas. (e.g. Ponman, Cannon \\& Navarro 1999; Ponman, Sanderson \\& Finoguenov 2003; Voit 2005; Pratt, Arnaud \\& Pointecouteau 2006). It is therefore essential to be able to accurately measure the gas distributions of clusters, not only so as to be able to make accurate inferences about their dark matter properties, but also to investigate the departures of the gas properties from the self-similar models expected in the absence of effects such as cooling and non-gravitational heating and feedback.  Surface brightness profiles of the X-ray emission from galaxy clusters have been the primary tool for studying their gas distributions since cluster data became available from early X-ray missions. Until recently, it was thought that (with the exception of central cooling flows) cluster profiles could be adequately fitted with a standard analytical $\\beta$ model profile in surface brightness (e.g. Neumann \\& Arnaud 1999). However, the substantial increase in the data quality of cluster observations in the era of {\\it Chandra} and {\\it XMM-Newton} has revealed that cluster gas distributions are considerably more complex than expected, e.g. from {\\it ROSAT} data. A variety of increasingly complicated analytical models have been used to fit cluster surface brightness profiles (e.g. Pratt \\& Arnaud 2002; Vikhlinin et al. 2006) and to calculate the corresponding gas density profile; however, the need for such a range of models, which are not generally physically motivated, is unsatisfactory. In addition, using an analytical model means that full information about real structure in the cluster profile is lost in converting the profile from surface brightness to emission measure.  One alternative is direct deprojection of measured surface brightness profiles (e.g. Fabian et al. 1981; Kriss et al. 1983; White et al. 1997). If we consider a spherical distribution of cluster gas, with a density distribution (or emission measure distribution, $S_{emit}(r)$) consisting of a series of concentric spherical shells of radii $r_0,r_1,..,r_i,..,r_n$, then using simple geometric considerations (e.g. McLaughlin 1999) we can calculate a 2-D matrix $[R_{proj}]$ whose elements consist of the contributions of each shell $i$ to the projected emission measure in a series of annuli on the plane of the sky $j$ having radii $R_0,R_1,..,R_i,..,R_n$ (which may be the same or different to the $r_i$), e.g. element $R_{proj,ij}$ is the fraction of the emission from shell $i$ that is observed in annulus $j$. For the purposes of this work, we consider a square matrix $R_{proj}$, with $R_0 = r_0$, $R_1 = r_1$, etc, so that the matrix has dimensions $n \\times n$, where $n$ is the number of bins in the observed surface brightness profile. The product $[R_{proj}][S_{emit}]$ is then the 2-D emission measure profile as would be observed by a perfect detector. We can then calculate a second redistribution matrix, $[R_{PSF}]$ that takes into account the effect of the instrumental PSF, i.e. $R_{psf,jk}$ is the fraction of counts from annulus $j$ of the `ideal' profile that are redistributed by the telescope into annulus $k$ in the final observed surface brightness profile (we assume that the energy dependence of the PSF is negligible). The relationship between an observed $S_X$ profile, $C_{obs}(R)$ and the originating 3D emission profile, $S_{emit}(r)$ can therefore be expressed as follows: \\begin{equation} [C_{obs}] = [R_{PSF}][R_{proj}][S_{emit}] \\label{deconv} \\end{equation} where $C_{obs}(i)$ is the surface brightness detected in an annulus $i$, $S_{emit}(j)$ is the emission measure produced by a 3-D shell $j$, $R_{proj}(i,j)$ represents the fraction of emission from shell $j$ that would be observed by a perfect detector to fall in annulus $k$, and $R_{PSF}(i,k)$ a second redistribution of counts from an annulus $k$ of the ideal profile to annulus $i$ of the actual profile resulting from the effect of the PSF. $R_{proj}$ depends only on the geometry of the cluster (see e.g. McLaughlin 1999) and can be easily calculated, and $R_{PSF}$ can also be calculated based on knowledge of the optical properties of a given instrument.  However, it is not straightforward to solve Equation~\\ref{deconv} directly, as it is an inverse problem: small amounts of noise in the data become greatly amplified in the solution for $S_{emit}$. This has limited the usefulness of such an approach for the analysis of cluster profiles, particularly for an instrument such as {\\it XMM-Newton} with a reasonably large PSF. The use of ``onion-skin'' deprojection techniques is common for spectral analysis, and has also been applied to surface brightness profiles from {\\it Chandra} (e.g. David et al. 2001); however, the ``onion-skin'' approach is heavily dependent on the choice of outermost bin. For Equation~\\ref{deconv} to be accurate, it is necessary to take account of the contribution to each annulus from shells outside the outermost annulus chosen for the analysis (we discuss this further in Section~\\ref{method}). Finally, it is necessary to convert the resulting emission measure profiles to gas density profiles taking into account the variation of the cooling function $\\Lambda(T,Z)$ with radius.  This paper describes the application of a regularisation technique to the direct deprojection and PSF-deconvolution of X-ray surface brightness profiles that has allowed us to overcome these limitations. Our method is motivated by, and adapted from, work by Bouchet (1995) on the deconvolution of gamma-ray spectra. In the following section, we describe the regularisation procedure in more detail. We then present in Section~\\ref{simulations} the results of a range of tests using simulated data that demonstrate the reliability of our technique and its applicability to a wide range in data quality and cluster properties. In Section~\\ref{data}, we apply our technique to real data from {\\it XMM-Newton} and {\\it Chandra} and demonstrate that it performs well in comparison to other methods. Finally, in Section~\\ref{dlogne}, we investigate the effect of using gas density profiles obtained from our method on the calculation of the logarithmic slope and thus total mass profiles for clusters. ",
        "conclusions": "We have described a method for obtaining gas density profiles from X-ray surface brightness profiles of galaxy clusters, with the aim of improving constraints on the entropy and total mass distributions of clusters. We first showed using simulated profiles that this method performs well in a range of conditions: \\begin{itemize} \\item The effect of the shape of the density profile was tested using four models: the $\\beta$ model, AB model, KBB model, BB model, as well as for profiles with deviations from a smooth shape. We found that in all cases the agreement between the output density profile and the model was good, with slightly poorer performance in the central regions for the KBB and BB models, due to the choice of an AB model as a scaling function. \\item The global signal-to-noise of the input profiles does affect the accuracy of the output density profiles, with poorer quality data giving less accurate results; however, this effect was fairly small, and a good agreement ($\\chi^{2}$ of 36/24) was obtained for a S/N of 15, corresponding to a profile containing $\\sim 2500$ net counts. \\item The profile binning did not appear to have an important effect on the accurate recovery of the input density profile. \\end{itemize} We then tested the code performance on real {\\it XMM-Newton} data for four clusters: three nearby clusters with published {\\it Chandra} gas density profiles, and one distant cluster. We found that our method resulted in gas density profiles in better agreement with the higher resolution {\\it Chandra} profiles. We also found that our method performed well in the case of the distant cluster, CL0016+16, reproducing a central excess and change of slope at large radii that could not be taken into account using a $\\beta$-model fit. Finally we demonstrated that the mass profiles obtained from our gas density profiles are consistent with those obtained from other methods. We therefore find that our method is suitable for obtaining gas density profiles both from high signal-to-noise observations of nearby clusters and also for distant clusters. This model-independent inversion method will be extremely useful as a consistent and reliable means of obtaining gas density profiles (and subsequently entropy and mass profiles) for ongoing studies of large, unbiased samples of nearby and distant clusters. The deprojection code is available on request to the authors.  \\begin{figure*} \\vbox{ \\hbox{ \\epsfig{figure=a478_mass.ps,width=7.0cm} \\epsfig{figure=a1413_mass.ps,width=7.0cm}} \\epsfig{figure=a1991_mass.ps,width=7.0cm} } \\caption{Mass profile comparison for Abell 478 (top left), Abell 1413 (top right) and Abell 1991 (bottom). In all cases, red + symbols indicate the {\\it XMM} mass profile obtained by deprojection, blue squares the {\\it XMM} mass profile obtained by fitting a model to the surface brightness distribution (taken from Pointecouteau et al. 2004, Pratt \\& Arnaud 2002, and 2005, respectively). Dashed lines indicate NFW fits to the profiles using deprojection, and dotted lines best fits to the profiles using analytic density profiles.} \\label{massplots} \\end{figure*}  We would like to thank the referee for useful comments."
    },
    "0608/astro-ph0608585_arXiv.txt": {
        "abstract": "We present new high resolution ($100\\arcsec$) neutral hydrogen (\\HI) self-absorption images of the Riegel-Crutcher cloud obtained with the Australia Telescope Compact Array and the Parkes Radio Telescope. The Riegel-Crutcher cloud lies in the direction of the Galactic center at a distance of $125\\pm25$ pc.  Our observations resolve the very large, nearby sheet of cold hydrogen into a spectacular network of dozens of hair-like filaments.  Individual filaments are remarkably elongated, being up to 17 pc long with widths of less than $\\sim 0.1$ pc.  The strands are reasonably cold, with spin temperatures of $\\sim 40 $ K and in many places appearing to have optical depths larger than one.  Comparing the \\HI\\ images with observations of stellar polarization we show that the filaments are very well aligned with the ambient magnetic field.  We argue that the structure of the cloud has been determined by its magnetic field. In order for the cloud to be magnetically dominated the magnetic field strength must be $>30~{\\rm \\mu G}$. ",
        "introduction": "\\label{sec:intro} Recent high resolution surveys of neutral hydrogen (\\HI) in the Galactic Plane show that the structure of the warm atomic medium varies significantly on arcminute scales.  The detailed structure of cold \\HI, however, is much more elusive.  Most of our knowledge of cold \\HI\\ comes from observations of \\HI\\ absorption towards continuum sources \\citep[e.g.][]{heiles03b,heiles05}.  Although these measurements are effective for determining gas temperature, they have limited applicability to structural studies because they only trace gas along single lines of sight towards an irregularly distributed array of background sources. Another probe of cold \\HI\\ is \\HI\\ self-absorption (HISA) towards bright background \\HI\\ emission.  This is not, strictly speaking, self-absorption, but absorption of background \\HI\\ emission by cold foreground \\HI.  Surveys like the Canadian Galactic Plane Survey \\citep[CGPS;][]{taylor03} and the Southern Galactic Plane Survey \\citep[SGPS;][]{mcgriff05} have revived studies of HISA, showing that it is a good probe of the structure of cold \\HI\\ on small scales \\citep[e.g.][]{gibson00,kavars05}.  Although these studies are able to image HISA over large areas, and have better estimates of the background emission than lower resolution surveys, it remains difficult to unravel the structure of the cold foreground \\HI\\ from variations in the background emission.  One of the most famous examples of a large scale HISA feature is the Riegel-Crutcher \\citep[hereafter R-C;][]{riegel72} cloud in the direction of the Galactic center.  The R-C cloud was first discovered by \\citet{heeschen55} in his survey of \\HI\\ at the Galactic center and he suggested that the cloud was very extended.  Subsequent observations by \\citet{riegel69} showed that the cold cloud extends over $\\sim 40$ degrees of Galactic longitude and $\\sim 10$ degrees of latitude.  It has even been suggested that the cloud is part of a much larger structure, Lindblad's Feature A, spanning the full 360 deg of Galactic longitude \\citep{lindblad73}.  Estimates of the temperature of the cloud show that despite its striking appearance it is not exceptionally cold, with most estimates of spin temperature ranging between 35 - 45 K \\citep{montgomery95,riegel72,riegel69,heeschen55}. By measuring \\ion{Ca}{2} lines in the spectra of OB stars \\citet{crutcher74} determined that the cloud is between 150 and 180 pc distant.  Further \\ion{Na}{1} observations by \\citet{crutcher84} helped constrain the distance to $125\\pm25$ pc, with an estimated thickness of 1 to 5 pc.  The R-C cloud provides us with one of the best opportunities for imaging the cold neutral medium (CNM) over a large area.  The \\HI\\ background towards the center of the Galaxy is amongst the brightest anywhere in the Galaxy; this direction provides an ideal bright, high column density background against which to image the CNM.  Here we present new high resolution ($\\sim 100\\arcsec$) \\HI\\ images of the R-C cloud, which allow us to study the structure of the CNM on scales of $0.06$ pc to 22 pc.  In this paper we explore the relationship of the R-C cloud to the ambient magnetic field and discuss its implications for the structure of the CNM throughout the Milky Way.  This paper is structured as follows: in \\S \\ref{sec:obs} we discuss the data taking and reduction strategy used.  In \\S \\ref{sec:results} we present images of the R-C cloud and in \\S \\ref{subsec:tempstruc} we derive optical depth and column density maps.  We present evidence of an associated magnetic field in \\S \\ref{sec:bfield}. In \\S \\ref{subsec:bstrength} we discuss the strength of the magnetic field and its effect on the structure of the cloud.  In \\S \\ref{subsec:cnm} and \\ref{subsec:molecules} we comment on the structure of the R-C cloud, its similarity to molecular clouds and the implications for models for the structure of the CNM. ",
        "conclusions": "\\label{sec:conclusions} We have presented new high resolution \\HI\\ self-absorption images of part of the Riegel-Crutcher (R-C) cold cloud at a distance of 125 pc in the direction of the Galactic center.  These images, which are part of the SGPS Galactic Center survey, cover the area $-5\\arcdeg \\leq l \\leq 5\\arcdeg$, $-5\\arcdeg \\leq b \\leq 5\\arcdeg$ with an angular resolution of 100 arcseconds.  The data reveal extraordinary filamentary structure within the R-C cloud.  The cloud seems to be composed of a network of aligned narrow filaments that merge into a smooth mass at the upper left of the cloud.  The filaments are very elongated, with the most extreme example only $0.07$ pc wide and over 16 pc long.  Using saturated absorption profiles in the cloud we estimate a spin temperature for the cloud of $T_s \\sim 40$ K, which is similar to values previously estimated by \\citet{crutcher74} and \\citet{montgomery95}.  Adopting a spin temperature of 40 K for the entire cloud we created images of the optical depth and column density of the cloud, finding an average column density of $\\sim 2 \\times 10^{20}~{\\rm cm^{-2}}$.  If the individual filaments of the cloud are approximated as cylinders, then their average thermal pressure is $\\sim 2 \\times 10^{4}~{\\rm K~cm^{-3}}$.  From measurements of the polarization of starlight from \\citet{heiles00} we have shown that the filaments in the R-C cloud are aligned with the magnetic field of the cloud.  Using the Chandrasekhar-Fermi method we have roughly estimated the magnetic field strength to be $\\sim 60~{\\rm \\mu G}$.  By comparing ratios of gravitational, kinetic and magnetic energy densities we show that the filaments are not self-gravitating.  We suggest that the excellent alignment between the filaments and the magnetic field, as well as the straightness of the filaments, is evidence that the magnetic field dominates over thermal and turbulent motions within the cloud. For the magnetic energy density to exceed the kinetic energy density the magnetic field must be $>30~{\\rm \\mu G}$.  Given the uncertainties associated with the Chandrasekhar-Fermi method the two field strengths are fully consistent.  Directly imaging the structure of the CNM is always difficult.  The R-C cloud, because it is in front of the bright, high column density \\HI\\ background towards the Galactic center, provides us with a unique opportunity to directly image the CNM.  In this case the CNM appears highly filamentary and most closely resembles the high aspect ratio filaments suggested by \\citet{heiles97} to explain Tiny Scale Atomic Structure.  There is no reason to believe that the structure of the R-C cloud is unique, it may in fact be typical for the CNM.  We point out that there are morphological similarities between the R-C cloud and filamentary molecular clouds, which are also characterized by strong magnetic fields.  If gas traced by \\HI\\ self-absorption is either a precursor to molecular gas or formed from dissociation of molecular clouds the morphological similarity is not surprising."
    },
    "0608/astro-ph0608066_arXiv.txt": {
        "abstract": "X-ray spectral state transitions are a key signature of black hole X-ray binaries and follow a well-defined pattern.  We examined 12 {\\it XMM-Newton} observations of the nearby spiral galaxy NGC 1313, which harbors two compact ultraluminous X-ray sources (ULXs), X-1 and X-2, in order to determine if the state transitions in ULXs follow the same pattern.  For both sources, the spectra were adequately fitted by an absorbed power-law with the addition of a low temperature ($kT = 0.1 \\sim 0.3$ keV) disk blackbody component required in 6 of the 12 observations.  As the X-ray luminosity of X-1 increases to a maximum at $3\\times10^{40}$ ergs s$^{-1}$, the power-law photon index softens to 2.5--3.0. This behavior is similar to the canonical spectral state transitions in Galactic black hole binaries, but the source never enters the {\\it high/soft} or {\\it thermal dominant} state and instead enters the {\\it steep power-law} state at high luminosities.  X-2 has the opposite behavior and appears to be in the {\\it hard} state, with a photon index of $\\Gamma = 1.7 - 2.0$ at high luminosity, but can soften to $\\Gamma = 2.5$ at the lower luminosities. ",
        "introduction": "Ultraluminous X-ray sources (ULXs) are point-like, non-nuclear X-ray sources with luminosities above the Eddington limit of a 10 $M_\\sun$ black hole \\citep[$2\\times10^{39}$ ergs s$^{-1}$, for review see][]{fab06}.  Many ULXs show strong variability and are believed to be black hole binaries.  Their X-ray luminosities \\citep{col99,kaa01}, thermal disk emissions \\citep{kaa03}, variation time scales \\citep{str03,dew06} and surrounding emission line nebulae \\citep{kaa04b} suggest that they might have masses of 20--10$^3$ $M_\\sun$, falling into the class of intermediate mass black holes (IMBHs).  However, if the emission is beamed \\citep{kin01,kor02} or exceeds the Eddington limit \\citep{beg02,ebi03}, some or all ULXs may be stellar-mass black holes.  A comparison of the properties of ULXs with the well-known phenomenology of stellar-mass black hole X-ray binaries should help reveal the nature of ULXs.  Spectral state transitions are essential characteristics of black hole binaries \\citep[for a review see][]{mcc06}.  Most X-ray binaries have been observed in both the {\\it hard} and {\\it high/soft} states.  Typically, the intrinsic energy spectrum in the {\\it hard} state is described as a power-law with a photon index $\\Gamma = 1.5 \\sim 2.1$.  The energy spectrum in the {\\it high/soft} consists of a thermal disk component and a power-law with a photon index $\\Gamma = 2.1 \\sim 4.8$.  Some sources also show the {\\it quiescent} state with very low luminosity and the {\\it steep power-law} state with $\\Gamma>2.4$.  The {\\it steep power-law} state differs from the {\\it high/soft} state in that the disk is the dominant spectral component in the {\\it high/soft} state while the power-law component is dominant in the {\\it steep power-law} state.  The luminosity and spectral shape are correlated with higher luminosities corresponding to softer photon indexes, e.g., in Cyg X-1 \\citep{wil06} and XTE J1550$-$564 \\citep{kub04}, and to higher disk temperatures, e.g., in XTE J1550$-$564 \\citep{kub04} and 4U 1630$-$47 \\citep{tom05}.  However, in the {\\it steep power-law} state, the disk luminosity remains constant or decreases while the disk temperature increases \\citep{tom05,mcc06}.  If ULXs are binary systems similar to Galactic black hole binaries, then they should exhibit similar spectral state transitions, unless they are super-Eddington sources in which case the spectral state(s) may differ from the standard ones.  State transitions following the canonical {\\it high/soft} versus {\\it hard} pattern have been reported for two ULXs in IC 342 \\citep{kub01} and Holmberg IX X-1 \\citep{lap01}.  However, the temporal coverage for IC 342 was poor, only two observations, while the coverage for Holmberg IX X-1 was poor and had to be stitched together from several different observatories.  Multiple high-quality {\\it XMM-Newton} spectra have been obtained only for a few sources: two very similar spectra for Holmberg IX X-1 \\citep{wan04} and three observations of Holmberg II X-1 with one showing an unusual low/soft state \\citep{dew04}.  Repeated {\\it Chandra} observations of the Antennae revealed state transitions of several ULXs \\citep{fab03}, but the photon counts were insufficient to permit spectral modeling and the source spectra were studied using hardness ratios -- which are insufficient to disentangle the disk versus power-law components of the spectra.  The availability of 14 archival {\\it XMM-Newton} observations of ULXs NGC 1313 X-1 and X-2 offers an opportunity to investigate state transitions of ULXs.  NGC 1313 is a spiral galaxy at a distance of 4.13$\\pm$0.11 Mpc \\citep{men02}.  Besides the supernova remnant SN 1978K, two ULXs, X-1 and X-2, are found to be associated with NGC 1313.  Both sources exhibit disk emission with temperatures lower than those found in stellar-mass black holes \\citep{mil03}, suggesting they could be IMBHs.  In the luminosity versus temperature ($L$--$kT$) diagram of bright ULXs in nearby galaxies \\citep{fen05}, they appear in the low temperature, high luminosity class, which contains sources most likely to be IMBH candidates.  The maximum luminosities of X-1 and X-2 exceed the Eddington limit for a 20 $M_\\sun$ black hole by a factor of 10, which is difficult to produce in super-Eddington models. \\citet{ram06} find a unique optical counterpart to NGC 1313 X-2 in {\\it HST} images and suggest that the companion star is a B1-B2 giant.  NGC 1313 X-2 resides in an optical nebular supershell \\citep{pak03}, whose kinetic energy is much larger than that could be produced by a typical supernova explosion.  The shell is either a hypernova remnant or a continually powered nebula - possibly by  an outflow from the ULX.  X-ray luminosities and spectral features, as well as optical counterparts, suggest these two ULXs are interesting IMBH candidates.  In the following, we present the details of the observations, our analysis, and our results on the spectral variation in \\S\\ref{sec:data} and discuss the nature of state transitions of these two ULXs in \\S\\ref{sec:diss}. ",
        "conclusions": "\\label{sec:diss}  The strong correlation between X-ray luminosity ($\\log(L_{\\rm X})$) and power-law photon index ($\\Gamma$) for NGC 1313 X-1 is similar to the behavior found in Galactic black hole binaries, e.g., Cyg X-1 \\citep{wil06}, and suggests that ULX NGC 1313 X-1 is a black hole X-ray binary.  The disk temperature in X-1 is constant within the uncertainties.  However, this is consistent with the behavior seen from stellar-mass black hole binaries because the range of the luminosity variation in X-1 is small.  In Galactic black hole X-ray binaries, over 75\\% of the total flux in 2-20 keV band in the {\\it high/soft} state arises from the disk emission component -- the {\\it high/soft} state is a {\\it thermal dominant} state \\citep{mcc06}.  To compare the disk versus power-law luminosity ratios for ULXs, we choose to use the 0.3--10~keV band rather than the 2--20 keV band.  This is motivated because the the disk temperature in ULXs is a factor of 5--10 smaller than that in stellar-mass black holes.  We also compared the luminosity ratios in the 0.1--10 keV and in the 0.4--4 keV band, as suggested by the referee, and found they were very similar to those in the 0.3--10 keV band.  The fraction of total flux arising from the disk for X-1 is always below 50\\%, even allowing for the uncertainties in the measurement of the disk luminosity.  Therefore, X-1 appears to never enter the {\\it high/soft} state and is in either in the {\\it hard} state or the {\\it steep power-law} state. \\citet{mcc06} define the {\\it steep power-law} state by the presence of a power-law component with a photon index larger than 2.4.  In Fig.~\\ref{fig:lg}, observation (10) with $\\Gamma=3.08$ and a disk component of 41\\% fractional flux, showing up as the rightmost point in Fig.~\\ref{fig:lg}, is well consistent with the {\\it steep power-law} state and there are observations ((3), (4), (7), and (9)) with a single power-law spectrum and $\\Gamma = 2.3 \\sim 2.4$ at the hard edge of the {\\it steep power-law} state.  The non-detection of timing noise, except a little non-constant emission, and the absence of quasiperiodic oscillations are consistent with the {\\it steep power-law} state when the power-law flux is greater than 50\\% of the total \\citep{mcc06}.  The luminosity of X-1 appears to saturates when the source enters the {\\it steep power-law} state even though the photon index can soften significantly.  The {\\it steep power-law} state is often associated with the highest luminosities seen from Galactic black hole X-ray binaries.  The five observations that we classify as the {\\it steep power-law state} ((3), (4), (7), (9), (10)) all have luminosities near the maximum observed from X-1.  The other observations of X-1 show photon index $\\Gamma < 2.2$, reduced luminosity, and sometimes a disk component with a fractional flux of $\\lesssim$30\\%.  These properties are all consistent with classification as the {\\it hard} state.  We note that relative small gap in luminosity between the {\\it steep power-law} and the {\\it hard} state is consistent with the behavior seen from some Galactic black hole binaries \\citep{fen04,hom04,mcc06}.  The detection of state transitions from X-1 may help determine whether or not the emission is beamed.  Two models with anisotropic emissions, the mechanical beaming \\citep{kin01} and relativistic beaming \\citep{kor02}, have been raised to interpret the high luminosity of ULXs as radiations from stellar-mass black holes.  The same behavior of state transitions in X-1 as in stellar-mass black holes indicates the emission is from the accretion flow and not a relativistic jet.  The mechanical beaming model needs a near- or super-critical accretion rate in the binary system.  Whether or not such system can produce spectral state transition as observed in normal black hole binaries is unclear.  Other models involving super-Eddington accretion rates, like the advection-dominated optically thick disk \\citep{ebi03} or the radiation-dominated accretion disk \\citep{beg02}, must also address the spectral state transition behavior.  The transition to the {\\it hard} state at luminosities only a factor of 3 below the maximum observed luminosity, as well as a maximum luminosity which exceeds the Eddington limit of a 20 $M_\\sun$ black hole by a factor of 10, may be difficult to explain in such scenarios.  The main physical distinction between the {\\it steep power-law} state and the {\\it high/soft} state (in most models) is the presence of a highly energized corona in the {\\it steep power-law} state.  The absence of the {\\it high/soft} state, or so called {\\it thermal dominant} state, may suggest that NGC 1313 X-1 always has an energetically important corona.  \\citet{goa06} and \\citet{sto06} have suggested that ULXs have optical thick coronae, but the data for NGC 1313 X-1 suggest an optically thin corona.  The nature of the corona in X-ray binaries is still a mystery.  Some models suggest that magnetic fields play an important role in connecting the accretion disk and the corona through magnetorotational instability \\citep[cf.][]{mil00} and the corona acts like a magnetic reservoir \\citep{mer01}.  \\citet{mer03} proposed a coupled magnetic disk-corona solution, in which the disk is stable when it is corona-dominated at high accretion rates. Most bright ULXs have spectra dominated by a soft power-law component \\citep{fen05}, and are thought to be accreting from high-mass companion stars via Roche-lobe overflow \\citep{liu04,kaa04a,ram06,kaa06a,kaa06b}.  Detailed modeling of accretion flows which produce powerful coronae at high accretion rates should be important in understanding ULXs and also in determining whether super-Eddington accretion occurs \\citep{mer03}.  We note that, based on the spectral properties alone, we identify the state of Holmberg II X-1 reported by \\citet{goa06} as the {\\it steep power-law} state since the photon index is 2.6 and less than 15\\% of the flux arises from the disk.  The low level of timing noise is consistent with the state definitions of \\citet{mcc06}, but should be considered in more detail in relation to the timing properties of Galactic black hole binaries in the {\\it steep power-law} state.  At luminosities above $6 \\times 10^{39} \\rm \\, erg \\, s^{-1}$, NGC 1313 X-2 appears to be in the {\\it hard} state, with a photon index of $\\Gamma = 1.7 - 2.0$, see the bottom panel of Fig.~\\ref{fig:lg}.  At lower luminosities, the photon index can reach much softer values.  When fitting with the simple power law model, the points from two clusters with high luminosity associated with hard photon index and low luminosity with soft photon index. The high/hard vs low/soft behavior of X-2 is unlike that of typical Galactic black hole binaries, but a similar, unusual, low/soft state has been observed from GRS 1758-258 \\citep{smi01}.  Galactic black holes in the {\\it hard} state typically produce compact radio jets.  Scaling using the radio to X-ray flux ratio of GX 339-4 \\citep{cor03}, the predicted radio flux density would be $\\sim 5 \\, \\mu$Jy at 8.6~GHz and would be difficult to detect.  However, such a jet could power the optical nebula surrounding X-2."
    },
    "0608/astro-ph0608299_arXiv.txt": {
        "abstract": "\\noindent Solutions of the stationary axisymmetric Einstein equations describing the interior of circularly rotating dust are investigated in order to study their potential applicability as galaxy models. It is shown that such interior solutions cannot be used as global metrics without becoming unphysical in certain regions of space. Although definite results concerning the non-existence of a continuation into a vacuum can only be derived for interior solutions describing rigidly rotating dust, the present analysis exhibits that the case of non-rigidly rotating dust would in general also be inadequate as a physically reasonable galaxy model. ",
        "introduction": "There has been some recent interest in a general relativistic model of Cooperstock \\& Tieu \\cite{2005astro.ph..7619C}, which attempts to fit rotation curves of spiral galaxies without invoking dark matter. This model is constructed such that it gives an excellent fit to observed rotation curves. However, there have been serious concerns about the consistency of the model in general \\cite{2006astro.ph..3791B,2006astro.ph..1191C}.  By using some basic results presented in \\cite{1975JMP....16.1806W}, where solutions of the Einstein equations describing axisymmetric rotating dust were calculated, the applicability of the corresponding solutions to serve as galaxy models is studied in this paper. Such axisymmetric solutions have been used as toy models for galaxies during the last years, e.g. in \\cite{2005astro.ph..7619C} and \\cite{2006astro.ph..2519B}, in order to investigate whether general relativity can give new insights into the problem of the distribution of dark matter in our universe.  The crucial question concerning the usefulness of axisymmetric dust solutions is whether the assumption of galaxies being stationary axisymmetric objects is in some way realistic. Since a galaxy is a dynamic object and a spiral structure is surely not invariant under arbitrary rotations, the above seems barely reasonable. Furthermore, circular flow of dust makes only sense for the galactic disk, but far off the galactic plane such dynamics would become rather unphysical. This has already been pointed out in \\cite{2006astro.ph..1191C}, where it was also shown, that the models in \\cite{2005astro.ph..7619C,2006astro.ph..2519B} indeed describe rigid rotation. Nevertheless, for a slowly rotating galactic disk where details of the spiral structure can be 'washed' out, i.e. where variations in the time coordinate $t$ and the angular coordinate $\\varphi$ are rather small in comparison to variations in the other coordinates, the assumptions described above might yield useful approximations. A restriction to the galactic disk may also be reasonable in connection with the debate about dark matter, since the main question is then whether an inner galactic disk solution can be continued into a vacuum without the need for a dark matter halo. Indeed it will be show below that solutions describing simple axisymmetric configurations of rigidly rotating dust cannot be continued into a vacuum solution without involving singularities.  Furthermore, also axisymmetric dust solutions with non-rigid rotation will be considered. In comparison to rigid rotation models, these solutions have a further degree of freedom expressible by an arbitrary function of one variable. A general answer about transitions into vacuum solutions will not be given, but according to what is currently known about vacuum solutions for Dirichlet boundary data \\cite{1998CMaPh.190..509S,2002PhRvD..65d4006A,1998MPLA...13.1509M}, such transitions will in general exhibit singular behavior on the boundary. ",
        "conclusions": "Besides the restriction that physically realistic solutions describing axisymmetric rotating dust should be restricted in the $Q$ direction because, as already stated in \\cite{2006astro.ph..1191C}, the assumption of circular flow is unrealistic far off the galactic plane, the calculations have shown that it is also required to restrict the solution in the $R$ direction. This unmasks the solutions in \\cite{2005astro.ph..7619C,2006astro.ph..2519B} once again as unphysical, though they seem to work quite well for the interior of the galactic disk. However, according to the expansions in \\ref{SUBSEC:slowrot}, the latter is not accidental, but simply because those metrics approximate non-rigid rotation solutions to first order.  Furthermore, solutions describing rigid rotation cannot simply be matched with exterior vacuum solutions without becoming singular, making them a quite useless mean to claim the non-existence of dark matter, besides the fact that rigid rotation might possibly be quite inappropriate to describe the galactic rotation law, even though the measured rotation curves look like differential rotation due to the relativistic dragging effect observed by rotating objects.  Even when non-rigid rotation of matter is allowed as in \\ref{SEC:arbirot}, things do not really improve in an optimal way. Though, in comparison to the case of rigid rotation, the interior solution can now indeed have an energy density $\\rho$ that vanishes on a compact surface, the resulting problem to find a matching vacuum solution is over-determined. Thus, the resulting metric would in general contain an infinitesimal mass shell enveloping the galactic disk. Nevertheless, it cannot be excluded at the current time, that smooth transitions could exist. Of course, also transitions into other solutions, especially such with more degrees of freedom should be analyzed in the near future, e.g. a dust solution where the $4$-velocity $Z$ is not restricted to be a linear combination of just $\\partial_t$ and $\\partial_\\varphi$. But still, already this assumption would require that metrics of a more general form than \\eref{EQ:metric} have to be analyzed, making the chance to obtain some reasonable results that would finally conclude the debate about the existence of dark matter, quite difficult if not impossible at the current moment.  Nevertheless, the results exhibit that galaxies should be treated as much more complex objects than just a circularly rotating cloud of dust and a final solution to match different solutions of the Einstein equations would also give some definite answers concerning the existence of dark matter. Though this answer is still outstanding, the apparent incapability of solutions containing circularly rotating dust as the only source of gravity to serve as global metrics might give a further indication that more physics is needed to model galaxies, like pressure related phenomena, dark matter halos, effects caused by a massive central black hole or even non-axisymmetric dynamics.  \\appendix"
    },
    "0608/astro-ph0608250_arXiv.txt": {
        "abstract": "{}{ We study the spectrum of the cosmic X-ray background (CXB) in energy range $\\sim$5-100 keV.}{ Early in 2006 the INTEGRAL observatory performed a series of four 30ksec observations with the Earth disk crossing the field of view of the instruments.  The modulation of the aperture flux due to occultation of extragalactic objects by the Earth disk was used to obtain the spectrum of the Cosmic X-ray Background(CXB).  Various sources of contamination were evaluated, including compact sources, Galactic Ridge emission, CXB reflection by the Earth atmosphere, cosmic ray induced emission by the Earth atmosphere and the Earth auroral emission.}  {The spectrum of the cosmic X-ray background in the energy band 5-100 keV is obtained.  The shape of the spectrum is consistent with that obtained previously by the HEAO-1 observatory, while the normalization is $\\sim$10\\% higher.  This difference in normalization can  (at least partly) be traced to the different assumptions on the absolute flux from the Crab Nebulae. The increase relative to the earlier adopted value of the absolute flux of the CXB near the energy of maximum luminosity (20-50 keV) has direct implications for the energy release of supermassive black holes in the Universe and their growth at the epoch of the CXB origin.}{}{} ",
        "introduction": "It is well established \\cite[e.g.][]{giacconi01} that the bulk of the CXB emission below $\\sim$5 keV is provided by the numerous active galactic nuclei (AGN) - accreting supermassive black holes, which span large range of redshifts from 0 up to 6. At these low energies X-ray mirror telescopes can directly resolve and count discrete sources. The resolved fraction drops from $\\approx$80\\% at 2-6 keV to 50-70\\% at 6-10 keV (see Brandt \\& Hasinger, 2005 for review).   At energies above 10 keV the efficiency of X-ray mirrors declines and at present it is impossible to resolve more than few percent of the CXB emission in this regime. It is believed that AGN still dominate CXB at higher energies (at least up to hundreds of keV), although the extrapolation of the low energy data is complicated by the presence of several distinct populations of AGN with different spectra and intrinsic absorption column densities (Sy I, Sy II and quasars, blazars, etc., see e.g. \\citealt{setti89, comastri95,zdziarski96}). At the same time the peak of the CXB luminosity is around 30 keV and accurate measurements of the CXB flux at high energies (even if we can not resolve individual objects) are important to understand the energy release in the Universe and the contribution of various types of objects to it.  Such measurements are complicated because instruments working in the energy range from tens to hundreds keV are often dominated by the internal detector background, caused by the interactions of charged particles with the detector material. To decompose the total background into particle-induced background and the CXB contribution one needs either a very good model of the internal detector background or two observations having different relative contributions of these two components. The latter approach was behind the INTEGRAL observations of the Earth which uses the Earth disk as a natural screen to modulate the CXB flux coming on to the detectors.  A similar approach has already been used for the same purpose in the analysis of early space X-ray experiments. In particular the HEAO-1 observatory used a movable 5 cm thick CsI crystal to partly block the instrument field of view and to modulate the CXB signal (Kinzer et al. 1997, Gruber et al. 1999). Here we report the results of the first INTEGRAL observations of the Earth performed in 2006. During these observations the Earth was drifting through the field of view of the INTEGRAL instruments producing a modulation of the flux with an amplitude of the order of 200 mCrab at 30 keV.  The structure of the paper is as follows. In Section 2 we present the details of the Earth observations with INTEGRAL in 2006. In Section 3 we introduce various components contributing to the light curves recorded by the different instruments. In Section 4 we describe the CXB spectrum derived from the INTEGRAL data. The last section summarizes our findings. ",
        "conclusions": "Using the modulation of the aperture flux by the Earth disk, the INTEGRAL observatory measured the spectrum of the cosmic the X-ray background in the energy range $\\sim$5-100 keV.  The observed flux near the peak of the CXB spectrum (in the $\\nu F_\\nu$ units) at 29 keV is 47 ${\\rm keV^2 cm^{-2} s^{-1} keV^{-1} sr^{-1}} \\pm 0.5 ~(\\pm 1.5 )$. The qouted uncertainties are pure statistic errors and (in parentheses) systematic errors, excluding the uncertainty associated with the choice of the model Crab spectrum. This value is $\\sim$10\\% higher than suggested by Gruber et al., 1999. Note that in the present analysis the absolute flux calibration was done assuming that the spectrum of the Crab nebula in the relevant energy band is described by the expression $dN/dE=10~E^{-2.1}$ phot~s$^{-1}$~cm$^{-2}$~keV$^{-1}$. Any changes in the CXB normalization directly translates into the changes in the energy release by the supermassive black holes in the Universe. These numbers are important for estimating the radiative efficiency of the growing black holes.  In the present analysis the observed level of the atmospheric emission was found to be very close (within 10\\%) to the results of the simulations. Since accurate measurements of the cosmic ray spectra are now available the Earth atmospheric emission could become a useful ``calibration'' source for the instruments operating in the few hundred keV range.  The present observations were made during Solar minimum when the expected level of the atmospheric emission (due to cosmic rays interactions with the atmosphere) is close to the maximum. The ``background'' field was close to the Galactic Plane and part of the data was discarded to avoid contamination of the signal by Earth obscuration of the unresolved Galactic sources.  Future (longer) observations during Solar maximum and with the pointing direction away from the Galactic Plane would be very useful to verify the robustness of the atmospheric emission simulations and to obtain the CXB spectrum in broader energy range."
    },
    "0608/astro-ph0608120_arXiv.txt": {
        "abstract": "{ In this paper we compute the emission coming from the direct recombination of free electrons to a given shell ($n\\geq 2$) during the epoch of cosmological hydrogen recombination. This contribution leads to a total of {\\it one} photon per recombined hydrogen atom and therefore a $\\sim 30-88\\,\\%$ increase of the recombination spectrum within the frequency range $1\\,\\text{GHz}\\leq \\nu \\leq 100 \\,\\text{GHz}$. In particular the Balmer-continuum emission increases the distortion at $\\nu\\sim 690\\,$GHz by $\\sim 92\\%$. With our 100 shell calculations for the hydrogen atom we find that a total of $\\sim 5$ photons per hydrogen atom are emitted when including all the bound-bound transitions, the 2s two-photon decay channel and the optically thin free-bound transitions. Since the direct recombination continuum at high $n$ is very broad only a few $n$-series continuua are distinguishable and most of this additional emission below $\\nu\\lesssim 30\\,$GHz is completely featureless. } ",
        "introduction": "With the advent of accurate observations of the Cosmic Microwave Background (CMB) temperature and polarization anisotropies it becomes increasingly important to understand the dynamics of cosmological recombination on the percent level precision. Several authors \\citep{Leung2004, Dubrovich2005, Chluba2006, Kholu2006} have discussed physical processes leading to percent level corrections to the results of the standard computation \\citep{Seager2000} for the ionization history. Recently, it has been shown that also the treatment of the populations in the angular momentum sub-states of hydrogen has percent level impact on the ionization history, especially at redshifts $z\\lesssim 800-1000$ \\citep[][hereafter RMCS06 and CRMS06 respectively]{Jose2006, Chluba2006b} and it is expected that more physical processes altering the dynamics of cosmological recombination may be realized. \\citet{Lewis2006} have made some first steps towards quantifying the possible impact of percent level corrections to the ionization history on the estimation of cosmological parameters.  In the future it may become possible to directly observe the spectral distortions of the CMB arising during the epoch of recombination. This would in principle open an alternative way to determine cosmological parameters like the baryon and total matter density. Several authors have discussed the distortions arising due to the hydrogen higher level bound-bound transitions \\citep[][RMCS06, CRMS06]{Dubrovich1975,Liubarskii83,Rybicki93,DubroVlad95,Dubrovich2004,Kholu2005,Wong2006}.  In the standard computations within the cosmological context the {\\it direct recombination} to the ground state of hydrogen is neglected, since this transition is so optically thick during the whole epoch of recombination that escape of photons in the Lyman-continuum is considered impossible \\citep{Zeldovich68,Peebles68}. On the other hand electrons can reach {\\it any} of the other levels (with $n\\geq 2$) by {\\it direct recombination} and the emitted continuum photons escape freely, just because after their creation they will not encounter another optically thick transition within the hydrogen atom.  Every successful recombination (i.e. when the electron is reaching the ground-state and remains there) will therefore lead to the emission of {\\it at least} two photons, one from the direct recombination and (because direct recombinations to the ground-state are neglected) {\\it at least} one within the cascade to the ground-state. Here we show that for hydrogen a total of $\\sim 4$ photons per neutral hydrogen atom is produced within the bound-bound cascade and 2s two-photon decay channel. Therefore, one expects a $\\sim 25\\%$ addition to the total number of emitted photons during the epoch of recombination due to the {\\it direct recombination} to states with $n\\geq 2$.  In this paper we discuss the hydrogen recombination spectrum in the frequency range $100\\,\\text{MHz}\\leq\\nu\\leq 3000\\,\\text{GHz}$, with special emphasis on the contribution from the direct recombination lines (for more details on the bound-bound emission within this context see RMCS06 and CRMS06). We use the solution for the recombination history and evolution of the hydrogen populations as obtained in CRMS06. In those computations a maximum of 100 shells following the populations of {\\it all} the angular momentum sub-states separately and including $l$- and $n$-changing collisions was treated. We refer the reader to this paper for more details on the computations. ",
        "conclusions": "In Fig. \\ref{fig:DI_results} we present the full hydrogen recombination spectrum including the emission due to bound-bound transitions, the 2s two-photon decay and direct recombinations to shells with $n\\geq 2$. At high frequencies the free-bound emission shows narrow features corresponding to the Balmer, Paschen, Brackett and Pfund continuua. At low frequencies the free-bound continuua overlap very strongly and the total continuum emission becomes completely featureless. The inlay plot in the upper panel shows the free-bound continuum emission for three well separated high levels. These contributions are very broad and each of them lies by a factor of $\\sim 10$ below the total free-bound emission spectrum. The total level of the emission is only reached after summing over many shells. As the inlay plot in the lower panel illustrates the sum of all contributions (here mainly bound-bound and free-bound) still has a significant modulation ranging from $\\sim 4-35\\%$ in the frequency band $1\\,\\text{GHz}\\leq \\nu \\leq 30\\,\\text{GHz}$. The slope of the free-bound distortion for $1\\,\\text{GHz}\\lesssim \\nu\\lesssim 10\\,$GHz is $\\sim 0.6$ and therefore slightly steeper than for the contribution from the bound-bound transitions alone ($\\sim 0.46$). At $\\nu\\lesssim 1\\,$GHz one still expects an increase of emission when including more than 100 shells.  In Fig. \\ref{fig:DI_results} we also show the free-bound contribution for different values of $\\nmax< 100$. From this one can conclude that within our assumptions the free-bound emission spectrum is converged to a level of better than 1\\% at $\\nu\\gtrsim 5-10\\,$GHz and better than 10\\% for $1\\,\\text{GHz}\\lesssim \\nu\\lesssim 2\\,$GHz. We have also checked how the free-bound emission is depending on the treatment of the angular momentum sub-states and found that at high frequencies the solution, when assuming full statistical equilibrium within the shells for $n>2$, is not very different. At low frequencies, like in the case for bound-bound transitions (see CRMS06), the difference is much larger, again showing that it is important to follow the populations of all the angular momentum sub-states separately.  Figure~\\ref{fig:Photon_number} shows the total number of photons per hydrogen atom emitted during recombination for a given $n$-series (sum of all bound-bound transitions to one shell with fixed $n$) and its continuum (the corresponding bound-free contribution). In the considered case, one can see that for $n\\gtrsim n_{\\rm cr} \\sim 30$ the total emission is dominated by the contribution from direct recombinations. For $n\\lesssim n_{\\rm cr}$ the contribution from bound-bound-transitions dominates. This behavior shows that for the lower shells cascading electrons are very important, whereas for larger $n$ the electrons reach a given shell mainly via direct recombinations. Obviously including more than 100 shells will lead to an increase of $n_{\\rm cr}$, but one does not expect a significant addition to the total number of emitted photons. There is a strong difference in both curves when the computation is done assuming full statistical equilibrium for shell with $n>2$. In this case the curves intersect at $n_{\\rm cr}\\sim 59$ instead of $n_{\\rm cr} \\sim 30$. This shows that following all the angular momentum sub-states the lack of strong redistribution to states with $l\\gg 1$ disfavours the cascade with transitions $\\Delta n\\ll n$.  Knowing the asymptotic behavior of the distortion at low frequencies, $\\Delta I_\\nu=A_0\\,[\\nu/1\\rm GHz]^\\beta\\, 10^{-28}\\rm J\\,m^{-2}\\,s^{-1}\\,Hz^{-1}\\,sr^{-1}$, from Fig.\\ref{fig:DI_results}, one can find the number of photons emitted per hydrogen atom (present density $N_{\\rm H}=\\pot{1.9}{-7}\\,\\rm cm^{-3}$) at $\\nu\\leq\\nu_0$, i.e. $N_\\gamma/N_{\\rm H}=\\frac{4\\pi}{c N_{\\rm H}}\\int_0^{\\nu_0} \\frac{\\Delta I_\\nu}{h \\nu} \\id\\nu=\\pot{3.3}{-2}\\,\\frac{A_0}{\\beta}\\,[\\nu_0/1\\rm GHz]^{\\beta}$. Within the range $1\\,\\text{GHz}\\lesssim \\nu\\lesssim 10\\,$GHz we obtained $A_0^{\\rm bb}\\sim 2.2$ and $\\beta_{\\rm bb}\\sim 0.46$ for the bound-bound and $A_0^{\\rm fb}\\sim 0.59$ and $\\beta_{\\rm fb}\\sim 0.6$ for the free-bound distortion. Considering the bound-bound distortion assuming that the $\\alpha$-transitions give the main contribution and that most of the photons are released at $z_{\\rm em}\\sim 1300$, with $h\\nu_{n_\\alpha}^{\\rm obs}=2\\chi/z_{\\rm em}\\,n^3$, one finds $N^{\\rm bb}_\\gamma(n>n_0)/N_{\\rm H}=\\pot{3.3}{-2}[\\pot{6.6}{6}/z_{\\rm em}]^\\beta\\,\\frac{A_0}{\\beta}\\,n_0^{-3\\beta}\\sim 8.0\\,n_0^{-1.38}$ and $|\\id N^{\\rm bb}_\\gamma/\\id n_0|\\sim 11\\,n_0^{-2.38}\\,N_{\\rm H}$ for the contribution of the whole series for given shell $n_0$. Similarly, for the free-bound emission with $h\\nu_{n}^{\\rm obs}=\\chi/z_{\\rm em}\\,n^2$ one has $N^{\\rm fb}_\\gamma(n>n_0)/N_{\\rm H}\\sim 3.6\\,n_0^{-1.2}$ and $|\\id N^{\\rm fb}_\\gamma/\\id n_0|\\sim 4.3\\,n_0^{-2.2}\\,N_{\\rm H}$, although for $n_0\\gg 1$ due to the large width of the free-bound contributions one expects this estimate to be rather crude. We found that $|\\id N^{\\rm fb}_\\gamma/\\id n_0|\\sim 2.5\\,n_0^{-1.9}\\,N_{\\rm H}$ represents the result shown in Fig.~\\ref{fig:Photon_number} for large $n_0$ very well. One should mention that for $n\\gg 1$ and $|\\Delta n|\\ll n$ induced recombinations and stimulated transitions play an important role.  \\begin{figure} \\centering \\includegraphics[width=0.94\\columnwidth]{eps/Photon_number_comp_ns_2.eps} \\caption{Total number of photons per hydrogen atom released during recombination for a given $n$-series and its continuum for $\\nmax=100$. Also the results assuming full statistical equilibrium within the shells for $n>2$ ($\\nsplit=2$) is presented.} \\label{fig:Photon_number} \\end{figure} \\begin{table} \\caption{Total number of photons per hydrogen atom emitted during recombination within a given series and its continuum. Adding all the contributions one obtains $4.98$ photons per hydrogen atom.} \\label{tab:photon_number} \\centering \\begin{tabular}{@{}lcccc} \\hline \\hline Series & $n$ & bound-bound & continuum & sum \\\\ \\hline Lyman & 1  & $\\pot{4.28}{-1}$ & 0  & $\\pot{4.28}{-1}$ \\\\ Balmer & 2  & $\\pot{7.83}{-1}$  &  $\\pot{2.11}{-1}$ & $\\pot{9.94}{-1}$ \\\\ Paschen & 3  & $\\pot{4.91}{-1}$  &  $\\pot{1.37}{-1}$ & $\\pot{6.28}{-1}$ \\\\ Brackett & 4  & $\\pot{2.38}{-1}$  &  $\\pot{9.59}{-2}$ & $\\pot{3.34}{-1}$ \\\\ Pfund & 5  & $\\pot{1.62}{-1}$  &  $\\pot{7.22}{-2}$ & $\\pot{2.34}{-1}$ \\\\ All & 1-100  & 2.84  &  1.00 & 3.84 \\\\ \\hline \\hline 2s-decay & 1  & 0 & 1.14  & 1.14 \\\\ \\end{tabular} \\end{table} In Table~\\ref{tab:photon_number} we give the values of the total number of photons per hydrogen atom emitted in the Balmer, Paschen, Brackett and Pfund-series and their continuua. Also we listed the total number of emitted photons in all $n$-series and their continuua and the contribution from the 2s two-photon decay. As expected the total number of emitted photons due to free-bound transitions (within the accuracy of our calculation) is identical to the total number of hydrogen atoms. The number of photons emitted within the Balmer-series and Balmer-continuum is also very close to one photon per hydrogen atom. One can see that roughly $2.8$ photons per hydrogen atom are released within the bound-bound transitions. For the free-bound distortion 80\\% of the photon are emitted for $n\\lesssim 14-15$, 90\\% for $n\\lesssim 26$ and 95\\% for $n\\lesssim 41$. On the other hand within the bound-bound transitions one finds 80\\% for $n\\lesssim 6-7$, 90\\% for $n\\lesssim 11-12$ and 95\\% for $n\\lesssim 18$. Focusing on the Lyman-series and the 2s two-photon decay contribution one can see that $\\sim 43\\,\\%$ of all electron go through the Lyman-$\\alpha$ and $\\sim 57\\,\\%$ of electron reach the ground state via the 2s two-photon decay channel. {\\it During the epoch of hydrogen recombination a total of $\\sim 5$ photons per hydrogen atom are produced.} Therefore, recombination of hydrogen slightly increases the specific entropy of the Universe (photons per baryon).%  Observations of the recombination features at frequencies $\\nu\\gtrsim 1412\\,$MHz might become feasible since the strength of the signal under discussion is close to $\\pot{2}{-7}B_\\nu$ and still has variability with well defined frequency dependence on the level of several percent (Fig.~\\ref{fig:DI_results}). Also in this band ($\\lambda<21\\,$cm) one does not expect other sources with similar frequency dependence.  \\acknowledgement{ We wish to thank J.A. Rubi\\~no-Mart\\'{\\i}n for useful discussion on this problem at the initial stage. }"
    },
    "0608/astro-ph0608316_arXiv.txt": {
        "abstract": "Cosmological long-duration gamma-ray bursts (LGRBs) are thought to originate from the core collapse to black holes of stripped massive stars. Those with sufficient rotation form a centrifugally-supported torus whose collapse powers the GRB. We investigate the role of tidal locking within a tight binary as a source of the necessary angular momentum. We find that the binary orbit must be no wider than a few solar radii for a torus to form upon core collapse. Comparing this criterion to the observed population of binaries containing two compact objects suggests that rotation may have been important in the formation of up to 50\\% of the observed systems. As these systems created a neutron star and not a black hole they presumably did not produce highly luminous GRBs. We suggest instead that they make the subset of GRBs in the relatively local universe which have much lower luminosity. ",
        "introduction": "The link between long-duration gamma-ray bursts (GRBs) and stellar collapse is now firmly established (Woosley 1993; Hjorth et al. 2003;  Stanek et al. 2003).  In particular long-duration GRBs appear to originate in type Ic supernovae and frequently in hypernovae (core collapse events with an order of magnitude more energy than classical supernovae).  The nature of the short duration ($<2s$) bursts is far more uncertain,  although their origin in populations of all ages (e.g. Gehrels et al. 2005; Berger et al. 2005; Prochaska et al. 2006) can be explained if they are caused by the final merger of a tight binary system of neutron stars (NS) or black holes (e.g. NS-NS, NS-BH). In addition to these two main classes of bursts, which are distinguished primarily based on their observed duration, there is also evidence for further subtypes. For example as  well as the very energetic GRBs which originate from high redshifts  (e.g. a mean redshift of $\\sim 2.8$; Jakobsson et al. 2006) there is also  a further population of low luminosity events which can be seen  only in the relatively local universe. The prototype for this class is GRB~980425/SN~1998bw (Galama et al. 1998) which occurred only 35 Mpc away, while GRB~031203/SN~2003lw (at $\\sim$ 450 Mpc; Watson et al. 2004; Malesani et al. 2004) and most recently  GRB~060218/SN~2006aj (at $\\sim$ 130 Mpc; Modjaz et al. 2006; Pian et al. 2006) also lie in this class. These low luminosity events have isotropic equivalent luminosities  of only $10^{48 - 49}$ ergs compared with energies of up to $10^{54}$ ergs for the most luminous GRBs. Thus, while  the number of observed systems is significantly lower, their space density is likely to be much higher than for the more luminous bursts.  Regardless of the object responsible for the GRB (e.g. NS-NS binary or collapsing star) the most popular model for the creation of the burst is essentially the same - extreme accretion rates on to a newly--formed compact object. This accretion is thought to be fuelled by a massive (0.1-10 M$_{\\odot}$) torus which forms if the infalling material has too much angular momentum to accrete directly on to the central compact object. In NS-NS mergers, it is relatively common for mergers to produce a torus (see e.g. Ruffert \\& Janka 1999;Rosswog \\& Davies 2002).  For massive single stars, it is not clear whether sufficiently high central rotation rates may be maintained to produce a torus on core collapse (e.g. MacFadyen \\& Woosley 1999;Petrovic et al. 2005), although it has been suggested that rapidly rotating metal--poor stars can retain sufficient angular momentum (Yoon \\& Langer 2005; Woosley \\& Heger 2006). Alternatively, binary scenarios suggest a way of removing the envelope and providing a source of angular momentum (e.g. Izzard et al. 2004; Podsiadlowski et al. 2004). We explore this mechanism further in this paper.   Specifically we examine the idea that the core of a massive star may be spun up by tidal locking within a tight binary.  The necessary binary separations are a few solar radii or less.  The hydrogen envelope of the massive star would therefore have been lost earlier (probably in a common envelope phase). Systems sufficiently tight may undergo a further period of mass transfer, similar to that seen in Cyg-X2 (e.g. Davies, Ritter \\& King 2002). We subsequently examine known compact binary systems (NS-NS, NS-WD) and conclude that, at the time of the supernova, approximately 50\\% may have been sufficiently tight to necessitate the formation of a disc.  We suggest that low and high luminosity GRBs may be distinguished by the formation of either a NS (low luminosity) or BH (high luminosity). The only compact object -- compact object binaries discovered so far in the Milky Way consist of neutron stars and white dwarfs. However similar systems, with more massive cores could produce black holes, and thus luminous GRBs. ",
        "conclusions": "We have examined the role of tidal locking in a tight binary system and its implications for the formation of a centrifugally-supported torus upon core collapse. By examining the  distribution of separations of known compact object binaries (NS-NS, NS-WD), we conclude that up to $\\sim$ 50\\% of the systems could, at the time of the second  supernova, have been sufficiently rapidly rotating to create a torus upon collapse. We subsequently investigated the implications of these observations for the  formation of GRBs. None of the observed systems contain a BH, which is  thought to be essential to the formation of the highly luminous GRBs. We suggest that higher mass analogues of the observed systems can form GRBs, but that the collapse of less massive stars leading to neutron stars might result in  a lower luminosity GRB, similar to those that have been observed in the local  universe."
    },
    "0608/astro-ph0608510.txt": {
        "abstract": "% Context {$\\alpha$\\,Centauri is our closest stellar neighbor, at a distance of only 1.3\\,pc, and its two main components have spectral types comparable to the Sun. This is therefore a favorable target for an imaging search for extrasolar planets. Moreover, indications exist that the gravitational mass of $\\alpha$\\,Cen\\,B is higher than its modeled mass, the difference being consistent with a substellar companion of a few tens of Jupiter masses.} % Aims {We searched for faint comoving companions to $\\alpha$\\,Cen\\,B. As a secondary objective, we built a catalogue of the detected background sources.} % Methods {We used the NACO adaptive optics system of the VLT in the $J$, $H$, and $K_s$ bands to search for companions to $\\alpha$\\,Cen\\,B. This instrument allowed us to achieve a very high sensitivity to point-like sources, with a limiting magnitude of $m_{Ks} \\approx 18$ at 7\" from the star. We complemented this data set with archival coronagraphic images from the HST-ACS instrument to obtain an accurate astrometric calibration.} % Results {Over the observed area, we did not detect any comoving companion to $\\alpha$\\,Cen\\,B down to an absolute magnitude of 19-20 in the $H$ and $K_s$ bands. However, we present a catalogue of 252 background objects within about 15\" of the star. This catalogue fills part of the large void area that surrounds $\\alpha$\\,Cen in sky surveys due to the strong diffused light. We also present a model of the diffused light as a function of angular distance for the NACO instrument, that can be used to predict the background level for bright star observations.} % Conclusions {According to recent numerical models, the limiting magnitude of our search sets the maximum mass of possible companions to 20-30 times Jupiter, between 7 and 20\\,AU from $\\alpha$\\,Cen\\,B.} % ",
        "introduction": "Our closest stellar neighbor, the $\\alpha$\\,Cen visual triple star ($d = 1.34\\,$pc), is an extremely attractive target for an extra-solar planet search. The main components, $\\alpha$\\,Cen\\,A (\\object{HD 128620}) and B (\\object{HD 128621}), are G2V and K1V solar-like stars, (e.g. have solar-like asteroseismic oscillation frequencies), while the third member is the red dwarf {\\it Proxima} (M5.5V). % In all imaging planet searches, the main difficulty is in retrieving the planetary signal in the bright point-spread function (hereafter PSF) from the star. The proximity of $\\alpha$\\,Cen is a clear advantage as it allows a faint companion to be easily separated angularly from the star itself down to orbital distances as close as a few astronomical units. After a discussion of the potential for companions around $\\alpha$\\,Cen (Sect.~\\ref{pot}), we present our adaptive optics observations (Sect.~\\ref{obs}) and the existing data from the HST archive (Sect.~\\ref{hst_archive}), followed by the catalogue of the detected sources (Sect.~\\ref{cat}) and a discussion (Sect.~\\ref{discussion}).  %___________________________________ % ",
        "conclusions": "We have obtained deep adaptive-optics images of the close environment of $\\alpha$\\,Cen~A and B. From these images, we did not identify any comoving companion, but we assembled a catalogue of 252 faint background objects. Within the explored area, this negative result sets an upper mass limit of 20-30\\,$M_J$ to the possible substellar companions orbiting $\\alpha$\\,Cen\\,B. If companions of $\\alpha$\\,Cen~B exist, they are likely to orbit close to the star (within 5 AU) and to be less massive than a few times Jupiter (from radial velocity surveys). They could also be fainter than the current imaging search-limiting magnitude, but from Baraffe et al.~(\\cite{baraffe03}), a 5\\,Gyr-old, intermediate-mass exoplanet (5\\,$M_J$) around $\\alpha$\\,Cen has an apparent magnitude of $m_H \\approx 27$, and $m_K \\approx 39$, fully out of reach of the deepest imaging searches. Stars younger than $\\alpha$\\,Cen could be more favorable targets, as the brightness of massive exoplanets is predicted to decrease steeply with time. However, the faintness of the detected background sources confirms the capabilities of modern adaptive optics instruments like NACO for exploring the close environment of very bright stars and for searching for massive exoplanets. %__________________________________Acknowledgements"
    },
    "0608/astro-ph0608667_arXiv.txt": {
        "abstract": "We present Very Large Array observations at 1.3 cm of several water maser detections obtained by previous single-dish studies of Bok globules in the Clemens \\& Barvainis (1988; CB) catalog. We report water maser emission in CB 3 (CB3-mm), CB 54 (IRAS 07020$-$1618), CB 101 (IRAS 17503$-$0833), and CB 232 (IRAS 21352+4307), and non-detection towards CB 65 (IRAS 16277$-$2332) and CB 205 (IRAS 19433+2743).  These are the first reported interferometric observations of water masers in Bok globules of the CB catalog. We also present single-dish observations of millimeter and centimeter spectral lines  towards CB 101 (IRAS 17503$-$0833) and CB 65 (IRAS 16277$-$2332). All the maser emission seems to be associated with star forming regions hosting bipolar molecular outflows, except IRAS 17503$-$0833 in CB 101, which we suggest to be a possible Mira evolved star, and IRAS 16277$-$2332 in CB 65, of unknown nature. We have used the precise position of the maser emission to derive information about the powering source of the masers. By analyzing the spatio-kinematical distribution of the water masers, we confirm the millimeter source CB 3-mm as the most likely powering source of the CB 3 masers. We propose the near-IR source CB 232 YC1-I as the best candidate for pumping the maser emission observed in CB 232, while in CB 54, we suggest that the pumping source of the masers could be located at the position of an elongated feature observed in near-infrared maps. ",
        "introduction": "Bok globules \\citep{Bok47} are small ($\\la 1$ pc), isolated clouds, some of which are sites of low- and intermediate-mass star formation \\citep{Yun90}. They have been observed to be in different evolutionary stages, from starless dark cores \\citep{Yun96,Kan97} to sites hosting T-Tauri stars \\citep{Yun95,Lau97}.  These globules have traditionally been considered important laboratories for the study of star formation processes, since their small size and relative simplicity make these processes less likely to be affected by confusion from different generations of young stellar objects. Another important aspect of star formation in Bok globules is that it may explain the somewhat puzzling existence of pre-main-sequence objects apparently not related with known molecular clouds. At least some of these objects may have originated in already dispersed Bok globules \\citep{Lau97}.   Among the observational tools used in star-formation studies, water masers have proved to be a very powerful one. Water maser emission at 22 GHz is a good tracer of mass-loss activity in young stars \\citep{Rod80,Fel92,Xia95,Deb05}. In the low-mass young stellar objects (YSOs) this signpost of mass loss phenomena occurs in their most embedded phase, at the earliest stages of evolution. This period is characterized by the presence of the disk/young stellar object/outflow system \\citep{Shu87} and a large amount of circumstellar dust. In the low-mass YSOs framework, the youngest evolutionary stage is represented by Class 0 sources \\citep{And93}, which show the most powerful and collimated molecular outflows (e.g., \\citealt{Hir06}, and references therein), centimeter emission (e.g. \\citealt{Ang95}), and host strong submillimeter and millimeter continuum emission \\citep{And93,And94}.  Water maser studies have shown that this emission provides a good characterization of the age of low-mass YSOs (e.g. \\citealt{Fur01}), with Class 0 sources being the most probable candidates to harbor water maser emission, due to the interaction of powerful jets with a larger amount of circumstellar material. This fact makes sources that host water maser emission good candidates for being in a very early stage of its evolution. In addition, several water maser surveys towards low-mass YSOs have shown that this emission tends to be located close (within several hundred AU) to the central powering source \\citep{Che95, Cla98, Fur00, Fur03}. These characteristics make the water maser phenomenon well-suited to derive information about the location of the exciting source of the mass-loss phenomena observed in young star forming regions. Moreover, maser emission can be a useful tool for studying the physical conditions and kinematics of the gas surrounding the most embedded YSOs at very high angular resolution ($\\leq$ 1 mas), due to the very specific conditions which pump the maser emission (hundreds of K and n$_{\\rm H_{2}}\\simeq 10^{8}-10^{10}$ cm$^{-3}$) \\citep{Cla98,Fur00,Pat00,Tor03,God05,Mar05,Vle06}. These conditions can be generated in shocked gas compressed by winds \\citep{Che95,Tor97,Fur00,Mos00}, as well as in circumstellar disks \\citep{Fie96,Tor96,Tor98,Set02,Bra03}. This dichotomy has also been suggested to be a possible evolutionary sequence, with water masers tracing gravitationally bound material (e.g., circumstellar disks) in the youngest sources, and outflows in more evolved YSOs \\citep{Tor97,Tor98}. Recently, the hot dense infalling gas after the accretion shock has also been proposed as a good environment for pumping the maser emission \\citep{Men04}.  Therefore, there are different evolutionary aspects related to water maser emission (e.g., the likely occurrence of masers in Class 0 YSOs, and their association with either disks and jets). Bok globules are appropriate sites to study evolution of YSOs, since they are usually identified as dark patches in optical images, without any selection criterion related to possible star-forming activity (\\citealt{Cle88}, hereafter CB; \\citealt{Bou95}), and therefore they may span a large range of evolutionary stages. Surprisingly, studies of water masers in Bok globules are relatively scarce. Recently, our group undertook a sensitive and systematic single-dish survey for water masers in 207 positions within Bok globules \\citep{Gom06}, using NASA's 70m antenna in Robledo de Chavela (Spain), which provided six new detections in \\objectname{CB 34}, \\objectname{CB 54}, \\objectname{CB 65}, \\objectname{CB 101}, \\objectname{CB 199}, and \\objectname{CB 232}. Before this survey, only \\citet{Scap91} conducted a search for water masers specifically in Bok globules, with 80 target positions and only one detection towards \\objectname{CB 3}. Other surveys for water masers included some Bok globules within their target positions \\citep[e.g.,][]{Fel92,Pal93,Wou93,Per94,Cod95}, although they did not find any detection in these types of sources.   Using as a reference the catalogs of Bok globules of CB in the northern hemisphere, and \\citet{Bou95} in the southern one (probably the more complete ones available), we are only aware of  9 Bok globules in those catalogs that are apparently associated with water maser emission: The six new detections mentioned above in the survey by \\citet{Gom06}, the detection towards CB 3 reported by \\citet{Scap91}, as well as the detections obtained by \\citet{sch75} and \\citet{Nec85} in  \\objectname{CB 39} and \\objectname{CB 205} respectively. All these reported detections of masers in Bok globules present single-dish data. However, interferometric high-angular resolution observations are necessary to pinpoint accurately the actual pumping source of water maser emission among several candidates, and to determine whether maser emission in these Bok globules tends to trace collimated jets or circumstellar disks.  In this paper we present for the first time interferometric, high-angular resolution observations using the Very Large Array (VLA) of the National Radio Astronomy Observatory\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} of some of these water masers in Bok globules. This paper is structured as follows: in \\S 2 we describe our observations and data processing; in \\S 3 we present our observational results and we discuss them; we summarize our conclusions in \\S 4. ",
        "conclusions": "In this work we have presented interferometric observations using the VLA,  of several sources  of water maser emission detected as part of a single-dish survey performed with the Robledo-70m antenna towards Bok globules. The general conclusions we have obtained are the following:  \\begin{itemize} \\item We have detected water maser emission with the VLA in CB 3, CB 54, CB 101 and CB 232. No water maser emission was detected towards CB 65 or CB 205.  \\item These water masers  are associated with star forming regions showing bipolar molecular outflows, except for CB 101 (IRAS 17503$-$0833), which we propose is an evolved object, possibly a Mira star, and CB 65 (IRAS 16277$-$2332), whose nature is unknown.  \\item All of the Bok globules associated with star formation present multiple stellar systems in different evolutionary stages. Accurate positions of water masers has helped us to get information about the powering sources in the case of CB 3, CB 54, and CB 232. We propose the millimeter source CB 3-mm and the near-IR source CB 232 YC1-I as the best candidates for pumping the maser emission in CB 3 and CB 232 respectively. In the case of CB 54, we propose a new young object (CB 54 YC1-SW), seen as an  elongated feature in near-infrared images, as the powering source of the maser emission.  \\item The maser emission in CB 3 is distributed along the direction of the molecular outflow, and seems to be tracing the inner part of a collimated jet. The position-velocity distribution of the maser emission shows a wave-like structure, which is consistent with a precessing jet.    \\end{itemize}"
    },
    "0608/astro-ph0608498_arXiv.txt": {
        "abstract": "We present a search for galaxy clusters in the fields of three bona-fide short GRBs (050709, 050724, and 051221a) and the putative short burst GRB\\,050911 using multi-slit optical spectroscopy.  These observations are part of a long-term program to constrain the progenitor age distribution based on the fraction of short GRBs in galaxy clusters and early-type galaxies.  We find no evidence for cluster associations at the redshifts of the first three bursts, but we confirm the presence of the cluster EDCC 493 within the error circle of GRB\\,050911 and determine its redshift, $z=0.1646$, and velocity dispersion, $\\sigma\\approx 660$ km s$^{-1}$.  In addition, our analysis of \\swift/XRT observations of this burst reveals diffuse X-ray emission coincident with the optical cluster position, with a luminosity, $L_X\\approx 4.9\\times 10^{42}$ erg s$^{-1}$, and a temperature, $kT\\approx 0.9$ keV.  The inferred mass of the cluster is $2.5\\times 10^{13}$ M$_\\odot$, and the probability of chance coincidence is about $0.1-1\\%$, indicating an association with \\grb\\ at the 2.6-3.2$\\sigma$ confidence level.  A search for diffuse X-ray emission in coincidence with the fifteen other short GRBs observed with XRT and {\\it Chandra} reveals that with the exception of the previously-noted cluster ZwCl 1234.0+02916 likely associated with GRB\\,050509b, no additional associations are evident to a typical limit of $3\\times 10^{-14}$ erg s$^{-1}$ cm$^{-2}$, or $M\\lesssim 5\\times 10^{13}$ M$_\\odot$ assuming a typical $z=0.3$.  The resulting fraction of short GRBs hosted by galaxy clusters of about $20\\%$ is in rough agreement with the fraction of stellar mass in clusters of $\\sim 10-20\\%$. ",
        "introduction": "\\label{sec:intro}  Several lines of evidence indicate that the progenitors of short-duration, hard-spectrum gamma-ray bursts (GRBs) are related to an old stellar population.  These include the localization of GRB\\,050724 \\citep{bpc+05,pbc+06}, and most likely GRB\\,050509b \\citep{gso+05,bpp+06}, to bright elliptical galaxies, the lack of supernova emission in several low redshift short GRBs \\citep{hsg+05,ffp+05,bpp+06,sbk+06}, and the location of GRB\\,050709 outside of any star forming region in its host galaxy \\citep{ffp+05}. While this supports the popular notion that the progenitors are compact objects binaries (DNS or NS-BH; e.g., \\citealt{elp+89,npp92,rrd03}), the lack of direct observations (e.g., gravitational waves or a sub-relativistic, radioactive component: \\citealt{lp98,kul05}), suggests that a more detailed understanding of the progenitor population has to rely on statistical studies.  In this vein, \\citet{ngf05} and \\citet{gp06} argue that the redshift and luminosity distributions of \\swift\\ and BATSE short GRBs, are inconsistent with the nominal merger time distribution of DNS binaries in the Milky Way, $P(\\tau)\\propto \\tau^{-1}$ \\citep{clm+04}.  They further conclude that the typical age of the progenitors is old, $\\gtrsim 4$ Gyr \\citep{ngf05}.  Similarly, \\citet{gno+05} and \\citet{zr06} propose that the relative fractions of short GRBs in early- and late-type galaxies should constrain the progenitor lifetime distribution.  This test derives from the fact that, on average, stars in early-type galaxies form earlier than in late-type galaxies.  An an extension of this idea, the preponderance of short GRBs in galaxy clusters can also be used to constrain the age distribution and nature of the progenitor population.  In the framework of $\\Lambda$CDM cosmology, high-resolution numerical simulations suggest that the oldest stars reside in dense galaxy cluster environments, typically within $\\sim 150$ kpc of the cluster center \\citep{ws00,dsw+06}.  This is supported by observations, which indicate that the difference in formation epochs for stars in cluster and field ellipticals is $\\sim 1-3$ Gyr (e.g., \\citealt{brc+98,ksc+02,tmb+05}).  Since $M$* in clusters is larger than in the field (e.g., \\citealt{bbb+06}), this leads to an even more pronounced difference in star formation history than just the difference between early- and late-type galaxies discussed in \\citet{zr06}.  In addition, the early-type fraction in clusters, $\\sim 60\\%$, is twice as high as in the field \\citep{dre80,wgj93}, suggesting that the fraction of short GRBs in clusters is intimately related to their rate in early-type galaxies.  This is of particular importance in cases where the positional accuracy of the burst is not sufficient to associate it with a particular galaxy, but may be sufficient to associate it with a galaxy cluster.  Finally, the specific frequency of globular clusters is at least a factor of a few higher in bright cluster ellipticals than in field galaxies \\citep{har91}.  Thus, the fraction of short bursts in galaxy clusters may shed light on whether globular clusters are an efficient site for the formation of short GRB progenitors, as proposed by \\citet{gpm06}.  A complete search for clusters is also important from an observational point of view.  These associations can be made at high significance based on the prompt $\\gamma$-ray positions alone (typically, $2-3'$), whereas associations with individual galaxies require accurate positions ($\\sim 1-10''$) from the optical, radio, or X-ray afterglow. Since the afterglow brightness correlates with the circumburst density\\footnotemark\\footnotetext{The X-ray luminosity is usually considered to be insensitive to the circumburst density ($n$), but this is only true when the X-ray band is located above the synchrotron cooling frequency, requiring a relatively high density ($n\\gtrsim 10^{-3}$ cm$^{-3}$) to begin with.  At lower densities, typical of the intracluster medium, the afterglow luminosity in all bands is expected to be low.} the latter approach may produce an observational bias in favor of gas-rich or disk galaxies.  To date, three associations of short GRBs with galaxy clusters have been claimed.  GRB\\,050509b appears to reside in the cluster ZwCl 1234.0+02916 at $z=0.226$ \\citep{peh+05,bpp+06}, GRB\\,050813 is apparently associated with a cluster at $z\\sim 1.8$ (\\citealt{ber06}; Gladders et al. in prep.), and GRB\\,790613 may be associated with the cluster Abell 1892 at $z\\approx 0.09$ \\citep{gno+05}.  The statistical significance of these associations is $\\sim 3\\sigma$.  Motivated by these considerations we began the first systematic search for galaxy clusters hosting short GRBs, using multi-slit optical spectroscopy and archival X-ray observations.  This is part of a long-term program to constrain the age distribution of the progenitors using the properties of their large-scale environments.  Here we present spectroscopy in the fields of GRBs 050709, 050724, 050911, and 051221a.  We also re-analyze all of the available X-ray observations of short GRB fields to search for diffuse emission from hot intracluster gas associated with potential clusters.  The layout of the paper is as follows.  The optical observations are described in \\S\\ref{sec:spec}, and the X-ray analysis in presented in \\S\\ref{sec:xray}.  In \\S\\ref{sec:res} we summarize the results of our search, including a determination of the optical and X-ray properties of the cluster EDCC 493, which coincides with GRB\\,050911.  We draw initial conclusions in \\S\\ref{sec:disc}.  Throughout the paper we use the standard $\\Lambda$CDM cosmology with $H_0=71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.27$, and $\\Omega_\\Lambda=0.73$. ",
        "conclusions": "\\label{sec:disc}  We present the first systematic search for galaxy clusters hosting short GRBs using multi-slit optical spectroscopy in the fields of GRBs 050709, 050724, 050911, and 051221a, and a re-analysis of all publicly available X-ray observations.  Future papers in this series will present optical spectroscopy of additional short GRB fields and a detailed analysis of the galaxy cluster statistics.  No apparent clusters are found in the fields of GRBs 050709, 050724, and 051221a from optical and X-ray observations.  In the error circle of the putative short burst \\grb\\ we show that the cluster EDCC 493 has a mean redshift, $z=0.1646$, a velocity dispersion, $\\sigma=660$ km s$^{-1}$, an X-ray temperature, $kT=0.9$ keV, and a luminosity, $L_X=4.9\\times 10^{42}$ erg s$^{-1}$.  These values are typical for poor clusters, and the inferred mass is about $2.5\\times 10^{13}$ M$_\\odot$.  We estimate that the chance probability of finding such a galaxy cluster in the BAT error circle of \\grb\\ is about 0.1-1$\\%$. This result highlights the ability to associate short GRBs with galaxy clusters based on $\\gamma$-ray positions alone, thus removing a potential bias in favor of associations with gas-rich galaxies when relying on afterglow positions.  It has been suggested that the relative fraction of early- and late-type host galaxies of short GRBs can be used to constrain the age distribution of the progenitor population \\citep{zr06}.  Making the association between \\grb\\ and EDCC 493, and using all of the available observations to date, we find that of the \\swift\\ short GRBs five appear to be associated with early-type galaxies (050509b, 050724, 050813, 050911, 060502b) and two are associated with late-type galaxies (050709, 051221a).  Taken at face value, this would argue for an age distribution, $P(\\tau)\\propto \\tau^n$ with $n\\sim 2$ \\citep{zr06}.  Naturally, in making a more accurate derivation of this value, one has to take into account the respective probability of association for each burst\\footnotemark\\footnotetext{For example, the probability of association for GRB\\,050509b is estimated at 3-4$\\sigma$ \\citep{gso+05,peh+05,bpp+06}, while that for GRB\\,060502b is only $2\\sigma$ \\citep{bpc+06}.}.  In fact if we consider only secure associations, the relative numbers are $1\\!:\\!2$ instead of $5\\!:\\!2$, leading to $n\\sim -1$.  Independent of associations with individual galaxies, the fraction of short GRBs in clusters is also of interest in assessing the age distribution.  Of the sixteen available bursts, three have claimed cluster associations\\footnotemark\\footnotetext{This fraction does not change if we include the claimed cluster association for GRB\\,790613 out of four IPN short burst error boxes searched by \\citet{gno+05}.} (050509b, 050813, 050911).  Within the uncertainty, this fraction is in rough agreement with the value of $\\sim 10\\%$ for the overall fraction of stellar mass in galaxy clusters \\citep{fhp98}, or roughly $20\\%$ in clusters equal to or more massive than EDCC 493 \\citep{ebc+05}.  Of course, not all stars are capable of producing short GRBs, but assuming that the initial mass function and binary fractions are independent of galaxy properties, the total stellar mass provides a good proxy for the mass in short GRB progenitors.  Thus, at the present there is reasonable agreement between the fraction of short GRBs and the baseline fraction of stellar mass in galaxy clusters.  Finally, we note that the frequency of short GRBs in galaxy clusters may reflect a potential association with globular clusters.  The latter are thought to provide an efficient environment for the production of DNS binaries, and may account for a substantial fraction of all short GRB progenitors \\citep{gpm06}.  In particular, the specific frequency\\footnotemark\\footnotetext{The specific frequency is defined as the globular cluster number normalized to $M_V=-15$ mag \\citep{hb81}.} of globular clusters increases significantly from a value of $\\sim 1$ for S+Irr galaxies to $\\sim 4$ for E+S0 galaxies, and $\\sim 12$ for cD galaxies \\citep{har91}.  Given that massive ellipticals are over-represented in galaxy clusters compared to the field, we expect that an association with globular clusters will increase the fraction of galaxy cluster associations compared to the baseline level of $\\sim 10-20\\%$ indicated above.  Thus, continued searches for galaxy clusters hosting short GRBs, and more detailed predictions for the expected fraction as a function of cluster mass and redshift, may hold the key to a clearer understanding of the progenitor population."
    },
    "0608/astro-ph0608451_arXiv.txt": {
        "abstract": "{First steps have been taken in a more comprehensive study of the dependence of observables in Type I X-ray bursts on uncertain (p,$\\gamma$) reaction rates along the \\textsl{rp}-process path.  We use the multi-zone hydrodynamics code \\textsc{Kepler} which implicitly couples a full nuclear reaction network of more than 1000 isotopes, as needed, to follow structure and evolution of the X-ray burst layer and its ashes.  This allows us to incorporate the full \\textsl{rp}-process network, including all relevant nuclear reactions, and individually study changes in the X-ray burst light curves when modifying selected key nuclear reaction rates.  In this work we considered all possible proton captures to nuclei with $10<Z<28$ and $N\\le Z$.  When varying individual reaction rates within a symmetric full width uncertainty of a factor of $10^4$, early results for some rates show changes in the burst light curve as large as 10 percent of peak luminosity.  This is very large compared to the current sensitivity of X-ray observations.  More precise reaction rates are therefore needed to test current X-ray burst models, particularly of the burst rise, with observational data and to constrain astrophysical parameters. }  \\FullConference{International Symposium on Nuclear Astrophysics - Nuclei in the Cosmos - IX\\\\ 25-30 June 2006\\\\ CERN, Geneva, Switzerland}  \\begin{document} ",
        "introduction": "Type I X-ray bursts are explosive events on the surface of accreting neutron stars.  They occur for certain accretion rates at which the hydrogen and helium nuclear fuel are largely consumed through the rapid proton capture process (\\textsl{rp}-process) and the $\\alpha$p-process, respectively. The \\textsl{rp}-process takes place under conditions of extreme temperature and density in regions of proton excess and thus functions as a series of fast proton captures and weak decays involving isotopes on the proton-rich side of stability, including several near the proton drip-line.  The process is therefore sensitive to nuclear structure properties far from stability, mostly unknown to nuclear physics.  Multi-zone burst models allow the study of the evolution of the neutron star crust composition over time, from accretion through stable burning, through the burst, and following the burst.  This allows consideration of effects which are impossible to study using only single-zone models. Through compositional inertia, the ashes of earlier bursts affect the character of later ones.  Also, the reheating of burst ashes by successive bursts will further process the nuclear material.  The resulting abundances in deeper layers of the neutron star crust, specifically that of carbon, are relevant to superburst processes.  This work examines the sensitivity of Type I X-ray bursts to all possible proton captures (those which produce particle stable compound systems) on nuclei with $10<Z<28$ and $N\\le Z$, varying one rate at a time within an assumed uncertainty, with the goal of determining the extent to which our ability to model these systems is limited by our current knowledge of the nuclear physics inputs.  Similar studies conducted previously used postprocessing techniques~\\cite{hix} or analytical methods~\\cite{izz} in single-zone models, but because the \\textsl{rp}-process and $\\alpha$p-process to be studied are the main energy sources, such a treatment is inadequate here.  We used a specially adjusted, single-zone burst model to make a rough search for sensitivity in the burst light curve and nucleosynthesis in the full set of reactions.  These results were compared, in selected cases, to the results of more rigorous multi-zone simulations. ",
        "conclusions": "Based on the few reactions studied so far it seems that the one-zone approximation can help to identify critical rates in some cases.  While there can also be large differences in the sensitivities to rates between the two models, the one-zone approximation can be used as an initial guide to narrow down the list of potentially important rates to be studied with the multi-zone model.  This approach does not necessarily identify all rate sensitivities, but within the computing time limits it provides an effective way to determine some of them.  This observation further demonstrates the importance of considering effects which exist only in higher dimensions and will hopefully encourage work on codes that can consider Type I bursts in greater than one dimension.  In the future, the one-zone model analysis could be extended to the variation of multiple rates or to a full Monte Carlo analysis.  There is at least one case found in this preliminary analysis with good evidence of a strong dependence of the burst lightcurve on a single reaction rate.  The other uncertain reaction rates considered in the multi-zone code to date appear not to be such critical inputs to the process for the specific accretion model considered."
    },
    "0608/astro-ph0608517_arXiv.txt": {
        "abstract": "Detailed high resolution observations of the innermost central parts of nearby galaxies have revealed the presence of supermassive black holes\\cite{2000ApJ...539L..13G, 2001ApJ...547..140M, 2002ApJ...574..740T, 2005SSRv..116..523F}. These black holes may interact with their host galaxies by means of 'feedback' in the form of energy and material jets; this feedback affects the evolution of the host and gives rise to the observed relations between the black hole and the host \\cite{1998A&A...331L...1S}. Here we report observations of the ultraviolet emissions of massive early-type galaxies. We derive an empirical relation for a critical black-hole mass (as a function of velocity dispersion) above which the outflows from these black holes suppress star formation in their hosts by heating and expelling all available cold gas. Supermassive black holes are negligible in mass compared to their hosts but nevertheless seem to play a critical role in the star formation history of galaxies. ",
        "introduction": "The main strategy behind the modelling approach we follow is to first calculate the collapse and merging history of individual dark matter halos, which is governed purely by  gravitational interactions, and secondly to calculate the more complex physics of the baryons inside these dark matter halos, including  e.g. radiative cooling of the gas, star formation and feedback from supernovae by simplified prescriptions on top of the dark matter evolution. Each of the dark matter halos consist of three main components which are distributed among individual galaxies inside them. A stellar, cold and hot gas component, where the latter  is only attributed to {\\it central} galaxies,  the most massive galaxies inside individual halos. In the following sections, we will describe briefly the recipes used to calculate these different components which are mainly based on recipes presented in Kauffmann et al.\\cite{k99} (hereafter, K99) and Springel et al.\\cite{spr01} (hereafter, S01), and we refer readers for more details on model implementations to their work and references therein.  Throughout this paper we use the following set of cosmological parameters: $\\Omega_0=0.3$, $\\Omega_{\\Lambda}=0.7$, $\\Omega_b/\\Omega_0=0.15$, $\\sigma_8=0.9$ and $h=0.65$.  \\subsection{Dark Matter Evolution} We calculate the merging history of dark matter halos according to the prescription presented in Somerville et al.\\cite{som99a}. This approach has been shown to produce merging histories and progenitor distributions in reasonable agreement with result from N-body simulations of cold dark matter structure formation in a cosmological context\\cite{som00}. The merging history of dark matter halos is reconstructed by breaking each halo up into progenitors above a limiting minimum progenitor mass $M_{min}$. This mass cut needs to be chosen carefully as it ensures that the right galaxy population and merging histories are produced within the model. Progenitor halos with masses below $M_{min}$ are declared as {\\it accretion} events and their histories are not followed further back in time. Progenitors labelled as accretion events should ideally not host any significant galaxies in them and be composed mainly of primordial hot gas at the progenitor halo's  virial temperature. To achieve a good compromise between accuracy and computational time, we use $M_{min}=10^{10}$ $M_{\\odot}$, which ensures that the results presented here are unaffected by numerical resolution effects.  \\subsection{Baryonic Physics} As mentioned above, once the merging history of the dark matter component has been calculated, it is possible to follow the evolution of the baryonic content in these halos forward in time. We assume each halo  consists of three components: hot gas, cold gas and stars, where the latter two components can be distributed among individual galaxies inside a single  dark matter halo. The stellar components of each galaxy are further divided into bulge and disc, to allow morphological classifications of model galaxies. In the following, we describe how the evolution of each component is calculated.  \\subsection{Gas Cooling \\& Reionisation}\\label{cool} Each branch of the merger tree starts at a progenitor mass of $M_{min}$ and ends at a redshift of $z=0$. Initially, each halo is occupied by hot primordial gas which was captured in the potential well of the halo and shock-heated to its virial temperature $T_{vir}=35.9\\left[V_c/(\\mbox{km s}^{-1}) \\right]^2$ K, where $V_c$ is the circular velocity of the halo\\cite{k99,wf91}. Subsequently this hot gas component is allowed to radiatively cool and settles down into a rotationally supported gas disc at the centre of the halo, which we identify as the central galaxy\\cite{s77,wr78,wf91}. The rate at which hot gas cools down is estimated by calculating the cooling radius inside the halo using the cooling functions provided by \\cite{sd93} and the prescription in S01. In the case of a merger between halos we assume that all of the hot gas present in the progenitors gets shock-heated to the virial temperature of the remnant halo, and this gas can cool down only onto the new central galaxy which is the central galaxy of the most massive progenitor halo. The central galaxy of the less massive halo will become a satellite galaxy orbiting inside the remnant halo. In this way, a halo can host multiple satellite galaxies, depending on the merging history of the halo, but will always only host one central galaxy onto which gas can cool. The cold gas content in  satellite galaxies  is given by the amount present when they first became satellite galaxies and does not increase, but decreases due to ongoing star formation and supernova feedback.  In the simplified picture adopted above, the amount of gas available to cool down is limited only by the universal baryon fraction $\\Omega_b h^2=0.0224$\\cite{sper03}. However, in the presence of a photoionising background the fraction of baryons captured in halos is reduced \\cite{ns97,g00} and we use the recipe of Somerville et al.\\cite{som02}, which is based on a fitting formulae derived from hydrodynamical simulations\\cite{g00}, to estimate the amount of baryons in each halo. For the epoch of  reionisation, we assume $z_{reion}=20$, which is in agreement with observations of the temperature-polarisation correlation of the cosmic microwave background\\cite{ko03}.  \\subsection{Star Formation in Discs and Supernova Feedback}\\label{sf} Once cooled gas has settled down in a disc, we allow for fragmentation and subsequent star formation according to a parameterised global Schmidt-Kennicutt law\\cite{ken98} of the form $ \\dot{M}_{*}=\\alpha M_{cold}/t_{dyn,gal}$, where $\\alpha$ is a free parameter describing the efficiency of the conversion of cold gas into stars, and $t_{dyn,gal}$ is assumed to be the dynamical time of the galaxy and is approximated to be 0.1 times the dynamical time of the dark matter halo\\cite{k99}. As in K99 we allow star formation only in halos of $V_c < 350$ km/s to avoid excessively-massive central galaxies in clusters.  Feedback from supernovae plays an important role in regulating star formation in small mass halos and in preventing excessively-massive satellite galaxies from forming. We implement feedback based on the prescription presented in K99 with \\begin{equation} \\Delta M_{reheat}=\\frac{4}{3} \\epsilon \\frac{\\eta_{SN} E_{SN}}{V_{c}^{2}} \\Delta M_*. \\end{equation} Here we introduce a second free parameter $\\epsilon$ which represents our ignorance about the efficiency with which the energy from supernovae reheats the cold gas.  The expected number of supernovae per solar mass of stars formed is given by $\\eta_{SN}=5 \\times 10^{-3}$, taken as the value for the Scalo initial mass function\\cite{sca86}, and $E_{SN}=10^{51}$ erg  is the energy output from each supernova. We take $V_c$ as the circular velocity of the halo in which the galaxy was last present as  a central galaxy.  \\subsection{Galaxy Mergers} We allow for mergers between galaxies residing in a single halo. As mentioned earlier, each halo is occupied by one central galaxy and a number of satellite galaxies depending on the past merging history of the halo. Whenever two halos merge, the galaxies inside  them are going to merge on a time-scale which we calculate by estimating the time it would take the satellite to reach the centre of the halo under the effects of dynamical friction. Satellites are assumed to merge only with central galaxies and we set up their orbits in the halo according to the prescription of K99, modified to use the Coulomb logarithm approximation of S01.  If the mass ratio between the two merging galaxies is $M_{gal,1}/M_{gal,2} \\leq 3.5$ ($M_{gal,1} \\geq M_{gal,2}$) we declare the event as a {\\it major} merger and the remnant will be an elliptical galaxy and the stellar components and the gas will be treated according to the prescriptions below. In the case of {\\it minor} merger $M_{gal,1}/M_{gal,2} > 3.5$ the cold gas in the disc of the smaller progenitor is assumed to settle down in the gas disc of the remnant and its stars contribute to the bulge component of the remnant\\cite{k99}.  \\subsection{Formation of Ellipticals and Bulges} It is suggested that major mergers will lead to the formation of elliptical galaxies\\cite{tt72}. Indeed detailed numerical simulations in the last decade seem to support this hypothesis\\cite{ba92,bn03}, and we wil  assume in the following that major mergers disrupt the discs in progenitor galaxies as seen in various numerical simulations and relax to a spheroidal distribution. During the merger, any cold gas in the discs of the progenitor galaxies is assumed to be funnelled into the centre of the remnant where it ignites a starburst which transforms all of the cold gas into stars contributing to the spheroidal component\\cite{k99,spr01}. The second assumption is certainly a simplification of what might happen since we neglect the possibility  that not all of the cold gas is funnelled to the centre but some fraction of it may for example settle down onto an extended disk which continues to grow inside-out by fresh supply of gas from tidal tails\\cite{bh91,mih96,ba01,b02}. The results of Barnes\\cite{b02} indicate that $40\\% - 80\\%$ of the initial gas in the disc could end up in the central region of the remnant and be consumed in a starburst. The exact amount is somewhat dependent on the merger geometry and on the mass ratio of the merger. Unfortunately, a sufficiently large survey investigating the gas inflow to the centres of merger remnants is not available yet, so we use the simplified approach of assuming that all cold gas gets used up in the central starburst. This prescription for the fate of the cold gas  results in an overestimate of the spheroid masses and an underestimate of the secondary disc components in our model. This  is not very significant for massive ellipticals since they are mainly formed in relatively non-dissipative mergers\\cite{kb03}.  \\subsection{Formation and Growth of Super-Massive Black Holes} We here follow the model introduced in Kauffmann \\& Haehnelt\\cite{2002MNRAS.332..529K} in which super-massive black holes get formed and fed during major mergers of glaxies. The assumption is to say that a fraction of the the cold gas available in the progenitor discs, that is funnelled into the centre of the remnant, will be accreted onto the black hole. Kaufmann \\& Haehnelt introduced the following scaling law for the effectivness of this procces: \\begin{equation} \\dot{M_{\\bullet}}=\\frac{f_{bh} M_{cold}}{1+(280 \\mbox{km s}^{-1}/V_c)^2}, \\end{equation} where $M_{cold}$ is the amount of cold gas available in the progenitor disks, $V_c$ the circular velocity of the dark matter halo and $f_{bh}$ a free parameter. This scaling provides a good fit to the observed relation between the mass of the super-massive black hole and the velocity dispersion of a galaxy\\cite{2002MNRAS.332..529K}. Following Croton et al.\\cite{2006MNRAS.365...11C} we here assume the velocity dispersion of a galaxy to be identical to the circular velocity of the dark halo the galaxy is embedded in. This however is not strictly true as e.g. in the case of an isothermal sphere $V_c/\\sigma =\\sqrt{2}$. Several authors have investigated the correlation between $V_c$ and $\\sigma$\\cite{2002ApJ...578...90F,2005ApJ...631..785P} and found different correlations. We here note that using a different scaling between $V_c$ and $\\sigma$ in our models still produces the same $M_{\\bullet}-\\sigma-$relation if we adjust the free parameter $f_{bh}$ accordingly.  \\subsection{Feedback from Super-Massive-Black Holes} We here introduce a new prescription for the modelling of feedback from super-massive black holes, which is based on the results presented in this paper. According to our analyses a critical black hole mass exists for a given galaxy velocity dispersion $\\sigma$, at which feedback is so strong that it reheats all the available cold gas and hence prevents further star formation. This critical black hole mass is derived from our empirical fit to the residual star formation fraction found in our data as \\begin{equation} M_{\\bullet,c}=1.26 \\times 10^8 \\left(\\frac{\\sigma}{200} \\right)^{3.65} \\end{equation} In galaxies with black holes more massive than the corresponding critical black hole mass we stop any cooling of gas and star formation. We callibrate this relation to give the right fraction of RSF galaxies as seen in Figure 1.  \\subsection{Free Parameters} We normalise our two model parameters for the star formation efficiency $\\alpha$ and supernova feedback efficiency $\\eta$ by matching the $I-$band Tully-Fisher relation of Giovanelli et al.\\cite{gio97} and requiring that spiral central galaxies of halos with circular velocity $V_C=220$ km/s have on average $10^{11}$ M$_{\\odot}$ of stars and $10^9$ M$_{\\odot}$ of cold gas\\cite{som99b}. The third free parameter $f_{bh}$, which regulates the black hole growth is set to match the observed $M_{\\bullet}-\\sigma-$relation and is $f_{bh}=0.02$. ",
        "conclusions": ""
    },
    "0608/astro-ph0608047_arXiv.txt": {
        "abstract": "{Intercluster medium is expected to be turbulent with turbulence being superAlfvenic at large scales. Magnetic fields substantially modify the turbulent cascade when the turbulence reaches the scales at which the fluctuation velocity gets less than the Alfven velocity. At those scales it is possible to consider three cascades, of fast, slow and Alfven modes with little energy exchange between them. As Alfvenic and slow modes are anisotropic they marginally scatter and accelerate cosmic rays, while fast modes dominate the processes. However, in the presence of cosmic rays the turbulence is modified as cosmic rays transfer the energy of compressible motions (i.e. slow and fast modes) from large  scales to the scale of cosmic ray Larmor radius. This results in generation of a new small-scale  Alfvenic component which is not a part of the ordinary MHD cascade. This component does scatter and accelerate cosmic rays. In addition, magnetic reconnection in turbulent medium accelerates cosmic rays. The complexity of the interacluster turbulence calls for observational studies. A new technique Velocity Coordinate Spectrum (VCS) is particularly promising for studies of velocity fluctuations with a new generation of X-ray observatories. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608101_arXiv.txt": {
        "abstract": "%  Cosmic-ray photons above 10$^{19}$~eV can convert in the geomagnetic field and initiate a pre\\-shower, i.e.~a particle cascade before entering the atmosphere. We compare the preshower characteristics at the southern and northern sites of the Pierre Auger Observatory. In addition to a shift of the preshower patterns on the sky due to the different pointing of the local magnetic field vectors, the fact that the northern Auger site is closer to the geomagnetic pole results in a different energy dependence of the preshower effect: photon conversion can start at smaller energies, but large conversion probabilities ($>90\\%$) are reached for the whole sky at higher energies compared to the southern Auger site. We show how the complementary preshower features at the two sites can be used to search for ultra-high energy photons among cosmic rays. In particular, the different preshower characteristics at the northern Auger site may provide an elegant and unambiguous confirmation if a photon signal is detected at the southern site. ",
        "introduction": "\\label{sec-intro} Substantial fluxes of cosmic-ray photons at ultra-high energy (UHE), above 10$^{19}$~eV, are predicted by non-acceleration (top-down) models of cosmic-ray origin (for example, see~\\cite{bhat-sigl}). At a smaller level, UHE photons are also expected to be produced in acceleration (bottom-up) models~\\cite{models}. UHE photons can be detected through the particle cascades they initiate in the atmosphere. So far, upper limits on the photon flux were set~\\cite{ave,shinozaki,risse05,troitsky,augerphoton}. The large exposure expected to be collected during the next years, in particular by the Pierre Auger Observatory~\\cite{auger}, will enormously increase the sensitivity for detecting UHE photons~\\cite{augerphoton}.  Contrary to the case of hadron primaries, UHE photons around 10$^{20}$~eV can interact with the magnetic field of the Earth before entering the atmosphere~\\cite{mcbreen}. This process is commonly called geomagnetic cascading or preshower and leads to a dramatic change of the air shower development for primary photons~\\cite{mcbreen,collection,stanev97,bednarek99,bertou00,homola}.  The probability of magnetic $e^+e^-$ pair production (``photon conversion'') and, in case of conversion, the synchrotron emission by the produced electrons depend on the particle energy and on the transverse component of the local magnetic field~\\cite{erber,homola}. Photon conversion can occur at altitudes up to several thousand kilometers above the ground, i.e.~well above the atmosphere ($\\sim$100~km altitude). A preshower consists mostly of the secondary photons emitted by the electrons. A secondary photon of sufficiently high energy may additionally convert in the magnetic field and produce another electron pair. The number of preshower particles depends on the specific conditions and can fluctuate strongly; it can range from a few particles to a few thousand. Nevertheless, as we demonstrate in this paper, the preshower formation can be fairly well inferred from characteristics of the observed extensive air shower. The spread of the preshower particles in transverse distance from the shower axis and in arrival time at the atmosphere is well below the resolution of current shower experiments~\\cite{homola}. Hence, also in case of geomagnetic cascading, the subsequent air shower would be observed as one event by the experiments.  Since the probability of photon conversion and the synchrotron emission spectrum depend on the local transverse magnetic field, characteristics of the preshower process (such as conversion probability, or the energy spectrum of the preshower particles) depend strongly on the specific trajectory through the magnetosphere. For a given experimental site, this leads to a large directional dependence of preshower characteristics within the local coordinate system of zenith $\\theta$ and azimuth $\\phi$. In addition, as already noted in the work of McBreen and Lambert 25 years ago~\\cite{mcbreen}, a dependence is expected on the location of the experiment: experimental sites at different geomagnetic latitudes have different local magnetic field conditions, which can affect preshower characteristics for UHE photons.  The comparison of preshower characteristics of UHE photons expected at two sites located at different hemispheres of the Earth is of particular current interest due to the construction of the Pierre Auger Observatory. The southern part of the Observatory (``Auger South'') is situated near Malargue (Argentina) at 69.2$^\\circ$~W, 35.2$^\\circ$~S. The construction of Auger South is approaching completion and first results were obtained, including an upper limit to UHE photons~\\cite{augerphoton}. To achieve full sky coverage, the northern site (``Auger North'') is planned in Colorado (USA) at 102.7$^\\circ$~W, 37.7$^\\circ$~N. At Auger South, the magnetic field of $\\sim$24.6~$\\mu T$ points upward to $\\theta \\sim 55^\\circ ,\\phi \\sim 87^\\circ$. At Auger North, the magnetic field of $\\sim$52.5~$\\mu T$ points downward from $\\theta \\sim 25^\\circ ,\\phi \\sim 262^\\circ$.\\footnote{Azimuth is defined in this work counterclockwise from geographic East. For instance, $\\phi = 0^\\circ$ means East, $\\phi = 90^\\circ$ North etc.}  Previous works in the literature focused rather on the directional dependence of preshower characteristics of UHE photons for one given site than comparing the conditions at different sites. Examples of magnetic field strengths at various sites were given in~\\cite{stanev97}. In that work, the observable of shower size (total number of electrons at ground) was studied. It was noted that azimuth regions of smaller (or stronger) preshower effect can be identified that may differ between the sites. In~\\cite{bednarek99}, the altitude dependence of the transverse component of the magnetic field was compared for the two Auger sites for a few directions, and preshower particle spectra were studied. E.g.~in~\\cite{bertou00}, it was pointed out that at Auger North, due to the larger magnetic field, the preshower effect starts at a lower energy. A particular focus of that work was the study of observables for the ground array, where differences in the detector signals for converted and unconverted primary photons were found.  In this work, making use of our PRESHOWER code~\\cite{homola} linked to CONEX, a fast shower simulation program~\\cite{conex}, we perform a dedicated comparison of the preshower characteristics of UHE photons for the conditions of the Auger North and Auger South sites. The geomagnetic field in PRESHOWER is described by the IGRF model~\\cite{igrf,tsygan}. We find that the transition range in energy from small to large conversion probabilities is larger at Auger North, starting at smaller but ending at higher energy. We study how the differences of the preshower characteristics between the two sites can be used to perform a complementary search for UHE photons. In particular, the situation may arise where a possible UHE photon signal detected at Auger South requires an independent cross-check. As will be shown, such a cross-check could come from Auger North. In fact, making use of the UHE preshower effect by observing from a different experimental site seems the most elegant and natural way for cross-checking an UHE photon signal claim.  While the study is performed for the specific case of the two Auger sites, the general findings hold for any two sites with sufficiently different local magnetic field conditions. The results shown for Auger North hold, to a good extend, also for the HiRes~\\cite{hires} and Telescope Array~\\cite{telarray} sites, both of them being located in Utah.  The plan of the paper is as follows. In Section~\\ref{sec-preshower}, preshower features are discussed, and in Section~\\ref{sec-eas}, the corresponding air shower features are studied. In Section~\\ref{sec-scenarios}, UHE photon scenarios and their observation at the two different sites are investigated. In particular, it is worked out how observations at the two sites can complement each other. Conclusions are given in Section~\\ref{sec-conclusion}. In Appendix~\\ref{app-a} it is shown how the transverse component of the geomagnetic field in the IGRF model varies along different particle trajectories. In Appendix~\\ref{app-b} we show a collection of conversion probability sky maps supplementing Section \\ref{subsec-convprob}. ",
        "conclusions": "\\label{sec-conclusion} The difference of conversion probabilities is a convenient and effective parameter to estimate the complementarity between two sites for searching for UHE photons. Regarding the two sites of the Pierre Auger Observatory, Auger North and Auger South, significant differences in the preshower features of UHE photons exist. The sky patterns of the preshower features are shifted due to the different pointing directions of the local magnetic field vectors. In addition, Auger North is located closer to the magnetic pole. This leads to a larger transition range in energy from small ($<10\\%$) to large ($>90\\%$) conversion probabilities compared to Auger South. More specifically, photon conversion starts at Auger North at smaller energies due to the larger magnetic field, but conversion probabilities $>90\\%$ are reached for the whole sky at higher energies due to the smaller bending of the field line with distance.  These differences in the preshower characteristics result in different rates of (un-)converted photons from the same (both in local and astronomical coordinates) regions of the sky. Air showers initiated by converted and unconverted photons can be well distinguished by current experiments. The main difference is related to the position of depth of shower maximum $X_{\\rm max}$, which is typically $\\sim$200$-$300~g~cm$^{-2}$ smaller for converted photons.  For a variety of UHE photon flux scenarios (diffuse photon flux; photons from source regions; absence of photons), the different preshower characteristics at the experimental sites can be used for a complementary search for UHE photons. Most important, a possible detection of UHE photons at Auger South may be confirmed in an unambiguous way at Auger North by observing the well predictable change in the signal from UHE photon showers.    {\\it Acknowledgements:} This work was partially supported by the Polish State Committee for Scientific Research under grants No.~PBZ~KBN~054/P03/2001 and 2P03B~11024, by Polish Ministry of Science and Higher Education under grant N202 090 31/0623 and in Germany by the DAAD under grant No.~PPP~323. Helpful remarks of the unknown referee, especially those concerning the motivation of Subsection~\\ref{subsec-localized-source}, are kindly acknowledged. MR acknowledges support from the Alexander von Humboldt foundation.    \\appendix"
    },
    "0608/hep-th0608193_arXiv.txt": {
        "abstract": "{ In this work, we have continued the study of the Hawking radiation on the brane from a higher-dimensional rotating black hole by investigating the emission of fermionic modes. A comprehensive analysis is performed that leads to the particle, power and angular momentum emission rates, and sheds light on their dependence on fundamental parameters of the theory, such as the spacetime dimension and angular momentum of the black hole. In addition, the angular distribution of the emitted modes, in terms of the number of particles and energy, is thoroughly studied. Our results are valid for arbitrary values of the energy of the emitted particles, dimension of spacetime and angular momentum of the black hole, and complement previous results on the emission of brane-localised scalars and gauge bosons.} ",
        "introduction": "The formulation, a few years ago, of theories \\cite{ADD, RS} postulating the existence of additional spacelike dimensions in nature have radically changed the landscape in gravitational physics (for some early works towards this direction, see \\cite{early}). Driven by the motivation to address the hierarchy problem, these theories leave unchanged the usual 4-dimensional particle physics by localising the Standard Model (SM) particles on a 4-dimensional brane while allowing the gravitons to propagate in the bulk -- the whole of the $(4+n)$-dimensional spacetime. In the case of the scenario of Large Extra Dimensions \\cite{ADD}, in which the additional spacelike dimensions are compact with size $L$, the latter feature introduces a new fundamental higher-dimensional scale for gravity $M_*$, that can be significantly lower than the 4-dimensional scale $M_P$, thus increasing the fundamental Newton's constant -- and  the strength of gravitational interactions -- by orders of magnitude.  An immediate consequence of the above assumption is that particles with energies larger than $M_*$ can probe physics beyond the scale of quantum gravity while their collisions may trigger the appearance of strong gravitational phenomena and lead to the creation of extended objects ($p$-branes, string balls, string states, etc), rather than ordinary particles, including tiny black holes \\cite{creation}. This type of trans-planckian particle collision could be achieved in various environments, such as at ground-based colliders \\cite{colliders} or in high energy cosmic-ray interactions in the atmosphere of the Earth \\cite{cosmic} (for an extensive discussion of the phenomenological implications and additional references, see the reviews \\cite{Kanti, reviews, Harris}). The black holes produced with horizon radius $r_h \\ll L$ are submerged into the higher-dimensional spacetime, which modifies their properties (horizon radius, temperature, entropy, life-time etc.) \\cite{ADMR}. Nevertheless, they are still expected to go through the same stages in their life as their 4-dimensional counterparts: a short {\\it balding} phase, during which the black hole will shed all additional quantum numbers apart from its mass, angular momentum and charge, the more familiar {\\it spin-down} phase, during which the black hole will lose mainly its angular momentum, then the {\\it Schwarzschild} phase, when the black hole gradually loses its actual mass, and finally the {\\it Planck} phase, where the black hole has reduced to a quantum object with unknown properties.  During the two intermediate phases, the spin-down and Schwarzschild, the black hole loses energy through the emission of Hawking radiation \\cite{hawking} (and through superradiance during the former phase). This will take place via the emission of elementary particles both in the bulk and on the brane, and it will be characterised by a very distinct thermal spectrum. It is this emission spectrum that upon successful detection will provide the most convincing evidence not only for the creation of microscopic black holes but of the very existence of extra spacelike dimensions. For this reason, the emission spectra from a higher-dimensional black hole have been under intensive study over the last few years. Having the simplest gravitational line-element, the Schwarzschild phase was the first one to undergo a systematic study both analytically \\cite{kmr1, Frolov1, kmr2} and numerically \\cite{HK1}. Both bulk (scalars) and brane (scalars, fermions and gauge bosons) channels were studied, and the derived emission spectra were shown to depend strongly on the number of additional spacelike dimensions in nature, both in terms of the amount of energy and number of particles produced per unit time and of the type of particles produced. More general spherically-symmetric backgrounds were also studied, such as Schwarzschild-de Sitter \\cite{BGK}, Schwarzschild-Gauss-Bonnet \\cite{Barrau} and Reissner-Nordstrom \\cite{Jung-charge} backgrounds, as well as the effect of the mass of the emitted particles on the spectrum \\cite{mass}. Recently, the graviton emission rates by a higher-dimensional Schwarzschild black hole were also investigated by a number of authors \\cite{gravitons}.  After the spherically-symmetric Schwarzschild phase was almost exhaustively studied, the investigation of the form of the radiation spectra for a higher-dimensional rotating black hole was initiated. Two early analytical studies \\cite{Frolov2,IOP1} were supplemented, during the last year or so, by a number of works attempting to shed light on this more technically involved case. Exact numerical results for the emission of Hawking radiation from a $(4+n)$-dimensional rotating black hole, in the form of scalar fields on the brane, were first presented in \\cite{HK2}. In that work, the exact value of the angular eigenvalue, that couples the angular and radial part of the equation of motion of the scalar field and does not exist in closed form, was found numerically, and the radial equation was also solved by employing numerical analysis. A comprehensive study of the fluxes of particles, energy and angular momentum for the emission of scalar fields on the brane soon followed \\cite{DHKW}, that once again presented exact results for the emission rates including the angular distribution of particles. A similar extensive analysis, focussed on the emission of gauge bosons on the brane by a rotating black hole, appeared in \\cite{CKW}. In both of the latter works, it was shown that all emission rates for scalars and gauge fields are greatly enhanced by both the dimension of spacetime and the angular momentum of the black hole. Results produced independently in \\cite{IOP2, IOP3}, for these two species of particles, reveal a good agreement between the two sets of analyses. Additional studies in the literature featured the study of the superradiance effect for a higher-dimensional black hole \\cite{super} and of the emission in the bulk \\cite{rot-bulk}.  Recently, and while this work was in progress, a study \\cite{IOP3} produced power and angular momentum emission spectra from a rotating black hole for the last species of brane-localised particles, i.e. fermions. The radial part of the spinor equation was solved numerically with the angular eigenvalue given by an analytical expansion (with terms kept up to 6th order). The power and angular momentum emission spectra were found to exhibit the same characteristics as the ones for scalars and gauge bosons. The present work aims to complete the above analysis by presenting all emission spectra (particle, energy and angular momentum) as well as total emissivities (emission rates integrated over the frequency) and angular distribution of the emitted radiation. To this end, both the radial and angular part of the fermionic equation of motion will be solved numerically, and an exact numerical value for the angular eigenvalue will also be determined and employed in the calculation.   The outline of our paper is as follows: we start, in Section 2, with the presentation of the theoretical framework for our analysis, and the derivation of the radial and angular part of the spin-1/2 equation of motion on the brane. The expressions for the transmission coefficients and the various Hawking radiation emission rates, adopted for emission on the brane, are reviewed in Section 3. In Section 4, the numerical techniques used to integrate the angular and radial parts of the spin-1/2 equation of motion and determine the angular eigenvalue are described. Section 5 presents our numerical results for the transmission coefficient, the particle, energy and angular momentum fluxes, as well as the total emissivities, and their dependence on both the dimension of spacetime and angular momentum of the black hole is investigated. The angular distribution of the particle and energy emission rates are also studied by using the exact values of the spin-1/2-weighted spheroidal harmonics. We finish with a summary of our results and conclusions, in Section 6. ",
        "conclusions": "In this work, we have completed the study of the Hawking radiation on the brane from a higher-dimensional rotating black hole by investigating the emission of fermionic modes. A comprehensive analysis was performed that led to the particle, power and angular momentum emission rates, and shed light on their dependence on fundamental parameters of the theory such as the spacetime dimension and angular momentum of the black hole. In addition, the angular distribution of the emitted modes, in terms of number of particles and energy, was thoroughly studied. Our results are valid for arbitrary values of the energy of the emitted particles\\footnote{Using the results in \\cite{NS}, we have investigated at what value of $\\omega r_h$ is the contribution to the spectrum from ultra-short distances (where a full theory of quantum gravity would be required) of the same order as the total spectrum itself. For the main cases plotted in this paper, such value of $\\omega r_h$ is at the tail-end of the curve for each dominant mode, so that the main contribution to the spectrum is indeed from the semiclassical regime.} and angular momentum of the black hole, and complement previous results on the emission of brane-localised scalars \\cite{HK2,DHKW} and gauge bosons \\cite{CKW}.  As in the case of our previous analyses for spin-0 and spin-1 particles, the equation of motion for fermions propagating in the projected-on-the-brane gravitational background needed to be determined first, a task that was performed by employing the Newman-Penrose formalism. The derivation of exact results for the various emission rates demanded the numerical integration of both the angular and radial part of the equation of motion in exactly that order. The former calculation offered exact results for both the angular eigenfunctions, the spin-weighted spheroidal harmonics, and the angular eigenvalues, that also appear in the radial equation. The latter calculation led to exact results for the radial part of the wavefunction of the fermionic field, and eventually for the transmission coefficient.  Our numerical results for the transmission probability as well as for the particle, energy and angular momentum fluxes were presented in Section \\ref{num-results}. The differential emission rates per unit time and frequency, integrated over the azimuthal angle $\\theta$, were addressed first, and spectra for different values of the dimension of spacetime and the angular momentum of the black hole were constructed as a function of the energy of the emitted particle. As in the case of scalars and gauge bosons, it was revealed that any increase in either $n$ or $a_*$ significantly enhances all emission rates of the black hole. The enhancement is present in all frequency regimes, and significantly boosts the emission of higher energy particles compared to the case of a non-rotating black hole. Total emissivities, i.e. emission rates integrated over all frequency regimes, were also presented at the end of Section \\ref{num-results}; these total emission rates confirmed the enhancement of all spectra, in terms of $n$ and $a_*$, and led to the derivation of the corresponding enhancement factors: as $n$ increases, the various emissivities are enhanced by a factor that varies from 35 to almost 100, while, as $a_*$ grows, this factor is of the order of 3 or 6.  Throughout the study of the emission of various spin-$s$ particles on the brane by a higher-dimensional rotating black hole, a significant amount of energy was invested in the calculation of the exact spin-weighted spheroidal harmonics, that would allow us to determine the exact angular distribution of the emitted particles and energy. Contrary to the case of a spherically-symmetric black hole, the various emission rates for a rotating black hole bear a signature of a characteristic direction in space which is the axis of rotation. It was found that the majority of particles emitted by the black hole are restricted to do so along the equatorial plane due to the strong centrifugal force - this conclusion was reinforced after the total energy emissivities as a function of the azimuthal angle were also computed. Nevertheless, the spin of the particle can significantly affect the angular distribution pattern at the low-energy regime due to the spin-rotation coupling with the angular momentum of the black hole. For comparison, in Figs. \\ref{comparison1}(a,b,c) we present the angular distribution of the energy emission spectrum for a 6-dimensional ($n=2$) rotating black hole emitting radiation on the brane, in the form of scalars \\cite{DHKW}, fermions and gauge bosons \\cite{CKW}, respectively. From these, it becomes obvious that the angular variation in the pattern of the emitted radiation can be a distinctive signature of emission from a rotating black hole, but also of the spin and energy of the particles emitted.   \\FIGURE{ \\begin{tabular}{c} \\hspace*{-0.4cm} \\epsfig{file=pow-angs0n2a1.eps, height=4cm}\\hspace*{-1.25cm} \\epsfig{file=pow-angs05n2a1.eps, height=4cm}\\hspace*{-1cm}{ \\epsfig{file=pow-angs1n2a1.eps, height=4cm}}\\end{tabular} \\caption{Angular distribution of the power spectra for {\\bf (a)} scalars, {\\bf (b)} fermions, and {\\bf (c)} gauge bosons on the brane from a 6-dimensional black hole with $a_*=1$.\\label{comparison1}} }  The array of graphs displayed in Fig. \\ref{comparison1} can shed light also on the question of the relative emissivities of different types of particles on the brane from a rotating black hole. In the absence of angular momentum, it was shown \\cite{HK1} that the preference of the black hole for the emission of scalar fields in four dimensions is replaced by the one for gauge bosons when a number of additional spacelike dimensions is introduced. In \\cite{CKW}, we were able to show that a higher-dimensional rotating black hole also prefers to emit gauge bosons over scalars on the brane. With the results for fermions being now available, the complete picture has emerged. By merely comparing the vertical axes of the graphs in Fig. \\ref{comparison1}, one may see that a higher-dimensional rotating black hole emits approximately the same amount of energy on the brane in the form of scalars and fermions -- a more careful comparison of the emission rates integrated over the azimuthal angle, presented in Fig. \\ref{n1power}, with the corresponding ones in \\cite{DHKW} reveals that the fermions slightly dominate over the scalars. Nevertheless, as Fig. \\ref{comparison1}(c) reveals, a significantly larger amount of energy (by, at least, an order of magnitude) is spent by the black hole in the emission of gauge bosons -- the above result holds for all values of $n$.  The results presented in this work, addressing modifications of the properties of black holes submerged into a higher-dimensional brane-world spacetime, have an obvious theoretical interest. The potential detection of the Hawking radiation from a higher-dimensional black hole formed during a high-energy particle collision in a ground-based accelerator, in the near or far future, increases their significance as the corresponding spectra offer vital information about fundamental parameters of the system under study, such as the dimension of spacetime and angular velocity of the black hole, as well as for the creation of the black hole itself. This information is obtained by looking only at the emission channel along our brane where ordinary Standard Model particles, as well as the potential observers, are restricted to live. With the investigation of the spin-down and Schwarzschild phases of the life of the black hole now completed, what remains to be done is a realistic study of the dynamical evolution of the produced black hole where more advanced aspects, such as the variation of the temperature of the black hole as the decay progresses or the properties and interactions of the Standard Model particles, are properly taken into account. This goal will be achieved by entering our results for the emission of individual scalar, fermionic and gauge bosonic degrees of freedom from a constant-mass black hole into the CHARYBDIS black hole event generator \\cite{HRW}. We hope to report results from this analysis soon."
    },
    "0608/astro-ph0608115_arXiv.txt": {
        "abstract": "We describe a high-precision, deep (to $V \\sim 19-21$) absolute proper-motion survey that samples $\\sim 50$ lines of sight in the Kapteyn Selected Areas along declination zones $-15\\arcdeg$,  $0\\arcdeg$ and $15\\arcdeg$.  In many fields the astrometric baseline reaches nearly a century.  We demonstrate that these data provide typical per star precisions between $\\sim1$ and 3 mas yr$^{-1}$ to the above magnitude limits, with the absolute reference frame established by numerous extragalactic sources in each survey field.  Combined with existing and ongoing photometric and radial velocity surveys in these fields, these astrometric data will enable, among other things, accurate, detailed dynamical modeling of satellite interactions with our Galaxy. In this contribution we describe the astrometric part of our survey and show preliminary results along the trailing tail of the Sagittarius dwarf galaxy, and in the Monoceros ring region. ",
        "introduction": "Modern, deep, uniform, large-area photometric surveys have shown unambiguously that the Milky Way outer halo contains accretion-derived substructure (e.g., Ibata et al. 2001, Newberg et al. 2002, Majewski et al. 2003, Rocha-Pinto et al. 2003, 2004, Vivas et al. 2004, Conn et al. 2005a). These structures of known or assumed remnants of satellite accretion, have long-lived, coherent tidal features that can be used to model the Galactic gravitational potential (e.g., Ibata et al. 2001, Law et al. 2005), as well as the characteristics of the original satellite (e.g., Johnston et al. 1999). However, such modeling studies have been limited by the meager available kinematical data over large angles along the tidal features. While radial-velocity programs have just recently begun to address this problem for the few known Galactic tidal tails (Sagittarius - Majewski et al. 2004, Vivas et al. 2005, and the Monoceros ring - Crane et al. 2003), no systematic survey has begun to address the transverse (tangential) velocities (i.e., absolute proper motions). Without this information, dynamical models remain poorly constrained, and therefore limited to describing merely a range of possible events, rather than an accurate description of the real event.  The work described here is a proper-motion survey that provides high precision (1 to 3 mas/yr per well-measured star) absolute and relative proper motions down to a magnitude of $V \\sim 19$, and for a few selected fields down to $V \\sim 21$, in $\\sim 50$ lines of sight in the Selected Areas (SA) designed by Kapteyn for Galactic structure studies in 1906. Current proper-motion surveys do not achieve this precision at a similar magnitude limit and thus have limited capability to detect and characterize distant halo substructure. To date there are only a few similarly deep, precise, pencil-beam-type proper-motion data sets that are primarily centered on globular clusters or dwarf spheroidals, and even fewer that are focused on Galactic field stars (e.g., Chiu 1980, Majewski 1992, Guo et al. 1993, Dinescu et al. 2002). Obviously, ground-based studies of this type are limited to parts of the sky where suitable first epoch astrometric data exist.  Here we exploit a unique cache of nearly century old, deep photographic plates having good scale, as well as other plate material collected in Kapteyn fields.  We intend to complement the new proper motion data with radial velocities, distances and metallicity estimates from our own photometric and spectroscopic work as well as from overlapping surveys such as 2MASS, SDSS, QUEST and RAVE.  With this survey we aim to (1) determine the extent and orbital motion of the highly obscured Monoceros ring-like structure above and below the Galactic plane, and explore its possible relation to other low latitude structures, such as the Canis Major (CMa) overdensity, (2) characterize the transverse motion of the Sagittarius tidal streams, (3) search for and characterize additional substructures in the halo of the Milky Way and (4) determine the kinematical properties of numerous thick disk and halo stars in our fields as a function of Galactic position.  Future papers will include detailed kinematical analyses for specific regions, while in this paper we characterize the survey and show its potential in two specific regions where tidal streams have already been identified from other sources. In the following Section we will describe the survey in detail. In Section 3 we show results in a few SA fields, and a brief summary is presented in Section 4. ",
        "conclusions": "We describe our ground-based, photographic absolute proper-motion survey along three near-equatorial declination zones. Proper motions are derived in $40\\arcmin\\times40\\arcmin$ fields from a collection of 2.5-m Du Pont, 4-m Mayall, 60-inch Mt. Wilson and POSS-I plates. The time baseline varies between 40 and 85 years. We have demonstrated that we obtain proper-motion uncertainties between $\\sim 1$ and 3 mas/yr per star down to a magnitude of 19, and for a few fields down to 21. The typical uncertainty in the correction to absolute proper motions as given by galaxies and QSOs is between 0.2 and 0.8 mas/yr. The described proper-motion survey is complemented by our own ongoing radial-velocity and photometric follow-up programs as well as by current surveys such as SDSS, QUEST, etc. The characteristics of this proper-motion survey make it suitable to address many topics related to both the main Galactic components, and the tidal features seen in the halo and associated with streams from disrupted satellites. In this contribution, we present preliminary results in the southern trailing tidal tail of Sgr and in the Monoceros ring region.     We thank Stephen Levine from USNO-Flagstaff for making available the POSS-I USNO scans. SRM is grateful to the Carnegie Observatories and its former director Augustus Oemler for a Carnegie Visiting Associateship that made possible the collection of the Du Pont plates used in this survey. The financial support from NSF grants AST-0406884 and AST-0407207 for this research is acknowledged.  This publication makes uses of SDSS data products. Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. The SDSS Web Site is http://www.sdss.org/."
    },
    "0608/astro-ph0608323_arXiv.txt": {
        "abstract": "We have converted the Palomar 60-inch telescope (P60) from a classical night assistant-operated telescope to a fully robotic facility. The automated system, which has been operational since September 2004,  is designed for moderately fast ($t \\lesssim 3$ minutes) and sustained ($R \\lesssim 23$ mag) observations of gamma-ray burst afterglows and other transient events.  Routine queue-scheduled observations can be interrupted in response to electronic notification of transient events.   An automated pipeline reduces data in real-time, which is then stored on a searchable web-based archive for ease of distribution.  We describe here the design requirements, hardware and software upgrades, and lessons learned from roboticization.  We present an overview of the current system performance as well as plans for future upgrades. ",
        "introduction": "\\label{sec:intro} The field of optical transient astronomy has matured to produce numerous important scientific discoveries in recent years. Type Ia supernovae (SNe) have been used as standard candles to produce Hubble diagrams out to $z \\sim 0.5$, providing evidence that the expansion of the universe is accelerating \\citep{rfc+98,pag+99}. Observations of the broadband afterglows of long-duration ($t > 2$ s) gamma-ray bursts (GRBs) have revealed an association with the deaths of super-massive stars \\citep{gvp+98,smg+03,hsm+03}.  The discovery of the first afterglows and host galaxies of short-duration ($t < 2$ s) GRBs \\citep{gso+05,bpp+06,hwf+05,ffp+05} has possibly revealed a new class of GRB progenitors: compact binary coalescence \\citep{elp+89}.  As interest in the field has steadily grown, new, more powerful methods of identifying optical transients have been developed. The \\Swift\\ Gamma-ray Burst Explorer \\citep{gcg+04} is currently providing $\\sim$ 100 prompt GRB localizations per year, an order-of-magnitude improvement over previous missions.  Planned wide-angle, high-cadence surveys with large facilities, such as Pan-STARRS \\citep{kab+02} and LSST \\citep{t05}, promise to overwhelm our current follow-up capability, providing hundreds of variable optical sources each night.  Dedicated, robotic, medium aperture ($1 - 3$ m) telescopes have the opportunity over the next few years to play a crucial role in this field.  Like small-aperture ($< 0.5$-m), robotic facilities, they can respond autonomously to transient alerts, providing observations at early times.  And given the relative abundance of such telescopes, it is entirely feasible to focus predominantly on transient astronomy.  However, like larger telescope ($> 5$-m), interesting events can be followed for longer durations and in multiple colors. In this sense robotic, medium-aperture facilities can act to bridge the gap between the earliest rapid-response observations and deep, late-time imaging and spectroscopy.  \\begin{figure}[p] \\plotone{f1_short.ps} \\caption[Early Afterglows of pre-\\Swift\\ GRBs and P60 Response Capabilities] {Early Afterglows of pre-\\Swift\\ GRBs and P60 Response Capabilities. Regions with a white background are accessible for automated P60 observations: $t \\gtrsim 3$ minutes, $R \\lesssim 23$ mag. With only a handful of examples, the early optical afterglows of pre-\\Swift\\ GRBs show a marked diversity.  GRB\\,990123 \\citep{abb+99} and GRB\\,021211 \\citep{fps+03,lfc+03} exhibit the fast $t^{-2}$ early-time decay indicative of adiabatic evolution of the reverse shock. On the other hand, GRB\\,021004 \\citep{fyk+03,hwf+03,psr+03} shows a distinctive slow $t^{-0.4}$ decay that likely signifies continuing energy input to shock regions. Reverse shock emission from GRB\\,030418 \\citep{rsp+04} was not seen; the optical peak at $t=0.4$ hours is due to the forward shock component. As a proof of concept, the P60 was the first to report the afterglow of GRB\\,040924 \\citep{fm04,lfc+04,hlh+04,sab+04, kay+04}.  The early time behavior is quite similar to that of GRB\\,021211.} \\label{fig:earlylc} \\end{figure}  To this end, we have roboticized the Palomar 60-inch telescope (P60). As a dedicated, robotic facility, the P60 is capable of responding moderately fast ($t \\lesssim 3$ min) to transient alerts.  With the increased event rate of \\Swift, the P60 is providing observations of the poorly understood early afterglow phase (Fig. \\ref{fig:earlylc}). Additionally, as a 1.5-m telescope, the P60 can continue the sequence of observations longer than most robotic telescopes. As Figure \\ref{fig:latelc} shows, one day after the burst, most afterglows have faded below $R = 20$; however, for days or even weeks after that, they remain at levels of $R < 23$ accessible to P60 photometry.  In this work, we first outline the high-level design requirements of a robotic system optimized for observations of transient sources (\\S\\ref{sec:design}). \\S\\ref{sec:procedure} provides the details of the automation procedure, including both the hardware and the software efforts. \\S \\ref{sec:performance} describes the current system performance (as of May 2006), which will primarily be of use for those interested in observing with the P60. Finally, in \\S\\ref{sec:conclusions}, we conclude with a summary of the project status and a discussion of possible future improvements to the robotic system.  \\begin{figure}[p] \\plotone{f2_short.ps} \\caption[Late-time light curves of pre-\\Swift\\ GRB afterglows] {Late-time light curves of pre-\\Swift\\ GRB afterglows. The gray shaded region displays the phase space inaccessible to automated P60 observations.  Observations of most afterglows require $> 1$-m class facilities after the first night; investigation of optically-extinguished (\"dark\") or high-redshift bursts require such facilities merely to register detections or collect physically interesting upper limits.} \\label{fig:latelc} \\end{figure} ",
        "conclusions": "\\label{sec:conclusions}  In this paper we have presented our efforts to automate and roboticize the Palomar 60-inch telescope.  As of September 2004, all components of the system operate in a fully automated fashion, making P60 one of the few robotic, medium-aperture facilities in the world.  The P60 has been routinely responding to \\Swift\\ GRB alters over the last year and a half, and will continue to do so over the lifetime of the \\Swift\\ mission. The system is well-positioned for the plethora of optical transients that will be discovered in the upcoming years.  In addition to the current optical camera, we are planning several major upgrades to further improve the scientific capabilities of the system. In the near-term, our top priority is to add a near-infrared (NIR) camera to P60.  We have already acquired the NIR detector from the out-of-use Cerro Tololo Infrared Imager (CIRIM\\footnote{http://www.ctio.noao.edu/instruments/ir\\_instruments/cirim/cirim.html}) and upgraded the controller electronics. We are currently working on both optical design and software development, with the hope of having both cameras mounted and functional in the next year. We also plan to make the P60 fully compliant with the Virtual Observatory Event Network (VOEventNet\\footnote{http://voevent.net}) protocol. In this manner the system can communicate with other observatories around the world without any needed human intervention.  As longer term projects, we are exploring the possibility of adding either a polarimeter or a multi-band camera to the facility.  Regardless of the details, we are committed to making the P60 a scientifically productive facility in the years to come."
    },
    "0608/astro-ph0608609_arXiv.txt": {
        "abstract": "Spectra have been obtained with the Infrared Spectrograph (IRS) on the Spitzer Space Telescope for 20 sources in the Lockman Hole field of the SWIRE survey. The sample is divided between sources with indicators of an obscured AGN, based primarily on X-ray detections of optically-faint sources, and sources with indicators of a starburst, based on optical and near-infrared spectral energy distributions (SEDs) which show a luminosity peak from stellar photospheric emission.  Ten of the 11 AGN sources have IRS spectra which show silicate absorption or are power laws; only one AGN source shows PAH emission features.  All 9 of the sources showing starburst SEDs in the near-infrared show PAH emission features in the IRS spectra.  Redshifts are determined from the IRS spectra for all 9 starbursts (1.0 $<$ z $<$ 1.9) and 8 of the 11 AGN (0.6 $<$ z $<$ 2.5).  Classification as AGN because of an X-ray detection, the classification as AGN or starburst derived from the photometric SED, and the IRS spectroscopic classification as AGN (silicate absorption) or starburst (PAH emission) are all consistent in 18 of 20 sources.  The surface density for starbursts which are most luminous in the mid-infrared is less than that for the most luminous AGN within the redshift interval 1.7 $\\la$ z $\\la$ 1.9. This result implies that mid-infrared source counts at high redshift are dominated by AGN for f$_{\\nu}$ (24$\\mu$m) $\\ga$ 1.0\\,mJy. ",
        "introduction": "Imaging surveys with the Spitzer Space Telescope ($Spitzer$) using the Multiband Imaging Photometer (MIPS) \\citep{rie04} detect sufficient sources at 24 $\\mu$m to resolve the cosmic background \\citep{hau98} at mid-infrared wavelengths \\citep{pap04, dol06}.  The source counts have been modeled to explain the infrared background as arising from the evolution of luminous, star-forming galaxies (starbursts) whose luminosity is primarily in the infrared and which arises from emission by dust \\citep{lag04, chr04}. It is also known from various infrared and X-ray studies that there are many obscured Active Galactic Nuclei (AGN) \\citep{ale01}, so many of the Spitzer mid-infrared sources could be dusty AGN instead of starbursts. It is essential to understand how the mid-infrared sources divide between starbursts and AGN so that the cosmic star formation history and the evolution of AGN can be distinguished and described as a function of redshift. This is a necessary step in understanding the relationships among the formation and evolution of stars, galaxies, and the massive black holes powering AGN within dusty environments.  Extensive spectroscopic studies with the Infrared Space Observatory (ISO) indicated that the strength of emission features from polycyclic aromatic hydrocarbons (PAH emission) is the primary mid-infrared diagnostic of starbursts \\citep{gen98, lut98, lau00}.  Relative to the underlying continuum, the PAH features from starbursts are at least 10 times stronger than from AGN, and this diagnostic leads to the conclusion that most ultraluminous infrared galaxies (ULIRGS) are powered by starbursts \\citep{rig00}.  If the $Spitzer$ surveys at 24 $\\mu$m are dominated by faint and distant ULIRGS which are starbursts, this conclusion implies that strong PAH features should frequently be seen in these sources.  So far, this expectation has not been confirmed with $Spitzer$ spectroscopic observations.  To begin understanding the faint 24 $\\mu$m source population, spectra have previously been obtained with the Infrared Spectrograph on $Spitzer$ (IRS) \\citep{hou04} of faint sources discovered with MIPS surveys.  Some of these sources are invisible optically, to limits fainter than $\\sim$ 26 mag in $Bw$, $R$, and $I$, so this population of highly obscured sources had not been discovered before $Spitzer$. Initial results of this spectroscopy are in \\citet{hou05}, \\citet{yan05}, \\citet{lut05}, \\citet{wee06a}, and \\citet{wee06b}. Redshifts were determined from mid-infrared spectral features for 43 of 58 sources observed in the Bo\\\"{o}tes field \\citep{hou05,wee06b}, for 6 of 8 sources selected by infrared color criteria in the $Spitzer$ First Look Survey (FLS) field \\citep{yan05}, for 14 of 18 radio sources in the FLS \\citep{wee06a}, and for 2 of 2 sources selected as submillimeter sources by \\citet{lut05}.  In summary, of the 86 mid-infrared spectra of faint $Spitzer$ 24\\,\\um sources discussed to date (f$_{\\nu}$(24\\,\\um) $\\sim$ 1 mJy), redshifts were measured in the IRS spectra for 65 sources. The remaining 21 sources have featureless continua, to within the signal-to-noise ratio (S/N) of the IRS observations.  The most important results so far are that the sources are generally at high redshift and have infrared spectra dominated by a strong silicate absorption feature at rest frame 9.7\\,\\um rather than by PAH emission.  For the 65 sources with IRS spectroscopic redshifts, the redshift (z) ranges from 1.1 to 2.8, with a median z = 2.1.  59 of these 65 sources have strong silicate absorption, with only 6 having redshifts derived from PAH emission.  These results are dominated by sources from the Bo\\\"{o}tes field \\citep{hou05,wee06b}, which were all chosen only on the basis of extreme infrared to optical flux ratios.  (Defining IR/opt = $\\nu$f$_{\\nu}$(24\\,\\um)/$\\nu$f$_{\\nu}$($R$), the lower limit for this parameter in the Bo\\\"{o}tes  sources is 40.)  However, the optically-faint sources selected because they obey the radio-infrared correlation expected for starbursts also show IRS spectra characterised by silicate absorption \\citep{wee06a}.  A few sources selected using the color criterion f$_{\\nu}$(24\\,\\um)/f$_{\\nu}$(8\\,\\um) $>$ 10 show PAH emission at z $\\sim$ 2 \\citep{yan05}, but only sources selected using submillimeter detections \\citep{lut05} consistently show PAH emission characteristic of starbursts.  From existing observations, therefore, the optically faint, infrared-selected population appears to be primarily an AGN population, although this conclusion depends on the interpretation that sources with silicate absorption are obscured AGN.  Where are the high-redshift, luminous starbursts?  Why are they not common in samples observed so far?   The present study is designed to address these issues arising from the previous results. First, we attempt to validate the spectral classification that assigns AGN as the primary source of mid-infrared luminosity in sources observed to contain strong silicate absorption.  This is done by selecting a sample of sources independently classified as AGN, either because of X-ray luminosity or because their spectral energy distributions (SEDs) are power laws with no evidence of a stellar component.  Second, we attempt to verify that high redshift starbursts do indeed exhibit strong PAH emission by selecting a sample of starbursts independently of their mid-infrared spectra.  This is done by selecting sources which show a peak in the SED arising from stellar photospheric luminosity, as observed with the $Spitzer$ Infrared Array Camera (IRAC) \\citep{faz04}.  Mid-infrared spectra of these samples obtained with the $Spitzer$ IRS are presented, and the overall spectral energy distributions (SEDs) are also discussed.  Finally, the space densities of starbursts found in this way are compared to the space densities of AGN at similar redshifts. ",
        "conclusions": "Mid-infrared spectra obtained with the IRS on $Spitzer$ are presented for a sample of 20 sources with f$_{\\nu}$ (24$\\mu$m) $\\ge$ 1.0\\,mJy derived from the SWIRE survey of the Lockman Hole field.  Sources were selected as AGN because of X-ray luminosities (9 sources) or power-law SEDs (2 sources).  Sources were selected as starbursts (9 sources) because of a strong photospheric component in the near-infrared SEDs derived from $Spitzer$ IRAC observations which allowed a photometric redshift estimate.  New redshifts were determined from IRS spectra for 14 sources, with redshifts measured from strong silicate absorption features in 5 AGN sources (1.14 $<$ z $<$ 2.25) and from strong PAH emission features in 9 starburst sources (0.98 $<$ z $<$ 1.91) and one AGN source (z = 1.04). Ten of the 11 AGN sources have IRS spectra which show silicate absorption or are power laws; only one AGN source shows PAH emission features.  All of the 9 sources selected because of starburst classification in the near-infrared SEDs show PAH emission features in the mid-infrared IRS spectra.  These results indicate that optically faint sources ($R$ $\\ga$ 22\\,mag) with f$_{\\nu}$ (24$\\mu$m) $\\ge$ 1.0\\,mJy which contain either a dominant AGN or a dominant starburst as determined from independent classifications at X-ray, optical, and near-infrared wavelengths have a predictable spectral classification in the mid-infrared: the AGN sources have IRS spectra which show silicate absorption or are power laws, and the starburst sources show PAH emission features. The results also show that photometric SEDs derived with IRAC and MIPS fluxes are an efficient method to locate starbursts at z $\\sim$ 1.8.  The starburst galaxies found at high redshift are $\\sim$ 100 times more luminous in PAH emission than local prototypes.  The sample of starburst galaxies gives the result that for f$_{\\nu}$ (24$\\mu$m) $>$ 1.0\\,mJy and 1.7 $\\la$ z $\\la$ 1.9, starburst galaxies have a surface density between $\\sim$ 1 deg$^{-2}$ and 5 deg$^{-2}$, depending on the selection effects that should be applied. This is the redshift interval in which selection at 24\\,\\um optimises detection of starbursts based on strong PAH emission.  By comparing the starburst sample to previous samples of AGN detected by $Spitzer$ to comparable limits of f$_{\\nu}$ (24$\\mu$m) $\\ga$ 1.0\\,mJy, it is concluded that at z $\\sim$ 1.8, the most luminous mid-infrared sources are dominated by AGN.  Surface densities of AGN with f$_{\\nu}$ (24\\,\\um) $>$ 1 mJy and z $\\sim$1.8 exceed the surface densities of starbursts with f$_{\\nu}$ (24\\,\\um) $>$ 1 mJy, even at redshifts where starburst selection is favored because of strong PAH emission within the 24\\,\\um band.  For AGN, absorption in both X-rays and in the mid-infrared silicate feature is compared.  Column densities \\nh\\ derived from X-ray absorption are always found to be greater than the \\nh\\ derived from silicate absorption, by factors that can exceed 100.  There are no correlations within the small sample of sources among X-ray absorbing properties or X-ray luminosities compared to infrared spectral properties, such as infrared spectral slope or extinction in the silicate absorption feature.  However, the AGN have flatter infrared spectra compared to the starbursts as measured by the ratio $\\nu$f$_{\\nu}$ (24\\,\\um)/$\\nu$f$_{\\nu}$ (8\\,\\um)."
    },
    "0608/astro-ph0608579_arXiv.txt": {
        "abstract": "Hardness ratios are commonly used in X-ray photometry to indicate spectral properties roughly. It is usually defined as the ratio of counts in two different wavebands. This definition, however, is problematic when the counts are very limited. Here we instead define hardness ratio using the $\\lambda$ parameter of Poisson processes, and develop an estimation method via Bayesian statistics. Our Monte Carlo simulations show the validity of our method. Based on this new definition, we can estimate the hydrogen column density for the photoelectric absorption of X-ray spectra in the case of low counting statistics. ",
        "introduction": "In recent years, high angular resolution X-ray telescopes make it possible to detect X-ray sources with only a few counts. This is very different from the optical photometry. Because of these low counts, the Poisson processes in corresponding wavebands cannot be approximated to Gaussian distribution. Therefore the statistics will be very different in some estimations and calculations than used before. In recent years Bayesian method has gained many applications (e.g. van Dyk et al. 2001 and references therein) since it has more advantages in low count cases than traditional statistics.  Hardness ratios are widely used in high energy astrophysics since faint sources with only limited counts cannot give any satisfying spectral modeling. In X-ray detection, hardness ratios are normally used to show spectral properties roughly (e.g. Tennant et al., 2001; Sivakoff, Sarazin \\& Carlin 2004). Hardness ratios are usually defined as the ratio of counts in different wavebands ($HR=b/a$) or the ratio of the difference and sum of counts in two wavebands ($HR=(b-a)/(b+a)$), $a$ and $b$ are counts in two wavebands $A$ and $B$. On the other hand, for the spectra of X-ray sources, the photoelectric absorption, quantified with the hydrogen column density $N_H$, cannot be neglected. Hydrogen column density contains many kinds of important information, such as the radial distance and the interstellar circumstance of the sources. For low count sources, hydrogen column density is hard to know since no reliable spectral fitting can be made. However interstellar absorption is energy dependent. Consequently the information of hydrogen column density can be drawn from the hardness ratios. In this paper, we first give a new definition of hardness ratio and its estimation method, and then we discuss the procedure to estimate the hydrogen column density accordingly.   \\bigskip ",
        "conclusions": "First we give the conditional probability distribution of $\\lambda$ parameter under certain counts in a Poisson process using Bayesian statistics. According to this result we derive the probability density function of two kinds of hardness ratios. We take the most probable values as the estimate of hardness ratios and the HPD intervals as the error intervals. Then we use Monte Carlo simulations to investigate the statistical properties of our results, and find that our method is more reliable than the traditional method when the counts are low. Finally we show how to estimate the hydrogen column density using hardness ratios.  Our method developed in this paper provides a way to estimate the hydrogen column density of sources which are too faint to do spectral fitting. However the spectral shape for these sources must be assumed {\\it a prior}. This method is especially convenient for a sample of faint sources with similar spectra.  After this paper has been submitted initially on 06-03-29, we noticed another submitted paper (Park et al. 2006) which discusses the same statistical problem as we have done in this paper. In that paper the authors also used the Bayesian method to estimate the hardness ratio, and showed some applications on quiescent Low-Mass X-ray Binaries, the evolution of a flare, etc, therefore justifying the wide range of applications of such a statistical problem. Since the strict analytic solution of the hardness ratio distribution does not exist for general situations, the authors suggested methods by Monte Carlo and numerical integration to obtain the distribution in that paper. In our paper we find simple analytic solutions of the probability density functions of hardness ratios for the situations in which the background can be ignored. This will be useful and convenient for some applications, such as {\\it Chandra} data in which background can be ignored for hardness ratio estimation of point sources.  Finally, we note, under the advise of the referee, that in 1980s some studies have been done on the ratio of Poisson means both from a frequentist standpoint (James \\& Roos 1980) and from a Bayesian standpoint (Helene 1984 and Prosper 1985). In this paper we used the Bayesian method under the uniform prior and the Jeffreys prior, made extensive comparisons between this method and traditional method, aiming explicitly at applications in astrophysics."
    },
    "0608/hep-th0608072_arXiv.txt": {
        "abstract": "{ Recent perturbative studies show that in 4d non-commutative spaces, the trivial (classically stable) vacuum of gauge theories becomes unstable at the quantum level, unless one introduces sufficiently many fermionic degrees of freedom. This is due to a negative IR-singular term in the one-loop effective potential, which appears as a result of the UV/IR mixing. We study such a system non-perturbatively in the case of pure U(1) gauge theory in four dimensions, where two directions are non-commutative. Monte Carlo simulations are performed after mapping the regularized theory onto a U($N$) lattice gauge theory in $d=2$. At intermediate coupling strength, we find a phase in which open Wilson lines acquire non-zero vacuum expectation values, which implies the spontaneous breakdown of translational invariance. In this phase, various physical quantities obey clear scaling behaviors in the continuum limit with a fixed non-commutativity parameter $\\theta$, which provides evidence for a possible continuum theory. The extent of the dynamically generated space in the non-commutative directions becomes finite in the above limit, and its dependence on $\\theta$ is evaluated explicitly. We also study the dispersion relation. In the weak coupling symmetric phase, it involves a negative IR-singular term, which is responsible for the observed phase transition. In the broken phase, it reveals the existence of the Nambu-Goldstone mode associated with the spontaneous symmetry breaking. } ",
        "introduction": "Non-commutative (NC) geometry \\cite{Sny,Connes} has been studied extensively as a modification of our notion of space-time at short distances, possibly due to effects of quantum gravity \\cite{gravity}. It has recently attracted much attention since gauge theories on a NC geometry have been shown to appear as a low energy limit of string theories with a background tensor field \\cite{String}. At the classical level, introducing non-commutativity to the space-time coordinates modifies the ultraviolet dynamics of field theories (the interaction becomes non-local at the scale of non-commutativity), but not the infrared properties. The latter changes at the quantum level, however, due to the so-called UV/IR mixing effect \\cite{rf:MRS}. As a consequence, one cannot retrieve the corresponding commutative theory by sending the non-commutativity parameter to zero in general.  The UV/IR mixing poses a severe problem in the renormalization procedure within perturbation theory since a new type of IR divergences, in addition to the usual UV divergences \\cite{UV}, appears in non-planar diagrams \\cite{rf:MRS}. The finite lattice formulation \\cite{AMNS} based on twisted reduced models \\cite{GAO} and their new interpretation \\cite{AIIKKT} in the context of NC geometry, regularizes such divergences --- not only the ordinary UV divergences but also the novel IR divergences. It therefore provides a non-perturbative framework to establish the existence of a consistent field theory on a NC geometry. In ref.\\ \\cite{2dU1} a simple field theory, 2d U(1) gauge theory on a NC plane\\footnote{See refs.\\ \\cite{Mafia,Griguolo:2003kq} for related analytic works, including several attempts to actually solve this model.}, has been studied by Monte Carlo simulation, and the existence of a finite continuum limit has been confirmed. The ultraviolet dynamics is described by the commutative 2d U($\\infty$) gauge theory. On the other hand, the infrared dynamics would be described by the commutative 2d U(1) theory, were it not for the UV/IR mixing effects. The ``dynamical Aharonov-Bohm effect'' observed in this regime clearly demonstrates the non-equivalence to the commutative 2d U(1) theory.   As an interesting physical consequence of the UV/IR mixing in the case of scalar field theory, ref.\\ \\cite{GuSo} predicted the existence of a ``striped phase'', in which non-zero Fourier modes of the scalar field acquire a vacuum expectation value and break the translational invariance spontaneously. This prediction, which was based on a self-consistent Hartree-Fock approximation, was supported later by another study using an effective action \\cite{CZ}, and discussed also in the framework of a renormalization group analysis in $4-\\varepsilon$ dimensions \\cite{ChenWu}. Finally, the existence of this new phase was fully established by Monte Carlo simulations \\cite{Procs,AC,Martin:2004un,Bietenholz:2004xs,Julieta}. In ref.\\ \\cite{Bietenholz:2004xs}, in particular, the explicit results for the phase diagram and the dispersion relation were presented in $d=3$. The dispersion relation in the disordered phase shows a positive IR singularity due to the UV/IR mixing, which shifts the energy minimum to a non-vanishing momentum. As one changes the mass parameter in the action, the mode at the energy minimum condenses, which yields the corresponding stripe pattern. The existence of a sensible continuum limit is suggested by the scaling of various physical quantities. In particular the average width of the stripes stabilizes at a finite value in the continuum limit. In even dimensions with the non-commutativity tensor of maximal rank, the emergence of the striped phase can also be conjectured from the eigenvalue distribution of the matrix, which represents the scalar field in NC geometry \\cite{Martin:2004un,Steinacker:2005wj}.   In the case of 4d U(1) NC gauge theories, the IR singularity occurs in the one-loop calculation of the vacuum polarization tensor \\cite{hayakawa}. This changes the dispersion relation of photons at low momenta \\cite{mst}. In fact the IR singularity in the dispersion relation occurs with an overall {\\em negative} sign \\cite{LLT,ruiz,Bassetto:2001vf} unless one introduces sufficiently many fermionic degrees of freedom. It turns out that the manifestly gauge invariant effective action, which is derived from field theoretical calculations as well as from string theoretical calculations \\cite{rf:MVR,rf:AL}, involves the open Wilson line operators \\cite{IIKK}, which implies that the IR singularity is associated with their condensation. This causes the spontaneous breakdown of translational invariance since the open Wilson lines carry specific non-zero momenta dictated by the star-gauge invariance \\cite{IIKK}. In the case of a finite NC torus, the same phenomenon has been understood \\cite{Guralnik:2002ru} from its Morita dual {\\em commutative} SU($N$) gauge theory with twisted boundary conditions\\footnote{At high tempertature (on the scale of the non-commutativity) the ring diagrams dominate. Their evaluation to all orders implies a critical temperature, above which the magnetic photon mass squared turns negative in the NC planes \\cite{BFM}.}.  We should note, however, that these perturbative analyses merely imply that the {\\em perturbative vacuum} of 4d NC gauge theory has tachyonic instability, but they do not exclude the possibility that the system eventually finds a stable vacuum. (An analogous scenario known as the ``tachyon condensation'' has been proposed in the open string field theory \\cite{tachyon-condensation}\\footnote{This might be more than a plain analogy since the tachyonic instability of NC gauge theory has partial contribution from the closed string tachyon \\cite{rf:AL,Seiberg:2000zk}.}.) In order to investigate this issue, we definitely need to employ a non-perturbative framework.  In the current work, the lattice formulation \\cite{AMNS} is particularly important since it allows us to study the model from first principles by Monte Carlo simulation. We study pure U(1) gauge theory in four dimensions, where two directions are non-commutative, and we find a phase in which open Wilson lines acquire non-zero vacuum expectation values. In fact this phase extends towards weak coupling as the system size is increased, and it turns out that we always end up in this phase in the continuum and infinite-volume limits with a fixed non-commutativity parameter. This is consistent with the prediction from the one-loop calculations. We observe, however, that various physical quantities obey clear scaling behaviors in the above limit, which provides evidence for a possible continuum theory. We also study the dispersion relation. In the weak coupling symmetric phase, it involves a negative IR-singular term, which makes the energy vanish at a non-zero momentum as one approaches the critical point. In the broken phase, although the momentum components in the NC directions are no longer conserved, we can still study the relation between the energy and the momentum in the {\\em commutative} direction. This reveals the existence of the Nambu-Goldstone mode associated with the spontaneous symmetry breakdown.   This paper is organized as follows. In section \\ref{sec:review} we define a gauge theory in 4d NC space, and we briefly review the results obtained by perturbative calculations. In section \\ref{sec:lattice} we introduce the lattice formulation and discuss how to take the continuum limit. In section \\ref{sec:phase} we present the phase diagram of the lattice model. In section \\ref{sec:double-scaling} we investigate the existence of a sensible continuum limit. In section \\ref{sec:phys-extent-NC} we study how the space in the NC directions looks like in the broken phase. In section \\ref{sec:dispersion-rel} we study the dispersion relation in each phase to gain deeper understanding of the observed phase transition and the continuum limit. Section \\ref{sec:summary} is devoted to a summary and discussions. In the appendix we explain the algorithm used for our Monte Carlo simulations. Some preliminary results of this work have been presented in proceeding contributions \\cite{proc}. ",
        "conclusions": "\\label{sec:summary}  In this paper we studied four-dimensional gauge theory in NC geometry from first principles based on its lattice formulation. In particular we clarified the fate of the tachyonic instability encountered in perturbative calculations. The lattice formulation is suited for such a study since the IR singularity responsible for the instability is regularized in a star-gauge invariant manner, and we can trace the behavior of the system as the regularization is removed. This revealed the existence of a first order phase transition associated with the spontaneous breakdown of the U(1)$^2$ symmetry, which includes the translational symmetry in the NC directions as a discrete subgroup.  In the weak coupling symmetric phase, we studied the dispersion relation and confirmed the presence of an IR singularity. This IR singularity prevents us from taking the continuum limit in the symmetric phase. While we are able to observe scaling up to some finite $N$, we cannot increase $N$ further since the energy of the lowest momentum mode vanishes, and the vacuum becomes unstable. In the broken phase, however, we provided evidence for a sensible continuum limit. The phase is characterized by the condensation of open Wilson lines, which represent the ``tachyon'' in the unstable perturbative vacuum. By studying the dispersion relation in the broken phase, we confirmed the appearance of the Nambu-Goldstone mode associated with the spontaneous breakdown of the U(1)$^2$ symmetry.   By measuring the eigenvalue distribution of $V_\\mu$ in the NC directions, we find that the dynamical space in the NC directions has shrunk, but it has a finite {\\em physical} extent in the continuum limit. An analogous first order phase transition is found in gauge theories on fuzzy manifolds \\cite{ABNN}, where the fuzzy manifolds collapse at sufficiently large couplings. The instability in those cases is due to the uniform condensation of a scalar field on the fuzzy manifold\\footnote{The fuzzy manifolds can be stabilized by adding a sufficiently large mass to the scalar field \\cite{ydri}. }. The phenomenon that the space-time itself becomes a dynamical object is characteristic to gauge theories on NC geometry. This should be closely related to the dynamical compactification of extra dimensions\\footnote{As a closely related line of research, see refs.\\ \\cite{CSDR}, which use fuzzy spheres for compactified dimensions. } in string theory \\cite{SSB,related} based on the IKKT matrix model \\cite{Ishibashi:1996xs}. Let us recall that the NC geometry appears in string theory as a result of introducing a background tensor field \\cite{String}. Our conclusion therefore suggests that the dynamical compactification in string theory may be associated with the spontaneous generation of the background tensor field in six dimensions. This is reminiscent of the spontaneous magnetization in 4d non-Abelian gauge theory at finite temperature \\cite{sp-mag} or in 3d QED \\cite{Hosotani}.        On the other hand, if we wish to obtain a phenomenologically viable {\\em 4d model}, we may stay in the symmetric phase by keeping the UV cutoff finite and view the NC gauge theory as an effective theory of a more fundamental theory. A $\\theta$-deformed dispersion relation for the photon such as the one displayed in figure \\ref{eps:dispersion} should then have implications \\cite{Camel} on observational data from blazars (highly active galactic nuclei) \\cite{Aha}, which are assumed to emit bursts of photons simultaneously, covering a broad range of energy. Such experimental efforts will be intensified in the near future. For instance, the Gamma-ray Large Area Space Telescope (GLAST) project \\cite{GLAST} is scheduled to be launched in September 2007 and to monitor gamma rays from 20 MeV up to 1 TeV. In particular, a relative delay of these photons depending on the frequency \\cite{galax} could hint at a NC geometry."
    },
    "0608/astro-ph0608029_arXiv.txt": {
        "abstract": "We present a study of the type IIn supernova (SN) 2005gl, in the relatively nearby ($d\\approx66$ Mpc) galaxy NGC 266. Photometry and spectroscopy of the SN indicate it is a typical member of its class. Pre-explosion {\\it Hubble Space Telescope (HST)} imaging of the location of the SN, along with a precise localization of this event using the Laser-Guide-Star assisted Adaptive Optics (LGS-AO) system at Keck Observatory, are combined to identify a luminous ($M_V=-10.3$) point source as the possible progenitor of SN 2005gl. If the source is indeed a single star, it was likely a member of the class of luminous blue variable stars (LBVs). This finding leads us to consider the possible general association of SNe IIn with LBV progenitors. We find this is indeed supported by observations of other SNe, and the known properties of LBV stars. For example, we argue that should the prototypical Galactic LBV $\\eta$ Carina explode in a phase similar to its current state, it will likely produce a type IIn SN. We discuss our findings in the context of current ideas about the evolution of massive stars, and review the census of SNe with identified progenitors. We introduce the concept of the progenitor-SN map as a convenient means to discuss the present status and future prospects of direct searches for SN progenitors. We conclude that this field has matured considerably in recent years, and the transition from anecdotal information about rare single events to robust associations of progenitor classes with specific SN types has already begun. ",
        "introduction": "It is generally assumed that supernovae (SNe) can be divided into two physical classes. Type Ia SNe are assumed to result from the thermonuclear explosion of a degenerate white dwarf star, reaching the critical ignition density as it approaches the Chandrasekhar limit by accretion from, or merger with, a binary companion. Direct observational evidence shows that all other types of SNe result from the gravitational core-collapse of young, massive stars. These progenitors are expected to be relatively luminous, and are thus potentially detectable in images of sufficient spatial resolution and depth obtained before these core-collapse SNe explode.  The impact of the study of SN progenitors was poignantly illustrated by the watershed case of SN 1987A and its blue supergiant progenitor (White \\& Malin 1987). Initial surprise at the color (blue rather than red) and compactness of this progenitor led to revisions in our understanding of massive star evolution and SN explosion physics. During the next 15 years progress was slow, with but a single additional progenitor identified (SN 1993J; Aldering, Humphreys, \\& Richmond 1994; Van Dyk et al. 2002; Maund et al. 2004). Pioneering work by Van Dyk and collaborators (Barth et al. 1996; Van Dyk et al. 1999; Van Dyk, Li, \\& Filippenko 2003a) utilized a new resource - the sensitivity and resolution afforded by pre-explosion images obtained by the Hubble Space Telescope ({\\it HST}). Several possible progenitors have been identified, but these associations were often inconclusive due to the SN astrometry being limited by post-explosion ground based images of relatively poor quality.  In the last few years, breakthrough results were presented by two groups (the California group, e.g., Van Dyk et al. 2002; Van Dyk, Li, \\& Filippenko 2003b, Li et al. 2006, and the UK group, e.g., Smartt et al. 2004; Hendry et al. 2006) using mostly post-explosion {\\it HST} imaging to precisely determine the location of SNe and securely identify progenitor stars in pre-explosion {\\it HST} images.  Most recently, we have introduced the use of laser guide star assisted adaptive optics (LGS-AO) as an alternative means for precise SN localization that is independent of HST scheduling and operations, and not subject to saturation by the brightness of a young SN (Gal-Yam et al. 2005a,b). Here, we report the second result from our program at Keck Observatory, the discovery of a luminous point source in pre-explosion images of the type IIn SN 2005gl. The paper layout is as follows. In $\\S~2$ we present our observations, including analysis of archival pre-explosion {\\it HST} images of the location of this event, and our Keck LGS-AO post-explosion observations leading to precise localization of the SN on the pre-explosion grid. In $\\S~3$ we present the discovery of a point source consistent with being a very luminous LBV-type progenitor of this SN. We conclude in $\\S~4$ with a discussion of our result in the context of accumulated information about SNe IIn and all available SN progenitor identifications to date. ",
        "conclusions": "\\subsection{SNe IIn from LBVs}  We now examine the hypothesis that LBV stars explode as SN IIn. The class of type IIn SNe is known to be a heterogeneous group of events. At least two subsets represent specific and probably unrelated phenomena to the one we consider here, namely, the core-collapse-driven explosion of a massive LBV star, similar, for example, to the well-known $\\eta$-Carina in our Galaxy. The first unrelated class are the so-called ``SN impostors'', which are believed to be super-outbursts of LBVs which do not result in total disruption of the progenitor star (see, e.g., Van Dyk 2005 and Maund et al. 2006, and references therein). These events are typically faint (compared to a normal SN) with absolute magnitudes between $-10$ and $-14$ (Van Dyk 2005), while ``genuine'' type IIn explosions (Schlegel 1990), are much brighter (sometime reaching $M=-20$; SN 2005gl has $M\\sim-17$). The second group of events which are also probably irrelevant to our discussion are the recently discovered class of ``hybrid'' Ia/IIn SNe, including the prototype SN 2002ic (Hamuy et al. 2003; Wood-Vasey et al. 2004) along with its recent (SN 2005gj, Prieto et al. 2005; Aldering et al. 2006) and past (SN 1997cy, Germany et al. 2000, Turrato et al. 2000; SN 1999E, Rigon et al. 2003) clones (Hamuy et al. 2003, Deng et al. 2004, Wang et al. 2004). These SNe display spectral properties similar to those of thermonuclear SNe Ia, along with narrow Balmer lines that probably result from strong interaction with circumstellar material in the immediate vicinity of the exploding star. Such events have been suggested to be thermonuclear SNe Ia, exploding either in close proximity to recently stripped gas from a binary companion with intense mass loss (e.g., an AGB star, Hamuy et al. 2003); within a symbiotic system; or perhaps thermonuclear explosion that occur while the envelope of the exploding star is still intact (''SNe 1.5'', Iben \\& Renzini 1983). In any case, SN 2005gl does not share the spectroscopic properties of these hybrid events, and does not appear to be related to this IIn sub-class (Fig.~\\ref{specfig}). We note that both of the sub-groups discussed above comprise a minority within the observed population of SNe IIn.  So, we can reformulate our question to be: can LBVs be the progenitors of most SN IIn explosions? We focus on the best-studied example of an LBV in our Galaxy, $\\eta$ Carina. The envelope of this star still contains large quantities of hydrogen (e.g., Davidson et al. 1986), so, had it exploded now, it would result in a type II (H-rich) SN. In addition, we observe the results of copious mass loss around this system, of order several solar masses of ejected material, mostly H. Were an energetic explosion to occur at the center of such a huge debris cloud, we would expect strong interaction leading to strong narrow H lines - the hallmark of type IIn SNe. Thus, we conclude that an LBV exploding during the active mass-ejection phase, or shortly thereafter, would indeed appear to distant observers as a type IIn SN.  The rates of SNe~IIn may shed some light on our proposed progenitor association. Let us consider a simplistic massive star evolutionary scheme, for single stars with approximately solar metallicity, broadly following, e.g., Maeder and Conti (1994). We use the customary notations RSG for red supergiants, WN for N-rich (and usually also He-rich) Wolf-Rayet stars, WC for C-rich W-R stars, and WO for O-rich W-R stars.  \\begin{equation} 80 M_{\\odot}< M < 150 M_{\\odot}: O \\rightarrow LBV \\rightarrow {\\rm SN IIn} (?) \\end{equation} \\begin{equation} 40 M_{\\odot}< M < 80 M_{\\odot}: O \\rightarrow LBV \\rightarrow WN \\rightarrow WC/WO \\rightarrow {\\rm SN Ic} \\end{equation} \\begin{equation} 25 M_{\\odot}< M < 40 M_{\\odot}: O \\rightarrow LBV \\rightarrow (early) WN \\rightarrow {\\rm SN Ib} \\end{equation} \\begin{equation} 15 M_{\\odot}< M < 25 M_{\\odot}: O \\rightarrow RSG \\rightarrow (late) WN \\rightarrow {\\rm SN IIL/IIb} \\end{equation} \\begin{equation} 8 M_{\\odot}< M < 15 M_{\\odot}: B/O \\rightarrow RSG \\rightarrow {\\rm SN IIP} \\end{equation}  We note that SNe IIn can, in this picture, occur in two distinct cases. First, it is believed that massive stars undergo short LBV phases, involving rapid and strong mass loss which transform H-rich supergiants to W-R stars, as well as during the early W-R evolution stages (while the stars still have significant amounts of H, i.e., are WN stars; see, e.g., Maeder \\& Conti 1994). Furthermore, since the evolution of the inner core is decoupled from that of the envelope during these stages, in some cases the core might collapse during those short-lived transition phase, resulting in SNe IIn from stars in the mass range, comparable to that of the progenitors of SNe IIb/L and Ib/c. In this context we note the recent work by Chugai \\& Chevalier (2006) interpreting the observations of SN 2001em, which exploded as a SN Ic, and then developed strong interaction signatures (a bright radio signal accompanied by SN IIn-like narrow H$\\alpha$ lines, Stockdale et al. 2004; Soderberg, Gal-Yam, \\& Kulkarni 2004; Bietenholz \\& Bartel 2005). Models by these authors suggest that SN 2001em underwent a violent and intense, hydrogen-rich, mass ejection episode, shortly ($\\simlt 1000$) years before the SN explosion. It appears SN 2001em exploded shortly after it exited a recent LBV phase, and a similar event, occurring slightly (on stellar evolution timescales) earlier, would appear as a SN IIn.  The second case is a speculative option which we now describe. We hypothesize that in the case of the most massive stars (Eq. 1) the evolution of the core might be rapid enough to overtake the progress of envelope mass loss, and that these stars might undergo core-collapse while in their first LBV phase, before they are significantly stripped. SNe IIn in this case take the place of SNe Ic as the end products of the rare, most massive stars. It is interesting to note in this context that explosion models of SNe Ic, including the most energetic ones associated with GRBs (e.g., Deng et al. 2005; Mazzali et al. 2006) require massive progenitors (with zero-age mass $\\simlt50 M_{\\odot}$), but not as massive as the claimed masses of the most extreme LBVs ($>80 M_{\\odot}$), e.g., $\\eta$-Carina (Figer et al. 1998), LBV 1806 (Eikenberry et al. 2004) or the Pistol star (Figer et al. 1998). We note that this scenario, in which the core evolution overtakes the envelope stripping, is not in accord with current stellar evolution models. Still, we think it is an interesting option that merits further theoretical and observational examination.  We note that Chugai et al. (2004) interpret their observations of the type IIn SN 1994W as indicating that the progenitor underwent an explosive mass ejection shortly before its core collapsed and led to its final explosion as a SN. This would fit well with the above picture of SNe IIn resulting from LBVs like $\\eta$-Carina (which are known to undergo extreme and violent massive loss episodes) during or shortly after their active phase.  In either of the above cases, we expect these SNe from LBV progenitors to be rare (either because LBVs are short-lived phases, or because the progenitors are very rare stars). For example, the total number of stars with initial mass above $80 M_{\\odot}$ is $\\sim3 \\% $ of the total core-collapse progenitor population ($8M_{\\odot}<M<150M_{\\odot}$) for a Salpeter (1955) initial mass function. This is similar to the estimated fraction of SNe IIn of the total core-collapse population ($\\sim2 \\%$, Cappellaro et al. 1997). We conclude that the notion that some or most of SNe IIn result from explosions of LBVs appears broadly consistent with available observational data.  \\subsection{SN 2005gl in context: The SN-Progenitor map}  In Fig.~\\ref{mapfig} we summarize the accumulated knowledge from direct observations of SN progenitor stars in the form of a progenitor-SN map. As possible progenitor classes we consider red and blue supergiants (RSGs and BSGs), LBVs, early (He rich) and late (C/O rich) Wolf-Rayet (W-R) stars, and massive binaries. On the SN side we list all the major well-defined sub-classes of core-collapse SNe. Among the H-rich type II subclasses, SNe II-P have extended plateau optical light curves, while SN 1987A-like events are fainter and have a late hump-like peak in their optical light curves. SNe IIn (excluding the specific subsets described in $\\S~4.1$ above) generally have blue and rather featureless spectra with Balmer emission-line profiles which include a narrow component ($<1000$ km s$^{-1}$) and are often very luminous in the optical, consistent with a strong contribution from shocked CSM. The transition class of SNe IIL/IIb have rapidly declining light curves and often display strong He lines, developing into events with spectra similar to SNe Ib (see below). Of the H-poor type I SNe, type Ib events are dominated by He lines, while type Ic lack both H and He and are dominated by intermediate-mass elements such as O, Mg, Si and Ca. It must be stressed that this is, by necessity, a simplified presentation of the core-collapse SN ``zoo'', and that many peculiar objects which do not strictly fit into this picture, or transition objects that are either intermediate or evolve between classes, are known. Still, this scheme should include the majority of observed objects.  We suggest this is a useful tool to present and discuss this topic. Of course, we have no guaranties that this map is complete (perhaps more progenitor classes or SN types are required) nor that it is a 1-to-1 mapping. Most likely this version will evolve and change. Let us use it for the time being as a basis for discussion. We note that the association of type II-P SNe with low-mass RSG progenitors is well-supported by 4 observed events, with no counter examples, and maintain this association is now progressing from suggestive to robust (including also numerous supporting upper limits; Smartt et al. 2006). Studies of the (H and He-poor) type Ic SN 2004gt (Gal-Yam et al. 2005; Maund et al. 2005) and to a lesser degree SN 2002ap (Smartt et al. 2002) suggest these had highly evolved (i.e., optically less-luminous) W-R progenitors, perhaps of the WC or WO subclasses. This may fit well with the population of less-evolved W-R stars (late WN) exploding as He-rich type Ib events, consistent with recent explosion models by Tominaga et al. (2005). This picture is consistent with that expected from stellar evolution theory. The case of the transition class of type IIb objects (which evolve from having H-dominated type II spectra toward the He-dominated type Ib class) is less clear. Observations of the prototypical event SN 1993J in M81 suggest a possible massive binary progenitor (Aldering et al. 1994; Van Dyk et al. 2002; Smartt et al. 2004) with some support from the recent work on SN 2001ig (Ryder et al. 2004; 2006). However, it is unclear whether the nearby putative binaries were indeed physically important or even bound to the SN progenitors, or perhaps (some) SNe IIb arise from single WN (H rich) W-R stars (Soderberg et al. 2005).  Finally, we note that type IIn events, which show evidence for intense interaction with circumstellar material, suggesting heavy mass-loss from the progenitor, appear to require a different progenitor from those so far discussed. Red supergiants have little mass loss, and so had the blue supergiant progenitor of SN 1987A. While LBVs are not the only option, they provide an elegant solution, as discussed above. It would appear that our putative map now encompasses all main SN types. Of course, more data are required to establish a robust mapping, and we would like each line to be supported by several progenitor-SN associations (as is the case for SNe II-P and red supergiants). Additional cases may also allow an attempt to understand the effects of other parameters (beyond initial mass) such as metallicity, rotation and binarity.  \\begin{figure*} \\includegraphics[width=15cm]{f5.eps} \\caption{The progenitor-SN map, presenting associations of SNe with progenitor stars based on direct observations of the progenitors in pre-explosion images. The association of the most common type of core-collapse events, SNe II-P, with (relatively) low-mass ($8 M_{\\odot}< M <15 M_{\\odot}$) red supergiants (RSGs) appears to be quite robust. The connections between fainter SN II events (1987A-like) and the type IIn sub-class, with blue supergiants and LBVs, respectively, is but suggestive, based on a single event in each case. The progenitors of stripped-envelope events (type Ib/c) are yet to be observed, but analysis of available upper limits suggest that the progenitors of the more highly stripped (type Ic) SNe are more evolved (and less luminous) ``early'' W-R stars, while SNe Ib, which still retain much of their He envelope, result from somewhat less evolved (and often more luminous) ``late'' W-R stars. Data used are from White \\& Malin 1987; Aldering et al. 1994; Van Dyk et al. 2002; Maund et al. 2004; Smartt et al. 2004; Van Dyk et al. 2003b; Li et al. 2005; Maund et al. 2005a; Li et al. 2006; Hendry et al. 2006; Gal-Yam et al. 2005a; Maund et al. 2005b; and this work. } \\label{mapfig} \\end{figure*}  \\subsection{Conclusions}  We have identified a luminous point source in pre-explosion images of the location of the typical type IIn SN 2005gl. If this was a single star it was most likely an LBV, based on luminosity considerations. We have discussed a major caveat - the possibility that this point-like source is actually a compact luminous cluster - and ways by which this progenitor identification can be tested by future observations. We suggest that an association of SNe IIn with LBV progenitors accommodates much of what is currently known about both LBVs and SNe IIn. We introduce the progenitor-SN map and use it as a basis to discuss our finding. This field has rapidly progressed in recent years and should evolve from a collection of isolated fragmented pieces of information to become based on more solid concepts (such as the progenitor-SN map). These should soon begin to provide useful constraints on models, and drive additional progress in understanding the last stages of massive star evolution and the physics of SNe."
    },
    "0608/astro-ph0608503_arXiv.txt": {
        "abstract": "Wilkinson Microwave Anisotropy Probe (WMAP) has provided us with the highest resolution all-sky maps of the Cosmic Microwave Background. As a result of thermal Sunyaev-Zel'dovich effect, clusters of galaxies are imprinted as tiny, poorly resolved dips on top of primary CMB anisotropies in these maps. Here, I describe different efforts to extract the physics of Intracluster Medium (ICM) from the sea of primary CMB, through combining WMAP with low-redshift galaxy or X-ray cluster surveys. This finally culminates at a mean (universal) ICM pressure profile, which is for the first time directly constrained from WMAP 3yr maps, and leads to interesting constraints on the ICM baryonic budget. ",
        "introduction": "Since the discovery of dark matter by Zwicky in 1937 \\cite{1937ApJ....86..217Z}, which was done through study of virial velocities of galaxies in clusters, galaxy clusters have been in the forefront of cosmological studies.  One of the most promising tracers of galaxy clusters is the thermal Sunyaev-Zel'dovich (SZ) effect, induced in the Cosmic Microwave Background (CMB) sky through scattering of CMB photons off hot electrons in the Intracluster Medium (ICM) \\cite{sunyaev72}. Most significant properties of SZ effect as a tracer of clusters can be summarized as \\cite{2002ARA&A..40..643C}: \\begin{itemize} \\item It is proportional to the line of sight integral of the ICM pressure. \\item It has a characteristic frequency dependent signature which enables observers to extract it from the CMB background. \\item The total SZ flux of a cluster does not depend on redshift for $z\\gtrsim 1$, making SZ effect an ideal tool for making high redshift cluster surveys. \\end{itemize}  In contrast to free-free X-ray emission, which traces the square of electron density, the fact that SZ effect follows gas pressure makes it a more representative tracer of gas density, and less sensitive to small scale structure or astrophysical processes (e.g. \\cite{2006astro.ph..1133R}). Therefore, it is believed that SZ flux can be used as a reliable tracer of cluster mass with relatively low scatter.  This has motivated many SZ cluster surveys, which are expected to give a census of dark matter halo mass function at high redshifts, and thus enable novel constraints on the expansion history of the Universe \\cite{2006astro.ph..5575B}.  However, despite theoretical expectations, it is not quite clear how well SZ observations can be calibrated to trace dark matter halo properties. Therefore, complementary X-ray and SZ observations, as well numerical studies are necessary to obtain a realistic picture of the ICM properties.  In this note, I outline efforts to extract the SZ effect from the highest resolution all-sky maps of CMB produced by Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{bennett03,hinshaw06}, the resulting constraints on ICM pressure profile, and comparison with other X-ray and SZ observations, as well as numerical studies. ",
        "conclusions": "As the era of the next generation of CMB experiments is fast approaching, it is vital to understand the kind of signals that these experiments are expected to see, the most significant of which is the thermal SZ effect. Moreover, the SZ effect is believed to be the most efficient and robust way of detecting clusters at high redshifts, which can be used to constrain the history of dark energy in that era. To what extent these promises can be realized is the subject of on-going research and inquiry.  In this note, we argued that despite its low resolution, WMAP, as the highest resolution existing all-sky CMB map, can bring in a unique perspective into the understanding of cluster SZ properties, and ICM physics in general. However, this information can only be extracted through combination with a tracer of low redshift universe, such as galaxy surveys or X-ray cluster observations. We described different methods to exploit this SZ signal, which involve varying degrees of theoretical modeling. In particular, WMAP provides the only available observations of the ICM outskirts, and thus for the first time constrain the cluster total SZ flux (or thermal energy).  Both X-ray and SZ observations of the ICM are prone to systematics and technological limitations that could complicate or confuse their interpretation. While on-going theoretical/numerical studies go a long way in understanding these systematics, it would not be the first time that nature outdid the Astronomers' imagination. Therefore, complementary cluster observations (and surveys) in SZ, X-rays, and gravitational lensing (along with numerical studies) is the only guaranteed way to obtain a comprehensive picture of the ICM physics, and thus be able to reliably use them as a {\\it Dark Energy machine}.        The topics discussed here have been borne out of past, or on-going collaborations with Yen-Ting Lin, Yeong-Shang Loh, Daisuke Nagai, Alastair Sanderson, and Michael Strauss. I would like to thank them for continuous scientific discourse and education."
    },
    "0608/astro-ph0608445_arXiv.txt": {
        "abstract": "{} {We present \\emph{FUSE} spectra of the giant H\\2\\ region NGC604 in the spiral galaxy M33. Chemical abundances are tentatively derived from far-UV absorption lines and compared to those derived from optical emission lines.} {Absorption lines from neutral hydrogen and heavy elements were observed against the continuum provided by the young massive stars embedded in the H\\2\\ region. We derived the column densities of H\\1, N\\1, O\\1, Si\\2, P\\2, Ar\\1, and Fe\\2, fitting the line profiles with either a single component or several components. We used \\texttt{CLOUDY} to correct for contamination from the ionized gas. Archival HST/\\emph{STIS} spectra across NGC604 allowed us to investigate how inhomogeneities affect the final H\\1\\ column density.} {Kinematics show that the neutral gas is physically related to the H\\2\\ region. The \\emph{STIS} spectra reveal H\\1\\ column density fluctuations up to 1\\,dex across NCG604. Nevertheless, we find that the H\\1\\ column density determined from the global \\textit{STIS} spectrum does not differ significantly from the average over the individual sightlines. Our net results using the column densities derived with \\textit{FUSE}, assuming a single component, show that N, O, Si, and Ar are apparently underabundant in the neutral phase by a factor of $\\sim 10$ or more with respect to the ionized phase, while Fe is the same. However, we discuss the possibility that the absorption lines are made of individual unresolved components, and find that only P\\2, Ar\\1, and Fe\\2\\ lines should not be affected by the presence of hidden saturated components, while N\\1, O\\1, and Si\\2 might be much more affected.} {If N, O, and Si are actually underabundant in the neutral gas of NGC604 with respect to the ionized gas, this would confirm earlier results obtained for the blue compact dwarfs, and their interpretations. However, a deeper analysis focused on P, Ar, and Fe mitigates the above conclusion and indicates that the neutral gas and ionized gas could have similar abundances.} ",
        "introduction": "The interstellar medium (ISM) is mainly enriched by heavy elements produced by the young massive stars during many starburst episodes taking place over the star-formation history of the galaxy. The immediate fate of metals released by these massive stars in the H\\2\\ regions where stars recently formed, has not yet been settled. Kunth \\&\\ Sargent (1986) have suggested that the H\\2\\ regions of the blue compact dwarf galaxy \\IZW\\ enrich themselves with metals expelled by supernov\\ae\\ and stellar winds during the timescale of a starburst episode (i.e., a few $10^6$\\,yr). However, observational evidence (see e.g., Martin et al. 2002) shows that metals might be contained in a hot phase reaching the halo of galaxies, before they could cool down and eventually mix into the ISM. The issue of the possible self-enrichment of H\\2\\ regions is essential to understanding the chemical evolution of galaxies, since H\\2\\ region abundances derived from the optical emission lines of the ionized gas are extensively used to estimate the metallicity of galaxies. If there is self-pollution of these regions, the derived abundances would no longer reflect the actual abundances of the ISM.  One approach to studying possible self-enrichment and, more generally, the mixing of heavy elements in the ISM is to compare the abundances of the ionized gas to those of the surrounding neutral gas. A first attempt was made by Kunth et al. (1994) using the \\textit{GHRS} onboard the Hubble Space Telescope (HST). These authors derived the neutral oxygen abundance in \\IZW\\ using the O\\1\\ $\\lambda$1302 line arising in absorption and using the H\\1\\ column density and velocity dispersion from radio observations. They found an oxygen abundance that is 30 times lower than in the ionized gas. However, the result remains inconclusive since the O\\1\\ line is likely to be heavily saturated, so that its profile could be reproduced with a solar metallicity by choosing a different dispersion velocity parameter, as pointed out by Pettini \\&\\ Lipman (1994).  The \\textit{Far Ultraviolet Spectroscopic Explorer} (\\textit{FUSE}; Moos et al. 2000) gives access to many transitions of species arising in the neutral gas, including H\\1, N\\1, O\\1, Si\\2, P\\2, Ar\\1, and Fe\\2, with a wide range of oscillator strengths for most of the species. Hence it becomes possible to determine abundances in the neutral gas with better accuracy.  A surprise came from recent \\textit{FUSE} studies of five gas-rich blue compact metal-poor galaxies (BCDs): \\IZW\\ (Aloisi et al. 2003; Lecavelier et al. 2004), \\Mark\\ (Thuan et al. 2002), \\IZ\\ (Lebouteiller et al. 2004), \\SBS\\ (Thuan et al. 2005), and NGC625 (Cannon et al. 2004). In these galaxies, nitrogen seems to be systematically underabundant in the neutral gas with respect to the ionized gas of the H\\2\\ regions. The oxygen abundance seems either lower in the neutral gas or similar to the one in the ionized gas. This picture, if true, could offer a new view of the chemical evolution of the ISM in a galaxy, especially for the metal dispersion and mixing timescales. It could, however, suffer from uncertainties such as ionization corrections, depletion effects, or systematic errors due to the multiple sightlines toward individual stars and, possibly, multiple H\\2\\ regions within the spectrograph entrance aperture.  In this respect, nearby giant H\\2\\ regions in spiral galaxies provide an interesting case, since only one H\\2\\ region fits the aperture. Moreover, their resolved young stellar population can be investigated to analyze the stellar continuum and in order to understand and correct for the effects of having multiple sightlines contributing to the absorption line profiles. NGC604 is the first nearby giant H\\2\\ region we have investigated for this purpose. It is a young star-forming region in the dwarf spiral galaxy M33 with an age of 3-5\\,Myr (e.g., D'Odorico \\&\\ Rosa 1981; Wilson \\&\\ Matthews 1995; Pellerin 2006). NGC604 is the brightest extragalactic H\\2\\ region in the sky after 30~Doradus. It is located about 12' from the center of M33. At a distance of 840\\,kpc (Freedman et al. 1991), 1\"~corresponds to 4.1\\,pc. The nebula has a core-halo structure, where the core has an optical diameter of $\\sim 220$\\,pc and the halo a diameter of $\\sim 440$\\,pc (Melnick 1980). The optical extinction in NGC604 appears to be correlated with the brightness. It is highest toward the brightest regions ($A_V = 1.7 \\textrm{-} 2.8$\\,mag), while the average extinction over the nebula is $A_V \\sim 0.5$\\,mag, suggesting that dust is correlated with the ionized gas (Churchwell \\&\\ Goss 1999, Viallefond et al. 1992).  \\begin{figure} \\epsfxsize=7cm \\rotatebox{-90}{\\epsfbox{3161fig1.ps}} \\hspace{0.4cm} \\caption[]{\\small{\\textit{FUSE} and HST/\\textit{STIS} apertures are plotted over the optical image from the POSS2 survey (B band). We superimposed the most massive stars in black as observed with \\textit{STIS} at 2000\\,\\AA. The arrows show the dispersion direction for each slit. North is up.}} \\label{fig:slits} \\vspace{0.1cm} \\end{figure}  We first describe the observations in Sect.~\\ref{sec:obs} and the data analysis in Sect.~\\ref{sec:data}, with particular emphasis on the influence of the source extent on the absorption line profiles. In Sect.~\\ref{sec:mol} we analyze the diffuse molecular hydrogen content, while in Sect.~\\ref{sec:neutral} we infer the neutral hydrogen column density. Metal column densities in the neutral gas are discussed in Sect.~\\ref{sec:heavy}. We estimate the neutral gas chemical composition in Sects.~\\ref{sec:model} and~\\ref{sec:ab} and eventually compare it with the ionized gas of the H\\2\\ region. Finally, in Sect.~\\ref{sec:hidden}, we investigate how the presence of hidden saturated unresolved absorption components can be a major problem for determining the column density. ",
        "conclusions": "This study provides the first detailed analysis of interstellar lines of H\\1, N\\1, O\\1, Si\\2, P\\2, Ar\\1, and Fe\\2\\ in the neutral medium in front of a giant H\\2\\ region in the spiral galaxy M33. Since NGC604 is a nearby system, we have been able to perform the necessary critical tests for analyzing possible selection effects.  \\begin{itemize}  \\item  In the frame of a simple model, we have derived column densities of metals in the neutral gas of NGC604 from both \\textit{FUSE} LWRS and MDRS spectra. These independent estimates allowed us to quantify the effects of the source extent on the interstellar absorption line profiles.  \\item The continuum used for the profile fitting was checked by comparison with a theoretical stellar model. Besides, no significant contamination from stellar photospheric lines was found for the H\\1 absorption lines we investigated.  \\item A particular attention was given to the neutral hydrogen column-density determination. The Ly$\\beta$ line from the \\textit{FUSE} LWRS and MDRS spectra is contaminated by the O\\6\\ P~Cygni doublet. We modelled this contamination to obtain the H\\1\\ column density. Also, profile fitting of the Ly$\\alpha$ absorption in HST/\\textit{STIS} spectra toward individual stars in the cluster NGC604 reveals inhomogeneities in the neutral gas. We finally adopted the FUSE value of $\\log~N$(H\\1)=20.75 with a conservative error of $\\pm$0.3~dex to account for all the possible uncertainties. Within the errors, this H\\1\\ column density is consistent with the \\textit{IUE} determination using Ly$\\alpha$ and with the 21~cm line radio observations.  \\item By modelling the ionization structure of the H\\2\\ gas with the photoionization code \\texttt{CLOUDY}, we have shown that N\\1, O\\1, Ar\\1, and Fe\\2\\ are reliable tracers of the neutral gas, in contrast with Si\\2\\ and P\\2, which require ionization corrections to obtain final abundances in the neutral phase.  \\end{itemize}  Adjusting absorption lines with a single component, we find that N, O, Ar, and Si are underabundant in the neutral gas as compared to the ionized gas by factors $\\gtrsim10$, while Fe/H is similar in the two gaseous phases. This result is rather puzzling, since iron is expected to be equally or more depleted in grains in the neutral gas compared to the ionized one, and there is no reason it should be relatively more abundant in the neutral gas.  This led us to investigate in detail the influence of individual unresolved components in the analysis of absorption lines. Using the method and results of Lebouteiller (2005), we argue that N\\1, O\\1, and Si\\2\\ column densities can be severely underestimated if there are saturated hidden components, while Fe\\2\\ and, to some extent, Ar\\1\\ and P\\2\\ should be more reliably determined. Since Fe is the only element to show similar abundances in the neutral and ionized gas, it is possible that all elements indeed have similar abundances in both media. The underabundance of Ar would then be due to the fact that we used Ar\\1/H\\1\\ to estimate Ar/H, when (Ar\\1+Ar\\2)/H\\1\\ should be used instead."
    },
    "0608/astro-ph0608673_arXiv.txt": {
        "abstract": "We report here on the optimization of 0.5~cm thick pixelated Orbotech CZT detectors with regards to the best contacting materials and the use of steering grids. We evaluated the performance of different contacting materials. Our study differs from earlier ones in that we investigated the performance of different anode and cathode materials separately. We obtain the best performance with Au cathodes. For different anode materials Ti and In give the best energy resolutions. The detector (2.0$\\times$2.0$\\times$0.5~cm$^3$, 8$\\times$8 pixels) shows excellent 59 keV, 122 keV and 662 keV energy resolutions of 1.4 keV, 1.9 keV, and 7.4 keV, respectively. Furthermore, we report on using steering grids to improve on the performance of the pixelated detectors. Previously, the benefit of steering grids had been limited by additional electronic noise associated with currents between the negatively biased steering grids and the anode pixels. We are currently exploring the possibility to isolate the steering grid from the CZT substrates by a thin layer of Al$_2$O$_3$. We performed a series of measurements to determine by how much the isolation layer reduces the grid-pixel currents. Comparing the currents between two Au contacts before and after isolating one of the two contacts from the CZT with a 700~nm thick layer of Al$_2$O$_3$, we measure that the isolation layer reduces the currents by a factor of about 10 at 500~V. We present some results from a detector before and after deposition of an isolated steering grid. The grid indeed improves on the detectors energy resolution and detection efficiency. We show that simulations can be used to model the anode to cathode charge correlation in excellent agreement with the experimental results. ",
        "introduction": "\\label{sect:intro} Cadmium Zinc Telluride (CZT) has emerged as the detector material of choice for the detection of hard X-rays and soft gamma-rays with excellent position and energy resolution and without the need for cryogenic cooling. We report here on tests with CZT substrates from the company Orbotech Inc. \\cite{Orbotech} (formerly Imarad). They are grown with a Modified Horizontal Bridgeman (MHB) process. The MHB process results in a high yield of good crystals and in very uniform substrates. \\begin{figure}[bth] \\vspace{-0.4cm} \\centering \\includegraphics[width=8cm]{figure1.eps} \\vspace*{-0.6cm} \\caption{Potential distribution of a central pixel in a 0.5~cm thick detector with pixel pitch of 0.24 mm, pixel width of 0.16 mm and a steering grid width of 0.016 mm from a 3-D detector simulation. While the anode pixels were held at ground, the cathode was biased at -500~V and the steering grid at -300~V.} \\label{potential} \\end{figure} The bulk resistivity ($\\sim$ $10^{10}$ ohm cm) is lower than that of the more conventionally grown High-Pressure Bridgman CZT. Using high-workfunction cathode materials that form blocking Shottky contacts on the $n$-type CZT substrates, the bias currents can be reduced by more than one order of magnitude so that the associated shot noise no longer deteriorate the energy resolution \\cite{Vada,Henric04,Jung05}. Different cathode materials have been studied extensively. In this contribution, we give a brief report on results obtained with different anode materials. A full description of this study will be given elsewhere (Jung et al., in preparation).  Most current CZT detectors use ``electron-only'' detection strategies \\cite{Barret95,Luke95} to ameliorate the problems associated with the poor hole mobilities and short hole trapping times in CZT ($(\\mu_h\\:\\tau_h \\sim \\rm(0.2-5)\\cdot10^{-5}$cm$^2/V)$). One can use for example the ``small pixel effect'' \\cite{Barret95,Luke95} so that charge generated inside the detector induces pixel currents only briefly before impinging on the anode pixels. The anode signals become largely independent of the location of the charge generation and the ``depth of interaction'' (DOI). An additional DOI correction can be used to further improve the energy resolution. The DOI can be estimated from measuring the drift time of the electrons \\cite{Kalem02,Zhang04} and from measuring the anode to cathode charge ratio \\cite{Kraw04}. Some groups (see e.g.\\ \\cite{Zhang04}) use steering grids to improve the performance of pixelated detectors. Steering grids ameliorate the energy resolutions and detection efficiencies of pixelated detectors by steering electrons away from the areas between pixels towards the pixels \\cite{bolotnikov}. Furthermore, steering grids modify the weighting potential of the pixels. The result is a stronger small pixel effect and an additional enhancement of the energy resolution of the detectors. Most groups bias the steering grids at voltages of between -30 and -60~V relative to the anode pixels. Owing to the limited surface resistivity of CZT, large grid-pixel currents start to flow for higher grid bias voltages. However, simulations show that a grid bias between -100 V and -300 V is required to appreciably modify the electric field inside the detectors (see Fig.\\ \\ref{potential}). We are thus exploring the possibility to isolate the grids from the CZT substrates (Fig. \\ref{fig:Layer}). The isolation layer can reduce the grid-pixel currents while allowing the electric field to enter the detector. The second main topic of this contribution are studies of the effectiveness of Al$_2$O$_3$ isolation layers. Furthermore, we present some results from a detector before and after deposition of an isolated steering grid, and compare experimentally measured data with simulations. ",
        "conclusions": "In this report, we present the results from several activities to optimize Orbotech CZT detectors. We tested detectors with standard In anode contacts and different cathode contacts. All the cathode contacts reduced the leakage current by more than one order of magnitude compared to the original In cathode. Au, Cr, In and Ti cathodes gave very similar 662 keV energy resolutions. On a different CZT detector, we replaced both, anode and the cathode contacts with identical material. In this study, Au gave the best results, followed in order of performance by Pt, In, and Cr. We tested four different anode materials. Ti and In gave excellent almost identical results. For further improvements one can use steering grids. Our results show that one can use an $\\sim$700~nm thick Al$_2$O$_3$ isolation layer beneath the steering grid to reduce the current between grid and pixel. A first detector with isolated steering grid shows a significant improvement in energy resolution and detection efficiency. We plan to test other isolation materials and to continue to optimize the use of Al$_2$O$_3$. One can use isolation layers also for other CZT detectors designs such as Frish-grid and cross-strip detectors."
    },
    "0608/astro-ph0608634_arXiv.txt": {
        "abstract": "High resolution cosmological N-body simulations of four galaxy-scale dark matter halos are compared to corresponding N-body/hydrodynamical simulations containing dark matter, stars and gas. The simulations without baryons share features with others described in the literature in that the dark matter density slope continuously decreases towards the center, with a density $\\rho_{\\rm{DM}} \\propto r^{-1.3\\pm0.2}$, at about 1\\% of the virial radius for our Milky Way sized galaxies. The central cusps in the simulations which also contain baryons steepen significantly, to $\\rho_{\\rm{DM}} \\propto r^{-1.9\\pm0.2}$, with an indication of the inner logarithmic slope converging. Models of adiabatic contraction of dark matter halos due to the central build-up of stellar/gaseous galaxies are examined. The simplest and most commonly used model, by Blumenthal \\emph{et al.}, is shown to overestimate the central dark matter density considerably. A modified model proposed by Gnedin \\emph{et al.} is tested and it is shown that while it is a considerable improvement it is not perfect. Moreover it is found that the contraction parameters in their model not only depend on the orbital structure of the dark-matter--only halos but also on the stellar feedback prescription which is most relevant for the baryonic distribution. Implications for dark matter annihilation at the galactic center are discussed and it is found that although our simulations show a considerable reduced halo contraction as compared to the Blumenthal \\emph{et al.} model, the fluxes from dark matter annihilation is still expected to be enhanced by at least a factor of a hundred as compared to dark-matter--only halos. Finally, it is shown that while dark-matter--only halos are typically prolate, the dark matter halos containing baryons are mildly oblate with minor-to-major axis ratios of $c/a=0.73\\pm0.11$, with their flattening aligned with the central baryonic disks. ",
        "introduction": "\\label{sec:intro} There are still a multitude of open questions regarding the distribution of dark matter in galactic halos, the effect of baryons upon the structure of dark matter halos being one.  Many simulations of galactic halos only take into account the dark matter (e.g.\\ \\cite{astro-ph-9508025,astro-ph-0311231,astro-ph-9903164,astro-ph-9908159,astro-ph-0306203,astro-ph-0312544}) , but one would expect the baryonic component of the galaxy to behave very differently since it is able to cool (dissipate energy), and contract considerably.  This is indeed observed in simulations containing baryons and also in nature, where the baryons form a disk and/or bulge at the center of apparently much more extended dark matter halos.  It has long been realized that this ability of baryons to sink to the center of galaxies would create an enhanced gravitational potential well within which dark matter will congregate, increasing the dark matter density there. To model this effect it is common to use adiabatic invariants or some small modification of them \\cite{Eggen:1962dj,young,zeldovich,ASJOA.301.27,rydengunn,astro-ph-0406247,astro-ph-0501567,astro-ph-0507589,astro-ph-0601669,astro-ph-0604553}. Such models are frequently used and it is of particular interest to test the validity of these models \\cite{Barnes&White_84,astro-ph-0204164,astro-ph-0406247} currently because of recent advances in mapping the velocity field of the Milky Way \\cite{astro-ph-0605025} and other galaxies. Furthermore, upcoming gamma ray experiments will look for the flux due to the self-annihilation of weakly interacting dark matter candidates from, e.g., the galactic center.  Since this flux is proportional to the dark matter density squared, predictions for, and conclusions from, the data will depend strongly on the details of the effect of baryonic pinching upon dark matter halos.  In this work we aim to investigate the effects of baryons on dark matter halos in disk galaxies and test the most common models of adiabatic contraction by comparing recent cosmological N-body/hydrodynamical simulations of galaxies containing dark matter and baryons to results from N-body dark-matter--only simulations of the {\\it same} halos. Fully cosmological simulations starting at high redshifts are used. The simulation results are known from previous studies to produce overall realistic gas and star structures for spiral galaxies \\cite{astro-ph-0204366,astro-ph-0602595,astro-ph-0606531}, even though the numerical resolution is still far from being able to resolve any small scale features observed in real galaxies. The most important dynamical property in this paper is the creation of stable disk and bulge structures both for the gas and star components. The \\emph{angular momentum problem} is overcome by stellar feedback, implying that the matter is not too centrally concentrated; a generic problem of early galaxy simulations (see, e.g., \\cite{astro-ph-0602351} and references therein on forming disk galaxies in simulations). With the baryonic disks and bulges formed fully dynamically the surrounding dark matter halo response should also be realistically predicted inside the simulated spiral galaxies.  \\smallskip The paper is organized as follows. In section \\ref{sec:sim} we present the simulations, and then focus in section \\ref{sec:DM_profile}, in particular, upon profile fits to the density of the dark matter halos in the simulations with and without baryons. In section \\ref{sec:adiabatic} we investigate different prescriptions which aim to predict the effects of baryons upon dark matter profiles by testing whether they are able to reproduce the dark matter profiles observed in the simulations with baryons. We then comment in section \\ref{sec:WIMP} on how these results might change the expected flux of gamma rays from dark matter annihilation in spiral galaxies, including the Milky Way. Finally, we analyze in section \\ref{sec:non-sph} the nonsphericity of the different components of the simulations -- dark matter, gas and stars -- to attempt to further quantify their effects upon each other before we summarize our results in section \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We have presented results comparing the structure of galactic dark matter halos formed in N-body simulations including only dark matter to that of {\\it the same} halos formed in N-body/hydrodynamical simulations of galaxies, containing dark matter, stars and gas. From our selected high resolution galaxies three out of the four galaxies formed in the hydrodynamical simulations contain very distinct disks of gas and stars, and central stellar bulges, the fourth is strongly barred.  The central slope of the dark matter density profiles becomes significantly steeper when baryons are present, with the average logarithmic slope parameter $\\gamma=-\\textrm{d} \\log \\rho /\\textrm{d}\\log r$ increasing from 1.3$\\pm$0.2 to 1.9$\\pm$0.2 at about 1\\% of the virial radius.  The pinching of dark matter halos in response to the cooling and contraction of baryons was investigated further, to test adiabatic contraction models for the case where galaxies of realistic linear sizes and other properties are formed.  In relation to the orbital structure of dark-matter--only halos, the mean of the time averaged radius $\\langle\\bar{r}\\rangle$ of dark matter particles versus radius $r$ is very well described by a power law relation (specified by two parameters $A$ and $w$) as suggested by Gnedin \\emph{et al}. Moreover, it is found that the Gnedin \\emph{et al.}\\ \\cite{astro-ph-0406247} prescription for adiabatic contraction is much more successful at reproducing the density profile of dark matter in the simulations with baryons, than the standard scenario of Blumenthal \\emph{et al.} \\cite{ASJOA.301.27}, in which circular orbits are assumed.  However, the parameters of the Gnedin \\emph{et al.}\\ model, ($A$,$w$), which give the best fit for the baryonic contraction, are somewhat different from the ($A$,$w$) parameters describing the averaged orbital structure of the dark-matter--only halos. Given the ($A$,$w$) uncertainty estimates, described in the text, this difference appears to be significant. In addition, it is also found that the contraction reconstruction values of ($A$,$w$) also depend on the strength of the stellar feedback in simulations of otherwise identical halos, further indicating (perhaps as one might expect) that the effects of baryonic pinching are more complicated, than what can be captured in this two-parameter model.   Our results indicate that the amount of baryonic pinching of the dark matter halos are overestimated in earlier works applying the adiabatic compression model by Blumenthal \\emph{et al.}, at least for Milky Way sized disk galaxies. This has ramifications for predictions of the (putative) dark matter annihilation flux from the galactic center. It is found, that the flux can be reduced by several orders of magnitude, although baryonic contraction still boosts the signal significantly above the value one would expect on the basis of simulations containing only dark matter.  Finally, we have determined the triaxiality of the dark matter halos. Dark matter only halos were found to be significantly more prolate than halos containing baryons.  The influence from baryons flatten the dark matter into a slightly oblate halo ($c/a=0.73\\pm0.11$) aligned in the same plane as the stellar/gaseous disk.  Moreover, in the simulations containing baryons, galactic disks tend to form in the planes aligned with the flattest parts of the dark matter halos formed in the corresponding simulation without baryons.  {\\bf Acknowledgments} We are extremely grateful for conversations with Lars Bergstr\\\"om, Anne Green, Anatoly Klypin, Edvard M\\\"ortsell, George Rhee and Christian Walck. All computations reported in this paper were performed on the IBM SP4 and SGI Itanium II facilities provided by Danish Center for Scientific Computing (DCSC). MF receives funding from the Swedish Research Council (Vetenskapsr\\aa det). The Dark Cosmology Centre is funded by the DNRF.  \\bigskip"
    },
    "0608/astro-ph0608128_arXiv.txt": {
        "abstract": "{ROTSE-IIId observations of the Be/X-ray transient system KS 1947+300 obtained between September 2004 and December 2005 make it possible to study the correlation between optical and X-ray activity. The optical outburst of 0.1 mag was accompanied by an increase in X-ray flux in 2004 observations. Strong correlation between the optical and X-ray light curves suggests that neutron star directly accretes from the outflowing  material of Be star. The nearly zero time lag between X-ray and optical light curves suggests a heating of the disk of Be star by X-rays. No optical brightening and X-ray enhancement was seen in 2005 observations. There is no indication of the orbital modulation in the optical light curve. ",
        "introduction": "The transient X-ray source KS 1947+300 was discovered in 1989 through observations of the GC 2023+338 field (Borozdin et al. 1990). After its discovery Grankin et al. (1991) and Goranskij et al. (1991) identified an optical counterpart which shows H$\\alpha$ emission line in its spectrum. Negueruela et al. (2003) derived a spectral type B0Ve for the optical companion and a distance of $\\sim$ 10 kpc. Considering this distance, the peak X-ray luminosity during the 2000-01 outburst was $\\sim 2\\,10^{37}$ erg/s (typical for Be/X-ray transients during the Type II outburst). The equivalent width of H$\\alpha$ line remained almost the same during their one year observation period. They mentioned that the occurrence of several weaker X-ray outbursts has not been reflected in any change in the intensity of H$\\alpha$ line. They explained this fact with a low inclination angle of the system. Galloway et al. (2004), analyzing their observations obtained with the instruments aboard RXTE (the ASM, PCA, HEXTE), found that neutron star revolves in a nearly circular orbit (e = 0.033) with orbital period 40.415 d and projected semimajor axis 137\\,lt-s, based on the modulations of neutron star's pulse period of $18.7\\,\\mathrm{s}$ (Chakrabarty et al. 1995) during the 2000-01 outburst. They reported an increase in neutron star period indicating that the angular momentum was effectively transferred from the accreted material on to the neutron star which required an extended disk. They did not find eclipses in their data. Tsygankov$\\&$Lutovinov (2005) estimated the magnetic field strength of the pulsar as $2.5\\,10^{13}$ G and distance to the binary as $9.5\\pm1.1$ kpc. This system occupies the region of Be/X systems in the P$_{spin}$- P$_{orb}$ diagram (Reig et al. 2004). ",
        "conclusions": "Start time of our observations (MJD 53255) corresponds to the orbital phase of 0.42 which is calculated using T$_{\\pi/2}$ = 51985.31 MJD from Galloway et al. (2004) (T$_{\\pi/2}$ is the epoch at which the mean longitude is $\\pi$/2 ). Maximum of brightness occurs at $\\sim$ 0.25 orbital phase. Magnitude  change lasts about 70-80 d. The brightening and fading phases have similar time durations. If the magnitude change is attributed to the disk around Be star then the disk is lost in about 70-80 d. Hanuschik et al. (1993), in their study of H$\\alpha$ outburst of $\\mu$ Cen, propose that the outburst is caused by the ejection of photospheric matter into bound orbits forming a short lived (less than 100 d) quasi stable Keplerian disk. They did not find close relation between H$\\alpha$ emission activity and photometric data. Negueruela et al. (2003) observed almost no change in the strength of the H$\\alpha$ line (equivalent width:$\\sim$ 15 A) during their one year observation period which begins one month after the 2000-01 X-ray outburst. They argued that if Be star disk was lost then the strength of H$\\alpha$ line should fade away. It seems that the long period of X-ray activity in 2000-01 period and likewise remaining relatively constant the strength of H$\\alpha$ line oppose to the possibility of a short lived disk during this period. It is less likely to have short lived disk in this system.  The circular orbital velocity of X-ray component is 74 km/s according to the orbital parameters given by Galloway et al (2004). Using the orbital parameters of KS 1947+300, Galloway  et al., (2004) estimated the mass function as f$_{x}$= 1.71 M$_\\odot$.   For a neutron star of mass 1.4 M$_\\odot$, the mass of optical component is 3.4 M$_\\odot$ when the inclination angle i=90$^{o}$. Since we did not see eclipses either from RXTE/ASM or ROTSE observations, we assumed that inclination angles should be less than 85$^{o}$. If we assume masses of 15 and 20 M$_\\odot$ for the optical component the inclination angles are found as 28$^{o}$ and 26$^{o}$, respectively. The corresponding Roche Lobe radii are 73 and 83 R$_\\odot$ which are much larger than the expected radius of B0V star. It is difficult to explain the X-ray outburst of this system by mass transfer from Be star through the inner Lagrangian point of Roche Lobe. In the truncated viscous disk model (Okazaki$\\&$Negueruela 2001) the tidal torques on neutron star truncates the disk effectively in systems with low eccentricity therefore preventing the development of an extended and steady disk. It is difficult to explain the behavior of Be disk in this system with the viscous disk model.  The maximum X-ray luminosity obtained from ASM data (Fig.1) during our 2004 observation period is $\\sim$ $5\\,10^{36}$ erg/s which is calculated using the definition in Galloway et al. (2004) (1 count/s $\\sim$$3.6\\,10^{-10}$ erg/s/cm$^{2}$ in the energy range 1.5-12 keV). This luminosity is typical for TypeI outburst but its duration is longer than the typical value for TypeI outburst. The derived accretion rate on to the neutron star is calculated as $\\sim$ $4\\,10^{-10}$ M$_\\odot$/yr. It is known that the mass loss rate derived is in the order of $10^{-8}$ M$_\\odot$/yr for Be stars (Waters et al. 1987, Clark et al. 1999). Clark et al. (1999) find a mass loss rate of 1.5-9$\\,10^{-8}$ M$_\\odot$/yr for the Be/X-ray binary A0535+26. Hanuschik et al. (1993) calculated the typical mass transfer rate during the H$\\alpha$ outburst of $\\mu$ Cen (B2IV-Ve star) which lasted a few days as $\\sim$ $4\\,10^{-9}$ M$_\\odot$/yr. Such equatorial mass loss rates from Be star can produce the observed X-ray luminosity for KS 1947+300.  The increase in optical luminosity of the system is calculated $\\sim$ $10^{36}$ erg/s using the value of M$_{v}$ =-4.2 (Vacca et al. 1996). This increase can be explained as a heating of Be star equatorial disk by X-rays. The nearly zero time lag between X-ray and optical light curves suggests this result. The X-ray emission from the neutron star is absorbed in the stellar wind and equatorial disk of the Be star and reprocessed into the optical wavelengths (Kriss et al. 1983, Mendelson and Mazeh, 1991)."
    },
    "0608/astro-ph0608402_arXiv.txt": {
        "abstract": "The origin of hydrodynamic turbulence, and in particular of an anomalously enhanced angular momentum transport, in accretion disks is still an unsolved problem. This is especially important for cold disk systems which are practically neutral in charge and therefore turbulence can not be of magnetohydrodynamic origin. While the flow must exhibit some instability and then turbulence in support of the transfer of mass inward and angular momentum outward, according to the linear perturbation theory, in absence of magnetohydrodynamic effects, it should always be stable. We demonstrate that the three-dimensional secondary disturbance to the primarily perturbed disk, consisting of elliptical vortices, gives significantly large hydrodynamic growth in such a system and hence may suggest a transition to an ultimately turbulent state. This result is essentially applicable to accretion disks around quiescent cataclysmic variables, in proto-planetary and star-forming disks, the outer region of disks in active galactic nuclei, where the gas is significantly cold and thus the magnetic Reynolds number is smaller than $10^4$. ",
        "introduction": "Origin of turbulence and then the viscosity parameter $\\alpha$, while has been demonstrated for a hot Keplerian accretion disk by the magnetorotational instability (MRI; Balbus \\& Hawley 1991; Balbus, Hawley, \\& Stone 1996; Hawley, Balbus, \\& Winters 1999), is still not well understood for cold disks, e.g. accretion disks around quiescent cataclysmic variables (Gammie \\& Menou 1998; Menou 2000), proto-planetary and star-forming disks (Blaes \\& Balbus 1994), and the outer region of disks in active galactic nuclei (Menou \\& Quataert 2001), even after more than three decades of the famous discovery of the $\\alpha$-disk (Shakura \\& Sunyaev 1973). These cold systems are largely neutral in charge so that the MRI driven turbulence appears to be ruled out. However, to support accretion, there must appear some sort of turbulence (most likely pure hydrodynamic in nature) and then the corresponding turbulent viscosity, as the molecular viscosity is negligible. Richard \\& Zahn (1999) indeed showed by analyzing results of laboratory experiments for the classical Couette-Taylor flow that in the case where angular momentum increases outward, similar to the Keplerian flow, hydrodynamic turbulence may be sustained. Then they derived the corresponding turbulent viscosity which is very useful for the astrophysical purpose. Longaretti (2002) argued that the Keplerian accretion flow in the shearing sheet limit should be turbulent and the lack of turbulence in simulations is due to lack of their resolution. Then Lesur \\& Longaretti (2005) showed that the required resolution is not possible to achieve with the presently available computer resources. They also found that the efficiency of turbulent transport is directly correlated to the critical Reynolds number for the transition to turbulence. Inspite of all these analyzes, the actual origin of hydrodynamic turbulence in such systems still remains unclear. It is important to note that the numerical simulation has never been carried out with a {\\it very high} Reynolds number what the real accretion disk exhibits. Therefore, the hydrodynamic effects may compete with the magnetohydrodynamic effects (essentially MRI) at a high Reynolds number and, in a realistic system, turbulence may be due to hydrodynamic effects independent of whether the disk is cold or hot.  Recently, several authors including ourselves have put their efforts in achieving progress toward the solution of this difficult problem (e.g. Tevzadze et al. 2003; Chagelishvili 2003; Umurhan \\& Regev 2004; Yecko 2004; Mukhopadhyay, Afshordi \\& Narayan 2005; Afshordi, Mukhopadhyay \\& Narayan 2005; Johnson \\& Gammie 2005; Umurhan et al. 2006; Barranco \\& Marcus 2005; Papaloizou 2005). The main aim of these works is to demonstrate the {\\it pure} hydrodynamic turbulence by transient growth of energy with the suitable choice of two-dimensional initial perturbation. The idea is any large growth plausibly could switch the accretion disk into the nonlinear regime which might result its subcritical transition to turbulence if energy growth exceeds the threshold for turbulence. One might argue that transient growth, even if large, can not provide much clue on the existence and properties of a turbulent basin of attraction. Schmid \\& Henningson (2001) and Criminale, Jackson \\& Joslin (2003) have described in detail that transition to turbulence is not a unique process but it depends on the initial condition/disturbance and the nature of the flow. In fact, it is known that even in presence of secondary instability linearly unstable base flows may reach to a non-turbulent saturated state. However, turbulence definitely belongs to the nonlinear regime and it is exhibited only in the situations when large growth (or more precisely transient growth for the present purpose) switches the system over the non-linear regime. As our present goal is to understand the possible origin of hydrodynamic turbulence, we consider those situations when large transient growth governs non-linearity. At the present case where the accretion time scale is comparable (or shorter) to the time scale to achieve maximum growth (MAN), decaying nature of growth at a period longer than the accretion time scale does not matter. If growth (or maximum growth) becomes large enough that exceeds the threshold to trigger non-linearity and turbulence, then it does not matter whether growth evolves exponentially or transiently.  Dauchot \\& Manneville (1997) argued with a toy model that transient growth does not guaranty the transition to turbulence as the underlying phase portraits do not get deformed much in absence of any linearly unstable mode. However, that does not rule out the importance of transient growth as their model is very simplistic consisting of two variables only and growth is tinny to expect any deformation in the phase portraits. Other alternative methods were proposed to describe the subcritical transition to turbulence and then to investigate various underlying aspects in detail by e.g. Waleffe (1995, 1997), Brosa \\& Grossmann (1999), Waleffe (2003), Kerswell (2005), while last two works, that discuss non-rotating Couette flows and pipe flows respectively, are mostly relevant for the astrophysical purpose. Although the works bring some new insight into the subject, the main problem i.e. to understand the subcritical turbulence in three-dimension remains unsolved. This is indeed a non-trivial problem not only in astrophysics but also in fluid dynamics. In the present paper, we suggest a possible mechanism to govern non-linear effect into a Keplerian accretion disk. While our prescription does not solve the problem completely, it certainly opens up a new avenue to understand the physics behind this puzzle.  The problem with transient growth, what has been proposed as the mechanism to generate turbulence (e.g. Umurhan \\& Regev 2004; hereafter UR; Mukhopadhyay, Afshordi \\& Narayan 2005, Afshordi, Mukhopadhyay \\& Narayan 2005; hereafter MAN), is that in two-dimension underlying perturbations must ultimately decline to zero in the viscous flow. To overcome this limitation, it is necessary to invoke three-dimensional effects.  Various kinds of secondary instability, such as the elliptical instability, are widely discussed as a possible route to self-sustained turbulence in linearly perturbed shear flows (see, e.g. Pierrehumbert 1986; Bayly 1986; Hellberg \\& Orszag 1988; Craik 1989; Le Dize\\'s, Rossi \\& Moffatt 1996; Kerswell 2002). Therefore, we are motivated to see whether these three-dimensional instabilities are present in the Keplerian flows which consist of elliptical streamlines under two-dimensional perturbation. Goodman (1993) first explored the possible role of the elliptical instability in an accretion disk and Lubow, Pringle \\& Kerswell (1993) and Ryu \\& Goodman (1994) showed that angular momentum may be transferred from the disk to the tidal source by the instability effect. They essentially considered the case of forced disks. Later, Ioannou \\& Kakouris (2001) also examined transient growth being constantly re-excited by external noise.  The three-dimensional instability of a two-dimensional flow with elliptical vortices has been demonstrated by number of authors (e.g. Craik \\& Criminale 1986; Waleffe 1990) and has been proposed as a generic mechanism for the breakdown of many two-dimensional high Reynolds number flows. This result motivates us to investigate whether the similar scenarios exist in a Keplerian accretion disk. Therefore, essentially we plan to investigate the Keplerian flow with two consecutive perturbation effects. The primary two-dimensional perturbation, which drives transient growth (MAN), governs elliptical streamlines in the flow. Then we consider the further three-dimensional perturbation, that is the secondary one, to this two-dimensional flow, if that drives three-dimensional instability effect in presence of viscosity. Presumably, three-dimensional instabilities lead to nonlinear feedback and self-sustained turbulence.  Although the maximum growth in the Keplerian disk with finite vertical thickness and then with finite vertical perturbation is smaller compared to that in a two-dimensional Keplerian disk, as Tevzadze et al. (2003) and MAN argued, vertical stratification may cause to have a nonvanishing asymptotic growth that is a fraction of the maximum growth. However, they could not show whether this will significantly help the onset of turbulence. By three-dimensional hydrodynamic simulations, Barranco \\& Marcus (2005) studied the dynamics and formation of vortices in stably stratified proto-planetary disks and found that vortices are unstable under perturbation. Theoretically, the origin of these vortices is understood as perturbation in plane shear flow (UR, MAN). Continuing in the same line of thought, recently, Umurhan et al. (2006) have shown substantial growth resulting from initial perturbations in the linearly stable steady state and have concluded that significant perturbation energy amplification occurs in accretion disks on a global scale.  In the present paper, we perform a local linear analysis in the shearing box approximation. Although a small section of an accretion disk may not reproduce the global disk properties, we assume that if turbulence is exhibited at any section, then that sustains and eventually affects the entire disk. While the local result does not guaranty its global signature, absence of local instability and then possible turbulence, perhaps, does guaranty its global absence. Therefore, from the results of this paper we can determine the future avenues of the subject, whether the analyzes for secondary instability including transient growth is a fruitful path to solve the problem.  In the next section, we outline the basic model considered for this problem describing primary and secondary perturbation. In \\S3, we present the simple solution when the evolution of secondary perturbation is much rapid than that of primary one. Subsequently, in \\S4 we analyze the general solutions when both perturbations vary simultaneously. Finally we sum up the results with a discussion in \\S5. ",
        "conclusions": "We find that significant energy growth of perturbation is possible to govern in shear flow with the Coriolis force. This system is an idealized local analog of an accretion disk which, under secondary perturbation, definitely exhibits three-dimensional large growth of transient kind and, in addition, sometimes exhibits unbounded late-time growth at a large amplitude of primary perturbation. We have explicitly demonstrated the perturbation effects one by one. First, primary two-dimensional perturbation induces vortex into the flow that can be locally seen as elliptical streamlines. This system, which does not have any exponential growing eigenmode but does exhibit significant transient growth, has been extensively studied already (e.g. UR; Yecko 2004; MAN). In this situation, a plane wave perturbation, that is frozen into the fluid, is sheared along with the background flow. At $t=0$, the effective wave vector of perturbation is in the $x$ direction ($k_x>>k_y$) and is negative which provides very asymmetric leading waves. Therefore, the flow at this stage is highly eccentric.  As time goes on, the wavefronts are straightened out by the shear and $|k_x|$ decreases and transient growth increases. When $k_x\\sim 0$, i.e. the wavefronts become almost radial, transient growth is maximum. At this time, the streamlines of the flow are almost circular. At yet later time, growth decreases and the wave becomes trailing. Then it has been argued that if the maximum growth exceeds the threshold for inducing turbulence, then this mechanism could drive the system to a turbulent state. Presumably, once the system becomes turbulent, it remains turbulent as a result of nonlinear interactions and feedback among the perturbations. It can be reminded that our present aim is to understand and establish the origin of viscosity in the flow that must be due to turbulence. The transfer of mass inward and angular momentum outward in an accretion disk is difficult to explain in absence of turbulence. However, the accretion disk is quite a complex system with, possibly, continuous perturbed flow. Therefore, if we can address a mechanism to govern turbulence, then its recycling is not a difficult job.  Second, we consider further perturbation, namely secondary perturbation, into the flow described above. The primarily perturbed shear flow serves as a background for secondary perturbation whose eccentricity naturally varies with time due to the evolution of primary perturbation. In this paper, we have especially demonstrated the evolution of this secondary perturbation which exhibits three-dimensional large transient growth in a local Keplerian accretion disk. While primary perturbation itself can produce large transient growth at a high Reynolds number that might drive the non-linear effects into the system, the best perturbation responsible for this effect is two-dimensional. However, it is understood that perturbations must ultimately decline to zero in presence of viscosity (see e.g. UR), unless three-dimensional effects are invoked. Therefore, we have addressed the possible origin of three-dimensional effects that shows a clear route to three-dimensional hydrodynamic growth and then possible non-linear feedback and turbulence in accretion flows. Underlying growth arises due to the elliptical vortices present in the background, rather than due to the plane shear which exhibited growth under primary perturbation.  In the standard fluid literature, the elliptical instability has been widely discussed as a possible route to self-sustained turbulence in linearly perturbed shear flows, as mentioned in \\S1. However, usual emphasis of those investigations is on non-rotating flows. Craik (1989), while discussed about the elliptical instability in rotating shear flows, did not focus on a Keplerian flow, what is of astrophysical interest. Therefore, in the present paper, we have first discussed the growth rate in standard elliptical flows and compared the results for non-rotating flows with that of rotating ones in \\S3.1. We have shown that the growth rate in a Keplerian flow with constant elliptical streamlines is significantly large compared to that in a non-rotating flow, particularly at high eccentricity.  However, in reality, when a small section of a Keplerian accretion disk is considered under a two-dimensional linear perturbation, the flow governs distorted elliptical streamlines whose structure vary with time. Therefore, the growth rate due to secondary perturbation at a fixed $\\epsilon$ (instantaneous growth rate) as described in \\S3.2 decreases much compared to that in the flow with idealized elliptical streamlines, unless the amplitude of primary perturbation, $\\zeta$, is very large. At a $\\zeta>1/3$, a pure vertical secondary perturbation produces the best growing eigenmode. On the other hand, at a $\\zeta\\le 1/3$, the best growing eigenmode arises due to other three-dimensional perturbations with a significant, but not sole, vertical effect. Although the instantaneous growth rate appears to be small for a small $\\zeta$ (which is of particular interest), at least compared to the case with idealized elliptical streamlines, actual growth which is the result of simultaneous evolution of both perturbations as described in \\S4 can be large enough to exhibit non-linear effects if the time-scale for the evolution of perturbation is large. The time for the evolution of perturbation scales with the Reynolds number of the flow as $R^{1/3}$ (MAN). As perturbation evolves, $k_x$ varies from $-\\infty$ to $0$ and thus the eccentricity of the flow decreases from $1$ to $0$. Most of the important three-dimensional growing modes are generated at the high eccentricity regime when $0.995\\le\\epsilon \\le 1$ and therefore $10\\le |k_x|\\le \\infty$. Important fact to note is that growth maximizes for three-dimensional perturbation with a significant vertical effect.  We therefore conclude with an important caveat. UR already showed via two-dimensional simulations that chaotic motions can persist for a time much longer than the time scale needed for linear growth. However, the corresponding vorticity decays unless the vertical structure is there. In the present paper, we have shown the existence of three-dimensional perturbation effects and corresponding eigenmodes which governs large energy growth and then suggests the possible existence of non-linear effects and self-sustained hydrodynamic turbulence in accretion disks. Now one will have to verify our suggestion by numerical simulation with proper resolution and possibly also by nonlinear analytic asymptotic methods."
    },
    "0608/astro-ph0608258_arXiv.txt": {
        "abstract": "% The recent development of numerical schemes for Relativistic MHD (RMHD) allows us to model the acceleration and outflow properties of winds from compact sources. Theoretical models suggest that acceleration and collimation of the flow are extremely inefficient when the speed is close to $c$, in contrast with many observations. Numerical results for an axisymmetric rotator, both in the case of monopolar and dipolar magnetic field will be presented, suggesting that in ideal RMHD acceleration is indeed inefficient. I will also mention numerical challenges and stability problems of present simulations. Finally I will discuss the interaction of a pulsar wind with the surrounding SNR. I will present emission maps based on numerical simulations of the flow inside the nebula and how they can be used in comparison with observations to derive informations about the properties of pulsar winds far away from the source. ",
        "introduction": "%  Magnetically dominated outflows from stars and accretion discs, are central to the  evolution and properties of these objects. Because they can efficiently extract rotational energy, and convert it into wind kinetic energy, magnetic outflows are ubiquitous in powering a variety of astrophysical systems. Rotationally powered pulsars are among the best example of relativistic outflows produced by a fast compact rotator. The high magnetic field at the surface ($B\\sim 10^{10-13}$ G) and short rotation period ($P\\sim 0.001-1$ s), induce an electric field at the surface itself which is strong enough to lift electrons and ions and accelerate them to relativistic speeds. These accelerated particles, as they move along curved magnetic field lines, produce the pulsed emission that is observed. In the magnetosphere of the neutron star, these particles produce pairs. If the charge density exceeds the Goldreich-Julian value then conditions can be modeled in term of ideal RMHD.  The typical Lorentz factor of the plasma in the pulsar magnetosphere, within the light cylinder, is $\\gamma\\sim 100$. However from observations of the interaction of the pulsar wind with the environment, and especially the surrounding SNR, we know that far away from the source the wind is ultrarelativistic ($\\gamma\\sim 10^{4-6}$) and weakly magnetized \\citep{kc84a,kc84b,arons04}. The basic idea is that the wind is magneto-centrifugally accelerated, and that the high Lorentz factor is achieved by almost completely converting magnetic energy into kinetic energy.  In this sense, pulsars constitute the prototype of relativistic accelerators and the basic physics of magneto-centrifugally driven relativistic winds apply to a wide variety of astrophysical objects from accretion disks to GRBs and microquasars. In all of these systems a plasma is assumed at the base with small outflow velocity (compared to the asymptotic velocity), embedded in a strong magnetic field. Magnetorotation then accelerates the plasma and converts the magnetic energy into relativistic motion. The study of winds from neutron stars can give an useful insight into the general problem of relativistic outflows.  In the simple 1D radial model for relativistic winds from compact rotators by \\citet{michel69}, solutions can be parametrized in term of $\\sigma=\\Omega^2\\Phi^2/\\dot{M}$, where $\\Omega$ is the rotation rate, $\\Phi$ the magnetic flux and $\\dot{M}$ the mass flux. $\\sigma$ is the maximum Lorentz factor the wind can achieve if all magnetic energy is converted into kinetic energy, and at the base of a pulsar it is $\\sim 10^6$. The results by \\citet{michel69} show that, due to relativistic effects, the asymptotic Lorentz factor tends to $\\sigma^{1/3}$, and conversion of magnetic to kinetic energy is minimal. This is known as the $\\sigma-\\gamma$ problem \\citep{arons04}.  The low efficiency of the Michel model is due to the assumption of strictly radial flow. If this assumption is relaxed and the flow is allowed to diverge more than radially from the Light Cylinder up to the SNR, then acceleration is more efficient \\citep{begelman94}. It is known from the theory of classical winds that the presence of a magnetic field can cause collimation of the flow along the rotation axis, and divergence along the equator, so in principle one might expect magnetic collimation to be also responsible for the acceleration of the wind. However when the flow speed reaches values close to $c$ the Coulomb term in the electromagnetic force cannot be neglected (as in non relativistic MHD). The strength of the induced electric field is comparable with the magnetic field and the Coulomb term balances the magnetic hoop stress. Solutions of the relativistic Grad-Shafranov equation for a monopolar field in the trans-fast regime \\citep{beskin98} have shown that acceleration is too slow and the asymptotic results of the Michel model apply also in 2D axisymmetric geometry, meaning that the flow remain radial.  All simulations were performed using the code developed by \\citet{ldz03}. It is a high-order finite difference UCT scheme, that solves the equation of relativistic MHD in general curvilinear coordinates. The scheme is based on an approximate HLL solver, that avoids wave decomposition. ",
        "conclusions": "Recent numerical results have confirmed theoretical predictions about the acceleration of relativistic pulsar winds. The conversion of magnetic to kinetic energy appears to be minimal, the flow closely resembles the exact monopole solution, there is no evidence for the formation of a collimated energetic jet. The use of a more realistic dipolar field does not seem to change significantly the wind properties far from the light cylinder. On the contrary the energy and angular momentum losses seem to depend only on the integrated value of $\\sigma$, and for $\\sigma>1$ on the amount of open magnetic flux.  An alternative way to investigate the wind properties at least at large distances comes from the study of its interaction with the surrounding SNR. Modeling of the ``jet-torus'' structure in PWNe have shown that  such features are due to the details of the nebular flow, which in turn depends on the wind properties. Simulations show that an energy and magnetic field distribution in the wind resembling the monopole solution, can indeed explain the main details of the observations."
    },
    "0608/cond-mat0608023_arXiv.txt": {
        "abstract": "The nonrelativistic standard model for a continuous, one-parameter diffusion process in position space is the Wiener process. As well-known, the Gaussian transition probability density function (PDF) of this process is in conflict with special relativity, as it permits particles to propagate faster than the speed of light. A frequently considered alternative is provided by the telegraph equation, whose solutions avoid superluminal propagation speeds but suffer from singular (non-continuous) diffusion fronts on the light cone, which are unlikely to exist for massive particles. It is therefore advisable to explore other alternatives as well. In this paper, a generalized Wiener process is proposed that is continuous, avoids superluminal propagation, and reduces to the standard Wiener process in the non-relativistic limit. The corresponding relativistic diffusion propagator is obtained directly from the nonrelativistic Wiener propagator, by rewriting the latter in terms of an integral over actions. The resulting relativistic process is non-Markovian, in accordance with the known fact that nontrivial continuous, relativistic Markov processes in position space cannot exist. Hence, the proposed process defines a consistent relativistic diffusion model for massive particles and provides a viable alternative to the solutions of the telegraph equation. ",
        "introduction": "\\label{introduction}  It is well-known for a long time that the nonrelativistic diffusion equation (DE)~\\cite{Be67,Gardiner,HaTo82,Roepstorff} \\be\\label{e:diffusion} \\f{\\p}{\\p t} \\rho(t,x)=D\\,\\f{\\p^2}{\\p x^2} \\rho(t,x),\\qquad t\\ge 0,\\; D>0 \\ee is in conflict with the postulates of special relativity. A simple way to see this is to consider the propagator of Eq.~\\eqref{e:diffusion}, which for $d=1$ space dimensions, is given by \\be\\label{e:solution_non-rel} p(t,x|x_0)= \\f{\\Theta(t)}{\\left(4\\pi Dt\\right)^{1/2}}\\, \\exp\\biggl[-\\f{(x-x_0)^2}{4D t}\\biggr] \\ee [Heaviside's theta-function is defined by \\mbox{$\\Gt(t)=1,t\\ge 0$} and \\mbox{$\\Gt(t)=0,t< 0$}]. This propagator~\\eqref{e:solution_non-rel} represents the solution of Eq.~\\eqref{e:diffusion} for the initial condition \\mbox{$\\rho(0,x)=\\gd(x-x_0)$}. That is, if $X(t)$ denotes the path of a particle with fixed initial position $X(0)=x_0$, then $p(t,x|x_0)\\diff x$ gives the probability that the particle is found in the infinitesimal volume element $[x,x+\\diff x]$ at time~$t>0$. As evident from Eq.~\\eqref{e:solution_non-rel}, for each $t>0$ there is a small, but non-vanishing probability that the particle may be found at distances $|x-x_0|>ct$, where $c$ is the speed of light. \\par This problem has attracted considerable interest during the past 100 years~(see, e.g., \\cite{Go50,Ru57,Sc61,Ha65,Du65,MaWe96,1997DeMaRi,1998DeRi,BoPoMa99,HePa01,KoLi00,KoLi01,AbGa04,AbGa05,Ra05,OrHo05,DuHa05a,DuHa07,DuHa06b,2004De,Ko05,FrLJ06} and references therein). Nonetheless, it seems fair to say that a commonly accepted solution is still outstanding. Apart from the profound theoretical challenge of developing a consistent relativistic diffusion theory, several practical applications exist, e.g., the analysis of data from high energy collision experiments~\\cite{Wo04,AbGa04,AbGa05,2005RaGrHe,2006PhRvC..73c4913V}. In this context, an often considered alternative to Eq.~\\eqref{e:diffusion} is given by the telegraph equation~(TE)~\\cite{Ta22,Go50,RevModPhys.61.41,RevModPhys.62.375,MaWe96,KoLi00,KoLi01,Ko05,AbGa04,AbGa05} \\be\\label{e:telegraph} \\tau_d \\f{\\p^2}{\\p t^2}\\rho + \\f{\\p}{\\p t} \\rho=D\\,\\nabla^2 \\rho, \\ee where $\\tau_d$ is an additional relaxation time parameter.  As e.g. discussed by Masoliver and Weiss \\cite{MaWe96} or Abdel-Aziz and Gavin~\\cite{AbGa04}, the diffusion fronts described by Eq.~\\eqref{e:telegraph} spread at finite absolute speed $c_d=\\pm (D/\\tau_d)^{1/2}$; cf. Fig.~1 of Ref.~\\cite{AbGa04}. In particular, by fixing $c_d=c$, the solution of Eq.~\\eqref{e:telegraph} defines a simple   relativistic, \\emph{non-Markovian}~\\footnote{Because of the second order time derivative, the TE~\\eqref{e:telegraph} describes a \\emph{non-Markovian} diffusion process. A general non-existence theorem for Lorentz-invariant Markov processes in position space can be found in Dudley's seminal paper~\\cite{Du65}.} diffusion process without superluminal propagation. The \\lq nonrelativistic limit\\rq\\space corresponds to letting $c_d\\to \\infty$ in Eq.~\\eqref{e:telegraph}, which gives $\\tau_d\\to 0$ and thus leads back to Eq.~\\eqref{e:diffusion}. \\par However, apart from these appealing properties the solutions of Eq.~\\eqref{e:telegraph} exhibit a peculiar feature that is rather unexpected from a simple diffusion theory: The diffusion fronts are given by running $\\gd$-peaks whose intensity decreases as time grows; cf.~Eq.~(32) in Ref.~\\cite{MaWe96}. The appearance of such singularities can be understood by the fact that the one-dimensional TE~\\eqref{e:telegraph} can be derived from a so-called \\emph{persistent} random walk (RW) model~\\cite{Fu17,Fu22,Ta22,Go50,BoPoMa98}. This model assumes that a particle moves with constant absolute velocity between neighboring lattice points. At each lattice point, the particle is either back-scattered or transmitted, with the transmission probability being larger than the back-scattering probability (\\emph{persistence}). As Goldstein~\\cite{Go50} showed for the one-dimensional case, the continuum limit of this model leads to the TE (see Bogu\\~n\\'a et al.~\\cite{BoPoMa98} for a generalization to higher space dimensions). \\par Hence, in contrast to the ordinary DE~\\eqref{e:diffusion}, the TE~\\eqref{e:telegraph} relies on asymmetric transition probabilities, which among others lead to a non-vanishing probability concentration at the diffusion fronts. From the practical point of view, the concept of persistent diffusion can certainly be useful in some situations, e.g., if one aims at describing the propagation of photons in thin slabs or foams~\\cite{BoPoMa99,miri:031102, miri:031111,miri:031115}. In general, however, it seems unlikely that the relativistic generalization of the simple nonrelativistic diffusion theory would require conceptual modifications as severe as the introduction of persistence. In particular, for massive particles $\\gd$-peaked diffusion fronts are unlikely to occur in nature, because such fronts would imply that a finite fraction of particles carries a huge amount of kinetic energy (much larger than $mc^2$). One may, therefore, wonder if it is possible to construct alternative, non-persistent diffusion processes, which are consistent with the basic requirements of special relativity, and converge to the solutions of the nonrelativistic DE~\\eqref{e:diffusion} in a suitable limit case. This is the question the present paper aims to deal with. \\par Before outlining our approach, a general remark might be in order. Usually, a diffusion theory intends to provide a simplified phenomenological description for the complex stochastic motion of a particle in a background medium (e.g., on a substrate~\\cite{Fu17,Fu22,Ta22,Go50,BoPoMa98} or in a heat bath~\\cite{DuHa07}). Thus, there exists a \\emph{preferred} frame, corresponding to the rest frame of the substrate (or, more generally, the center-of-mass frame of the interaction sources causing the stochastic motion). It is therefore not expedient to look for Lorentz or Poincar\\'e invariant spatial diffusion processes (cf. Sec.~5 of Montesinos and Rovelli~\\cite{MoRo01}). Accordingly, we focus here on discussing simple diffusion models that comply with the basic requirements of special relativity in the rest frame of the substrate. \\par Very often, diffusion models are based on time evolution equations like the DE~\\eqref{e:diffusion}, derived either from an underlying microscopic model or as an approximation to a more complicated phenomenological model~\\cite{Roepstorff,Gardiner,LaRoTi82,HaTaBo90,Kramers,Be67}. Unfortunately, apart from the TE~\\eqref{e:telegraph}, the construction of relativistic DEs turns out to be very cumbersome. Among other things, this can be attributed to the fact that standard limiting procedures, as those leading  from discrete RW models~\\cite{Gardiner} to Eq.~\\eqref{e:diffusion}, typically result in superluminal propagation (cf. the discussion in Sec.~\\ref{massive_model} below). Therefore,  a different approach is pursued in the present paper: Instead of aiming at a relativistic evolution equation in the first place, we try to obtain a continuous, relativistic generalization of Eq.~\\eqref{e:solution_non-rel} by focussing on the structure of diffusion propagators. \\par The paper is organized as follows. First, some basic properties of Markovian diffusion models for massive particles are reviewed (Sec.~\\ref{massive_model}). In particular, it will be demonstrated in detail why it is impossible to construct a continuous relativistic Markov model in position space (Sec.~\\ref{relativistic}). Subsequently, we propose a simple non-Markovian generalization of the nonrelativistic Wiener process (Sec.~\\ref{axiomatic}). The specific functional shape of the relativistic propagator arises naturally, if one rewrites the Wiener propagator~\\eqref{e:solution_non-rel} in terms of an action integral. The paper concludes with a summary of the main results. ",
        "conclusions": "The search for a consistent description of diffusion processes in the framework of special relativity represents a longstanding issue in statistical physics. Several different relativistic generalizations of the nonrelativistic diffusion equation~\\eqref{e:diffusion} have been proposed in the past~\\cite{Ha65,MaWe96,BoPoMa99,HePa01,KoLi00,KoLi01,AbGa04,AbGa05}. The most prominent example is the telegraph equation~\\eqref{e:telegraph}. This second order partial differential equation (PDE) describes a non-Markovian diffusion process without superluminal particle propagation~\\cite{Ta22,Go50,RevModPhys.61.41,RevModPhys.62.375}. However, the solutions of the telegraph equation exhibit singular diffusion fronts propagating at light speed $c$, and may therefore be less appropriate for the description of massive particles. \\par To avoid such singularities, we pursue a complementary approach in this paper; i.e., instead of trying to find a relativistic diffusion PDE we focus on the transition probability density function (PDF) of a relativistically acceptable diffusion model. As the main result, an analytic expression for the transition PDF is proposed in Eq.~\\eqref{e:rel_PDF_best}. The simple relativistic diffusion process defined by this PDF avoids superluminal propagation, and reduces to the Wiener process~\\eqref{e:solution_non-rel} in the nonrelativistic domain. This analytically tractable relativistic diffusion model can be useful for practical calculations in the future (e.g., in high energy physics \\cite{AbGa04,AbGa05} and astrophysics~\\cite{RyLi85}). \\par Finally, it is worth emphasizing that, analogous to the solutions of TE but in contrast to the Wiener process, the proposed diffusion model is non-Markovian. This is not a drawback, as it is generally impossible to construct nontrivial continuous relativistic Markov processes in position space (cf. discussion in Sec.~\\ref{relativistic} and Ref.~\\cite{Du65}). In view of the close analogy between diffusion and quantum propagators in the nonrelativistic case, it would be quite interesting to learn, if this also means that it is impossible to formulate a relativistic particle quantum mechanics on the basis of standard path integrals~\\cite{Kleinert,St04}."
    },
    "0608/astro-ph0608222_arXiv.txt": {
        "abstract": "% The outburst of V838 Monocerotis in early 2002, and the subsequent appearance of its light echoes, occurred just before the installation of the Advanced Camera for Surveys into the {\\it Hubble Space Telescope\\/} (\\HST\\/)\\null. This fortunate sequence of events has allowed us to obtain spectacular \\HST\\/ images of the echoes, yielding not only pictures of extraordinary beauty, but also providing unique scientific information.  Our team has used the \\HST\\/ images to provide a direct geometrical distance to V838~Mon, based on polarimetric imaging, and limits on the distance based on the apparent angular expansion rates.  Several morphological features seen in the \\HST\\/ images strongly suggest that the illuminated dust was ejected from the star in a previous outburst, similar to the current one. In particular, a ``double-helix'' feature points exactly back to the star. Moreover, three-dimensional mapping of the outer edges of the dust suggests an overall ellipsoidal shape, centered on V838~Mon itself. And the appearance of the light echo in the most recent {\\it Hubble\\/} images is remarkably similar to that of a well-known planetary nebula, M27.  Future work on the \\HST\\/ images will include an analysis of interstellar-dust physics, in a situation where the scattering angle and illumination are unambiguously known, and visualization of a fully three-dimensional map of the dust distribution. ",
        "introduction": "The {\\it Hubble Space Telescope\\/} ({\\HST\\/}) has recently celebrated its 16th anniversary in orbit, having been launched aboard the NASA space shuttle in April 1990. Since then there have been four astronaut servicing missions between 1993 and 2002 (SM 1, 2, 3a, and 3b). We are currently hoping that there will be an SM4 in 2008, during which two new instruments would be installed: the Cosmic Origins Spectrograph (COS) and the Wide Field Camera 3 (WFC3). WFC3 will be a remarkable instrument that provides direct imaging at UV, optical, and near-IR wavelengths.  The \\HST\\/ is uniquely suited for high-resolution imaging. In this paper I would like to present some results from {\\it HST\\/} imaging of the remarkable light echoes around V838 Monocerotis. ",
        "conclusions": "V838 Monocerotis is illuminating the most spectacular light echoes seen in the history of astronomy. The high spatial resolution provided by the {\\it Hubble Space Telescope\\/} has yielded images of extraordinary beauty, as well as providing unique scientific information.  The \\HST\\/ images have provided a direct geometrical distance to V838~Mon, based on polarimetric imaging. Angular expansion rates have also given limits to the distance.  Several morphological features seen in the \\HST\\/ images strongly suggest that the illuminated dust was ejected from the star in a previous outburst, similar to the current one. The {\\it Hubble\\/} images show an extremely rich array of filaments and other features on small physical scales, suggestive of magnetohydrodynamic forces.  Future work on the \\HST\\/ images are expected to yield dust physics (i.e., the angular and color dependence of the scattering function) and a fully three-dimensional map of the dust distribution at high spatial resolution."
    },
    "0608/astro-ph0608708_arXiv.txt": {
        "abstract": "We report on spectroscopic observations of periodic comet 9P/Tempel~1 by the Alice ultraviolet spectrograph on the {\\it Rosetta} spacecraft in conjunction with NASA's Deep Impact mission. Our objectives were to measure an increase in atomic and molecular emissions produced by the excavation of volatile sub-surface material.  We unambiguously detected atomic oxygen emission from the quiescent coma but no enhancement at the 10\\% (1-$\\sigma$) level following the impact.  We derive a quiescent \\htwoo\\ production rate of $9 \\times 10^{27}$~\\mols\\ with an estimated uncertainty of $\\sim$30\\%.  Our upper limits to the volatiles produced by the impact are consistent with other estimates. ",
        "introduction": "The Deep Impact mission \\citep{A'Hearn:2005a} successfully placed a 364 kg impactor onto the surface of comet 9P/Tempel~1 at a relative velocity of 10.3~\\kms\\ on 2005 July 4 at 05:52:03 UT (as seen from Earth).  The event was observed by cameras aboard the mother spacecraft \\citep{A'Hearn:2005b} and by a large number of Earth- and space-based telescopes as part of an extensive campaign to study the comet prior to, during, and in the course of several days following the impact \\citep{Meech:2005}.  Several instruments on board the {\\it Rosetta} spacecraft, which is on a trajectory for a rendezvous with comet 67P/Churyumov-Gerasimenko in 2014, were programmed to observe comet Tempel~1 from June 28 through July 14, 2005.  During this period {\\it Rosetta} was located between 0.54 and 0.51~AU from the comet, which was observed at a solar elongation angle varying from 93.4\\deg\\ to 88.1\\deg\\ and solar phase angle varying from 65.1\\deg\\ to 72.5\\deg.  At the time of impact these parameters were 0.531~AU, 91.1\\deg, and 68.2\\deg, respectively.  Perihelion, at 1.506~AU, occurred on July 5. Imaging of the comet with the OSIRIS cameras on {\\it Rosetta} has been reported by \\citet{Kuppers:2005} and \\citet{Keller:2005}.  Alice was the only space-based spectrograph operating at wavelengths below 2000~\\AA\\ at the time of Deep Impact.  The spectral region spanned by Alice is rich is molecular and atomic emissions of the primary volatile components of the cometary coma.  These include several band systems of CO and the strongest resonance transitions of the principal atomic species, C, H, O, N, and S, that have been observed in many comets over the past 30 years using orbiting observatories and sounding rockets \\citep[e.g.,][]{Feldman:2004}. Alice is also sensitive to emissions of oxygen and carbon ions.  All of these species were expected to show an increase as a result of the excavation of volatile sub-surface material by the impactor.  Because of the comet's relatively large distance from the Sun and its low quiescent activity level, characteristic of Jupiter-family comets, Alice was only able to unambiguously detect the quiescent coma in \\ion{O}{1}~$\\lambda$1304, from which a quiescent water production rate of $\\sim9 \\times 10^{27}$~\\mols\\ is deduced.  \\ion{H}{1} Lyman-$\\beta$ is also marginally detected above the interplanetary background.  A slow secular decrease in the brightness of \\ion{O}{1}~$\\lambda$1304 with time was detected but we did not see any significant increase (i.e., 10\\%) as a result of the Deep Impact event.  Our upper limits on the increase in these emissions (and those of other species as well) are consistent with other reports of \\htwoo\\ and CO production from the impact. ",
        "conclusions": "Observations of periodic comet 9P/Tempel~1 were made with the Alice ultraviolet spectrograph on the {\\it Rosetta} spacecraft from six days before through ten days following the Deep Impact encounter with the comet.  \\ion{O}{1} emission from the quiescent comet was clearly detected and an \\htwoo\\ production rate of $9 \\times 10^{27}$~\\mols, with an estimated uncertainty of $\\sim$30\\%, was derived from the measured \\ion{O}{1} brightness profile.  Any enhancement either immediately following the impact or in the subsequent ten days was less than 10\\% (1-$\\sigma$).  Our upper limits to the volatiles produced by the impact are consistent with estimates from other observations.   \\begin{center}{\\bf ACKNOWLEDGMENTS}\\end{center}  We thank the ESA Rosetta Science Operations Centre (RSOC) and Mission Operations Center (RMOC) teams for their expert and dedicated help in planning and executing the Alice observations of comet Tempel~1.  We also thank Claudia Alexander (JPL), Kristin Wirth (ESTEC), Detlef Koschny (ESTEC), Paolo Ferri (ESOC), Maarten Versteeg (SwRI), and Gerhard Schwehm (ESTEC) for their continued support of Alice.  The work at Johns Hopkins University was supported by NASA's Jet Propulsion Laboratory through a sub-contract from Southwest Research Institute."
    },
    "0608/astro-ph0608364_arXiv.txt": {
        "abstract": "All the confirmed Soft Gamma-ray Repeaters\\linebreak have been observed with the EPIC instrument on the \\xmm\\ satellite. We review the results obtained in these observations, providing the most accurate spectra on  the persistent X-ray emission in the 1--10 keV range for these objects, and discuss them in the context of the magnetar interpretation. ",
        "introduction": "\\label{intro}   Soft Gamma-ray Repeaters (SGRs) were discovered as\\linebreak sources of short intense bursts of gamma-rays, and for a long time were considered as a puzzling category of Gamma-ray bursts. Their neutron star nature was immediately suggested by the 8 s periodicity seen in the famous event of 5 March 1979, but it was only with the discovery of their pulsating counterparts in the few keV region that this was finally proved. Although the main motivations for the magnetar model \\citep{duncan92,thompson95} were driven by the high energy properties of the SGRs bursts and giant flares, X-ray observations of the ``quiescent'' emission have provided fundamental information to understand the nature of these objects \\citep{woods04}.  Extensive observational programs have been carried out with the \\xte\\ satellite, focusing mainly on the SGRs timing properties. Long term studies based on phase-\\linebreak connected timing analysis revealed significant deviations\\linebreak from a steady spin-down \\citep{woods99d,woods00,woods02}, larger than the timing noise seen in radio pulsars and not linked in a simple way with the bursting activity. \\xte\\ has also been used to investigate the variations of the pulse profiles as a function of energy  and time \\citep{gogus02}, and to study the statistical properties of the bursts \\citep{gogus01}. However, the \\xte\\ observations are not ideal to accurately measure the flux and spectrum of these relatively faint sources located in crowded fields of the Galactic plane, since its non-imaging instruments suffer from source confusion and large uncertainties in the background estimate. Imaging satellites like \\sax, \\asca\\ and \\cha\\ have yielded useful spectral information, but it is only with the advent of the large collecting area \\xmm\\ satellite that high quality spectra of SGRs have been obtained, in particular with the \\epic\\ instrument \\citep{struder01,turner01}.   Here we review the results obtained with the \\xmm\\ satellite for the three confirmed SGRs in our Galaxy. There are also some \\xmm\\ data on \\zerocinque\\ in the Large Magellanic Cloud, but they are of limited use  due to the contamination from diffuse emission from the surrounding supernova remnant and will not be discussed here.        \\begin{table*}[t] \\begin{center} \\caption{\\xmm\\ observations and timing results for SGR 1806$-$20. } \\label{obs}       % \\begin{tabular}{cccccc} \\hline\\noalign{\\smallskip} Obs. & Date      & Duration &  Mode$^{(a)}$ and  exp. time   & Mode$^{(a)}$ and exp. time   & Pulse Period \\\\ &           &  (ks)        & PN camera   &  MOS1/2 cameras  &    (s) \\\\[3pt] \\tableheadseprule\\noalign{\\smallskip} A & 2003 Apr 3   & 32  & FF (5.4 ks)  &  LW (6 ks)    & 7.5311$\\pm$0.0003 \\\\ B & 2003 Oct 7   & 21  & FF (13.4 ks) &  LW (17 ks)    & 7.5400$\\pm$0.0003\\\\ C & 2004 Sep 6   & 51  & SW (36.0 ks) &  LW (51 ks)    & 7.55592$\\pm$0.00005 \\\\ D & 2004 Oct 6   & 18  & SW (12.9 ks) &  Ti (18 ks)    & 7.5570$\\pm$0.0003 \\\\ E & 2005 Mar 7   & 24  & SW (14.7 ks) &  Ti/FF (24 ks) & 7.5604$\\pm$0.0008\\\\ F & 2005 Oct 4   & 33  & SW (22.8 ks) &  Ti/FF (33 ks) & 7.56687$\\pm$0.00003\\\\ G & 2006 Apr 4   & 29  & SW (20.5 ks) &  Ti/FF (29 ks) & 7.5809$\\pm$0.0002\\\\ \\noalign{\\smallskip}\\hline \\end{tabular}\\\\ \\end{center} \\begin{footnotesize} $^{(a)}$ FF = Full Frame (time resolution 73 ms); LW = Large Window (time resolution 0.9 s);  SW = Small Window (time resolution 6 ms); Ti = Timing (time resolution 1.5 ms) \\end{footnotesize} \\end{table*} ",
        "conclusions": "Many of the results presented above fit reasonably well with the magnetar model interpretation. However, there are also a few aspects that require more theoretical and observational efforts to be interpreted in this framework, in particular when one considers the variety of different behaviors shown by these sources and their close relatives like the Anomalous X-ray pulsars \\citep{kaspi07}.   The long term variations seen in \\zerosei , the\\linebreak source observed more often with \\xmm , are qualitatively consistent with the predictions of the magnetar model involving a twisted magnetosphere \\citep{tlk02}. As mentioned above, according to this model, resonant scattering from magnetospheric currents  leads to the formation of a high-energy tail. A gradually increasing twist results in a larger optical depth that causes a hardening of the X-ray spectrum.  At the same time, the spin down rate increases because, for a fixed dipole field, the fraction of field lines that open out across the speed of light cylinder grows.  In addition, the stresses building up in the neutron star crust and the magnetic footprints movements lead to crustal fractures causing an increase in the bursting activity. Since spectral hardening, spin-down rate, and bursting rate increase with the twist angle, it is not surprising that these quantities varied in a correlated way in \\zerosei\\ (see Fig.~\\ref{1806} and \\citealt{gotz07}). However, as visible in Fig.~\\ref{1806}, while the spectral hardening took place gradually over several years, the spin-down variation occurred more rapidly in 2000. A recent analysis of \\xte\\ data around the time of the giant flare \\citep{woods06} shows that the correlation between spectral and variability parameters is indeed rather complex.  The long term flux evolution of  \\zerozero\\ is shown in Fig.~\\ref{1900}. It can be seen that, excluding the enhancements seen in correspondence of the flares, the luminosity remained always at the same level in the years 1997--2001, and then decreased slightly until the lowest value seen with \\xmm\\ in September 2005. The following observation of April 2006 showed that the decreasing luminosity trend has been interrupted by the recent onset of bursts emission. However, the moderate flux increase was not associated with significant changes in the X-ray spectral and timing properties, probably because the source is, up to now, only moderately active.  The  luminosity now reached by \\sedici ,\\linebreak \\mbox{$\\sim$3.5$\\times$10$^{33}$ erg s$^{-1}$} is the smallest ever observed from a SGR, and is similar to that of the low state of the transient anomalous X-ray pulsar XTE J1810$-$197 \\citep{ibrahim04,gotthelf04}. This low luminosity might be related to the long period ($\\sim$6 years) during which \\sedici\\ has not emitted bursts. However, the behavior of this source differs from that of the other SGRs that during periods of apparent lack of bursts changed only moderately their luminosity. In fact no bursts were detected from \\zerozero\\ in the three years preceding the \\xmm\\ observations: had its luminosity decreased with the same trend exhibited by \\sedici\\ it would have been much fainter than observed by \\xmm . Even more striking is the case of \\zerocinque , which has a high luminosity  ($\\sim2\\times10^{35}$ erg s$^{-1}$), despite no signs of  bursting activity have been observed in the last 15 years  (unless weak bursts from this source have passed undetected due to its larger distance and location in a poorly monitored sky region).  It thus seems that, similarly to the case of Anomalous X-ray Pulsars, SGRs comprise both  ``persistent'' sources (\\zerosei , \\zerozero\\ and \\zerocinque\\ ) and ``transients'' (\\sedici\\ ). The reasons behind this difference are currently unclear and not necessarily the same that differentiate between SGRs and AXPs."
    },
    "0608/astro-ph0608687_arXiv.txt": {
        "abstract": "The dwarf spheroidal (dSph) satellites of the Local Group have long been thought to be simple spheroids of stars embedded within extended dark matter halos. Recently, however, evidence for the presence of spatially and kinematically distinct stellar populations has been accumulating.  Here, we examine the influence of such components on dynamical models of dwarf galaxies embedded in cold dark matter halos. We begin by constructing a model of Andromeda~II, a dSph satellite of M31 which shows evidence for spatially distinct stellar components. We find that the two-component model predicts an overall velocity dispersion profile that remains approximately constant at $\\sim 10 - 11$\\,km\\,s$^{-1}$ out to $\\sim 1$\\,kpc from the center; this is despite wide kinematic and spatial differences between the two individual components. This prediction can be validated by detailed spectroscopic analysis of this galaxy.  The presence of two components may also help to explain oddities in the velocity dispersion profiles of other dSphs; we show that velocity dispersion profiles which appear to rise from the center outwards before leveling off---such as those of Leo I, Draco, and Fornax---can result from the gradual transition from a dynamically cold, concentrated component to a second, hotter, and more spatially extended one, both in equilibrium within the same dark halo. Dwarf galaxies with two stellar components generally have a leptokurtic line-of-sight velocity distribution which is well described by a double Maxwellian. This may be contrasted with other dynamical explanations such as a radially-dependent anisotropy in the stars' orbits. Interestingly, we find that multiple equilibrium components could also provide a potential alternative origin for ``extra-tidal'' stars (normally ascribed to tidal effects) in situations where corroborating evidence for tides --- such as elongation of the main body of the dwarf in the orbital direction or velocity gradients across its face driven by protruding tidal tails --- may be lacking. ",
        "introduction": "Dwarf spheroidal (dSph) galaxies are extremely faint systems whose relatively large size and velocity dispersion combine to make them premier examples of systems likely dominated by dark matter. They contain little or no gas, and individual stars are the only available tracer of their structure and kinematics. Although long regarded as relatively simple, single-component spheroids of stars, this view has been challenged by the advent of panoramic imaging techniques and multi-object spectrographs in large telescopes.  The analysis of the color-magnitude diagrams (CMDs) of dSphs shows, in many systems, compelling evidence for multiple and protracted episodes of star formation (eg. \\citealt{grebel1997,mateo1998a}, and references therein), as well as for the presence of spatially distinct stellar populations (\\citealt{dacosta1996,harbeck2001}). These populations often spill beyond the nominal ``tidal radius'' estimated from King-model fits to the inner surface brightness profile (eg. \\citealt{munoz2006} and references therein), a result which has brought into focus questions regarding the role of Galactic tides in limiting the spatial extent of dSphs, as well as the mass and spatial extent of their surrounding dark matter halos (\\citealt{johnston2002,penarrubia2007a}).  Spectroscopic studies have also uncovered remarkable complexity in the dynamics of dSphs. As sufficient data become available, a complex picture has arisen where kinematic oddities such as cold clumps (\\citealt{kleyna2003}), chemodynamically distinct stellar populations (\\citealt{tolstoy2004,battaglia2006,ibata2006}), and hints of rising dispersion profiles (\\citealt{wilkinson2004,munoz2005,koch2006a,walker2006a}), are common. This is at odds with the ``natural'' expectation for well-mixed stellar systems, where the velocity distribution is expected to be a smooth function of radius.  Observations thus suggest that it may be time to relax the standard assumption that dSphs are well described by a single well-mixed stellar system embedded within a dominant dark matter halo.  In \\cite{penarrubia2007a}, we examined the observational properties of single component stellar systems (described by King models) embedded in massive cold dark matter CDM) haloes (described by Navarro, Frenk \\& White (1996, 1997, hereafter NFW) profiles), and applied these models to the known population of Local Group dSphs. We examine here the general effect of allowing for multiple stellar components in these dynamical models. We begin in \\S 2 by motivating the modeling technique using the case of Andromeda~II (And~II; a dSph companion of M31), and then examine in general the velocity dispersion and surface density profiles of two-component stellar systems embedded within a CDM halo. We conclude in \\S3 with a discussion of our results and their wider implications.  \\begin{figure*} \\begin{center} \\includegraphics[angle=270, width=11.5cm]{figure1.ps} \\caption{Left panel: the density profile of Andromeda~II derived by McConnachie, Arimoto \\& Irwin (2007) (black squares). This is consistent with the sum of two spatially distinct stellar components; an extended component traced by horizontal branch stars (HB, blue diamonds/ dotted line) and an exponential component traced by the reddest red-giant branch stars (RRGB, red triangles/dot-dashed line).  Upper right panel: The predicted projected velocity dispersion profile for Andromeda~II (solid curve), together with the velocity dispersion profiles of the individual components, shown by dot-dashed and dotted curves. Lower right panel: the kurtosis of the line-of-sight velocity distribution as a function of radius for this model. Dwarfs with multiple equilibrium stellar components are expected to have a leptokurtic velocity distribution at all radii where both components contribute significantly.} \\end{center} \\label{fig:AndII} \\end{figure*} ",
        "conclusions": "In two component models, the centrally concentrated component is dynamically colder and has a decreasing influence on the overall dispersion profile as a function of radius. In the extreme case where each component dominates at different radii, we transition from a cold to a hot velocity dispersion profile. It is intriguing to note that such features in the velocity dispersion profile may not be unusual in Local Group dSphs. This is illustrated in Figure~\\ref{fig:ldf}, where we show data for Leo I, Fornax, and Draco, compiled from the literature. In each of these three cases, $\\sigma_p$ appears to rises by about a factor of $\\sim 2$ from its central value before starting to decline again. In the case of Fornax, we refer the reader to \\cite{battaglia2006}.  The velocity dispersion profile of Draco has previously been fitted by \\cite{mashchenko2006} using a single stellar component with a radially varying orbital anisotropy embedded in a massive halo, and \\cite{koch2006a} imply a similar anisotropy profile is required for Leo~I if it resides in an NFW halo.  \\cite{klimentowski2006} suggest that that these bumps in the velocity dispersion profiles are due to tidal effects. However, we consider this unlikely for two reasons (i) as concluded by \\cite{read2006a} and \\cite{klimentowski2006}, tides do not appear to affect dSph satellites inside of $1$\\,kpc, and so are unlikely to produce observable features at radii of typically $100 - 400$\\,pc (ii) the presence of tidally stripped stars manifest themselves as monotonically increasing features in the velocity dispersion profiles of dSphs (\\citealt{read2006a,klimentowski2006}). The features highlighted in Figure~3, however, rise and then level off or decline. Thus we conclude their origin is unlikely tidal.  For each galaxy in Figure~3 we have shown a possible decomposition of the velocity dispersion profile into a cold and hot component. These curves are inevitably degenerate, given that many parameters need to be fitted to describe two stellar components and a dark matter halo. In addition, we have not attempted to determine the uncertainties due to observation errors or bin sizes since these profiles have been taken direct from the literature.  Some caution is clearly required; for example, recent data for Leo I do not appear to show such as pronounced rise in the dispersion profile as found by Koch et al. (M. Mateo, private communication). As such, these curves merely demonstrate the {\\it plausibility} of multiple components in these systems and illustrate a few fairly general conclusions that can be drawn.  The first is that the rise in $\\sigma_p$ from the center outwards reflects the rise in the relative contribution of the more extended component, a fact that may be used to constrain the relative size of the two components and to pinpoint the radius where both components contribute more or less equally to $\\sigma_p(R)$.  This is where deviations from a simple Maxwellian-like velocity distribution are likely to be maximized. At this position, the line-of sight velocity distribution would be best characterized by a double Maxwellian distribution, distinguishing it from other proposals such as a radially-varying orbital anisotropy (see \\cite{mashchenko2006} for application of this idea to Draco).  The ratio of central-to-maximum velocity dispersion provides a lower bound to the ratio of the central velocity dispersions of the two components.  Interpreting Leo I as a two-component King model, the data shown in Figure~\\ref{fig:ldf} suggest that the central component has a central (projected) velocity dispersion of $\\sim 6$--$7$\\,km\\,s$^{-1}$, and that the outer component has $\\sigma_{p_2}(0) \\sim 13$-$14$ km s$^{-1}$.  If Leo I is a multiple component system, we expect that the line-of-sight velocity distribution of stars at about $R \\simeq 150-200$\\,pc should reveal evidence for the double Maxwellian distribution discussed previously. Similar testable inferences can be made for Fornax and Draco.  We note that a multiple component system is not required to exhibit differences in the stellar populations of its components. For example, one could build a dwarf by merging two progenitors of very different spatial extent and kinematics but similar age and metallicity. The remnant of such a merger would appear homogeneous to a CMD analysis but would retain the signature of the two progenitors in phase space. Within this context, evidence for a hierarchical origin of a dSph might be best appreciated in the dynamical evidence rather than in the CMDs.  Perhaps the most common deviation from single component models is the presence of stars that are clearly associated with the dSph, but which lie beyond the putative limiting radius of a King model fit. The presence of these ``extra-tidal'' stars can be caused by dynamical heating in the presence of the tidal field of the Galaxy and has attracted considerable attention because of its potential to constrain the mass and extent of a dSph's dark halo (e.g., \\citealt{johnston2002,read2006a,read2006b,klimentowski2006}).  As illustrated in Figure~2, the presence of a low density, extended second component, could in some instances be mistaken for ``extra-tidal'' stars, although in this scenario the stars are actually bound and in equilibrium with the dark matter halo. Note also that, in this interpretation, we would expect $\\sigma_p$ to {\\it decline} in this region, since these stars belong to the extended component. Thus, a multiple component scenario offers an alternative explanation for ``extra-tidal'' stars in situations where corroborating evidence for tides --- such as elongation of the main body of the dwarf in the orbital direction or velocity gradients across the face of the dwarf driven by protruding tidal tails --- may be lacking.  If the multiple component scenario is correct, a natural question which arises from this discussion is: {\\it how did these systems develop and preserve such complex structures}? A spatially varying star formation history, where subsequent epochs of star formation occurred in different volumes than previous ones, has been suggested by \\cite{kawata2006}.  Alternatively, \\cite{battaglia2006} suggest that there is tentative evidence for non-equilibrium kinematics in their data for Fornax, which might imply a merger-driven scenario. However, it is currently too soon to distinguish between these and other scenarios. Given the observational discovery of multiple structural components in dSph galaxies, it is important that the consequences for dynamical models of these systems are fully explored."
    },
    "0608/astro-ph0608146_arXiv.txt": {
        "abstract": "We study the dominant central giant elliptical galaxies in ``Fossil groups'' using deep optical (R-band) and near infrared (Ks-band) photometry. These galaxies are as luminous as the brightest cluster galaxies (BCGs), raising immediate interest in their link to the formation of BCGs and galaxy clusters. However, despite apparent similarities, the dominant fossil galaxies show non-boxy isophotes, in contrast to the most luminous BCGs. This study suggests that the structure of the brightest group galaxies produced in fossil groups are systematically different to the majority of BCGs. If the fossils do indeed form from the merger of major galaxies including late-types within a group, then their disky nature is consistent with the results of recent numerical simulations of semi-analytical models which suggest that gas rich mergers result in disky isophote ellipticals.  We show that fossils form a homogeneous population in which the velocity dispersion of the fossil group is tightly correlated with the luminosity of the dominant elliptical galaxy. This supports the scenario in which the giant elliptical galaxies in fossils can grow to the size and luminosity of BCGs in a group environment. However, the boxy structure of luminous BCGs indicate that they are either not formed as fossils, or have undergone later gas-free mergers within the cluster environment. ",
        "introduction": "It is believed that most of the large and luminous elliptical galaxies have formed via mergers of disk galaxies \\citep{toomre72,searle73}. This has been suggested by morphology density relation \\citep{dressler80}, the observed frequency of merging galaxies at high redshift and also extensively in computer simulations \\citep{barnes89}. Many of these luminous ellipticals ($M_B\\le$ -21) are found in rich galaxy clusters. This, however, does not imply that they are formed in cluster environment. The brightest cluster galaxies (BCGs) are of special interest as they reside close to the centroid of cluster X-ray emission \\citep{jf84}, and the centre of the dark matter distribution in clusters, as inferred from gravitational lensing \\citep{smith05} -- implying that they lie at the minimum of the cluster potential well. They also show various correlations with cluster properties.  In general, two main formation modes could be assumed for the hierarchical formation of BCGs according to their formation environment: {\\bf 1)} BCGs formed in the high velocity environment of clusters. Although the effectiveness of dynamical friction in bringing individual galaxies to the cluster centre via orbital decay will be reduced by the high velocity dispersion, infalling groups will still suffer rapid orbital decay if they survive long enough, and can then deposit their brightest galaxies in the cluster core, where they can merge and form a bright elliptical galaxy \\citep{mohr04,hausman78}. {\\bf 2)} BCGs are formed in the low velocity environment of groups, where dynamical friction causes the orbits of individual galaxies to decay, resulting in the merger of all large galaxies, if the group forms early and is left undisturbed for a sufficiently long period \\citep{dubinski98,ponman94,jones03}. The group containing this `ready made' BCG then provides the nucleus around which a cluster forms.  Elliptical galaxies show fine structures and are more complex than originally thought. These structures take the form of hidden disks, shells and bars, departures from pure elliptical isophotes \\citep{bender88} and variations in radial surface brightness profiles \\citep{kwkm00}, some of which are found to be environment dependent \\citep{habib04}. BCGs formed in the above two modes should display several observational signatures of how they formed. A useful probe is provided by galaxy morphology -- isophotal shapes, radial surface brightness profiles and the presence or not of multiple-nuclei. Different star forming histories are also expected. Some of these studies require space resolution data, but some can be studied using ground based observations, including the isophotes of elliptical galaxies.  Based on their isophotal shapes, elliptical galaxies can be classified as disky or boxy \\citep{bender88}. Low mass ellipticals, which are fast rotating, are usually disky isophote galaxies with positive fourth-order Fourier coefficient $B_4>0$. Some disky-isophote ellipticals might contain faint disks similar to S0 galaxies \\citep{scorza95}. With negative $B_4$, boxy isophote ellipticals are less rotationally supported. They generally contain flat cores \\citep{faber97,laine03} and show complex internal kinematics \\citep{emsellem04}. The observations of \\citet{rest01} show that it is very unlikely to find disky ellipticals which are also core galaxies. The majority of the BCGs are found to be core galaxies \\citep{laine03}. The distinct observed properties of disky and boxy isophotes of elliptical galaxies, such as their radio properties show that they are more than just an artifact of viewing angle or the projection on the plane of sky \\citep{bender89}.  It is important to understand the origin of isophotal shapes before they can be used to trace the formation history of ellipticals. \\citet{naab03} performed a large survey of dissipationless merger simulations of disk galaxies and found that unequal-mass 3:1 to 4:1 mergers lead to fast rotating disky ellipticals while equal-mass 1:1 to 2:1 mergers produce slowly rotating, pressure-supported ellipticals. However, \\citet{khochfar05} showed that the above scenario is not able to reproduce the observation that the fraction of boxy and disky ellipticals depends on galaxy luminosity. They argued that equal-mass mergers lead to boxy ellipticals and unequal-mass mergers produce disky ellipticals. However, major mergers between bulge-dominated galaxies result in boxy ellipticals, independent of the mass ratio, while merger remnants that subsequently accrete gas, leading to a secondary stellar disk with more than 20 per cent of the total stellar fraction, are always disky. More recently they showed that mergers of spiral galaxies alone cannot reproduce the kinematic and photometric properties of very massive elliptical galaxies \\citep{naab06}, nor can they reproduce the observed correlation between isophotal shapes and the luminosity of ellipticals.  Here we study the isophotes of brightest group galaxies (BGGs) in fossils.  A brief introduction to fossils is given in section 2 where we also describe the sample and observations. Our analysis and the results are presented in section 3. A discussion and concluding remarks are in section 4. We assume $H_0=70$ km~s$^{-1}$~Mpc$^{-1}$ and $\\Omega_m=0.3$ with cosmological constant $\\Omega_\\Lambda=0.7$ throughout.  \\begin{figure} \\center \\epsfig{file=new-j1119.ps,width=1.3in,height=1.2in} \\epsfig{file=new-j1256.ps,width=1.3in,height=1.2in} \\epsfig{file=new-j1331.ps,width=1.3in,height=1.2in} \\epsfig{file=new-j1340.ps,width=1.3in,height=1.2in} \\epsfig{file=new-j1416.ps,width=1.3in,height=1.2in} \\epsfig{file=new-j1552.ps,width=1.3in,height=1.2in} \\epsfig{file=new-n6482.ps,width=1.3in,height=1.1in} \\caption{Ellipse fits to the surface brightness distribution of the central galaxy in fossil groups. All the images are from the Ks-band observations with the exception of the most distant fossil, RX J1256.0+2556, which is observed in R-band. 5 arcsec scale bars are also shown.} \\label{chansoft} \\end{figure} ",
        "conclusions": "This analysis shows that, despite apparent similarities, the dominant giant elliptical galaxy in fossil groups are different in their isophotal shapes, compared to the brightest central galaxies in non-fossil systems, especially in rich clusters. Luminous elliptical galaxies in non-fossil groups and clusters do not present disky isophotes. Less luminous BGGs show disky and boxy isophotes in similar proportions. However, the observed frequency of boxy-isophote dominant fossil galaxies is apparently zero.  If the central galaxies of fossil groups have indeed been formed from the merger of all major galaxies within the inner regions of the group, as we suppose, then some of these mergers would have been gas-rich, as group spirals are incorporated into the central merger-remnant. The disky character of central fossil galaxies is then consistent with the findings of numerical simulations, that disky isophotes result from gas rich mergers. \\citet{khochfar05} highlight the importance of the role of gas in galaxy mergers and show that the isophotal shapes of merger remnants are sensitive to the morphology of their progenitors and to subsequent gas infall. In contrast, boxy isophote ellipticals are formed by equal-mass mergers of bulge dominated galaxies. Such mergers are likely to occur at the core of clusters where most of the galaxies are gas poor.   Almost 90\\% of BCGs studied by \\citet{laine03} are core galaxies, ie. with flatter slope near the nucleus in their radial surface brightness profiles. The study by \\citet{faber97} shows that a large fraction ($\\sim$70\\%) of core galaxies have boxy isophotes, with as low as $\\sim$10\\% with disky isophotes, implying a strong association of boxy isophotes with core galaxies. Taking into account the conversion of cuspy cores into flat low-density cores by black hole merging, \\citet{khochfar05} find that disky ellipticals should contain central density cusps whereas boxy ellipticals should in general be characterised by flat cores. Only rare low-luminosity boxy ellipticals, resulting from equal-mass mergers of disk galaxies, could have power-law cores. Space resolution data is needed to study the core of the galaxies to verify the core and power- law property of dominant fossil galaxies. If future observations show that fossil group dominant galaxies are power-law galaxies, as expected from their isophotes in light of the above argument, then they will be the first of such entities to grow to the size of BCGs.  \\begin{figure} \\center \\epsfig{file=lbggsigma.ps,width=2.6in,height=2.3in} \\caption{A comparison between the correlation of group/cluster velocity dispersion and absolute B-magnitude of the dominant galaxy. The data points from Table 1 (except RX J1119.7+2126) are shown with filled circles. Diamonds represent two of the fossil candidate OLEGs, out of four, in \\citet{yoshioka04}. The correlation is much tighter in fossils than in non-fossil groups (circles) and clusters (squares). The comparison non-fossil groups are GEMS X-ray selected groups with early-type BGG and with group scale X-ray emission (G-sample in \\citet{osmond04}). The cluster sample is selected from \\citet{girardi02} for which the luminosity of the central galaxy was available in \\citet{mohr04}.} \\label{sigma} \\end{figure}  The link between fossil galaxy groups and the BCGs is further motivated by observational space density estimates of fossils which is found to be 8\\% to 20\\% of X-ray luminous systems \\citep{jones03} and as large as the space density of poor and rich galaxy clusters combined. This means that there are enough fossils to provide BCGs to clusters. A recent theoretical study \\citep{milos06} predicts 5\\%-40\\% of galaxy groups and 1\\%-3\\% of galaxy clusters to be fossils. In the context of BCG formation, fossils appear to be suitable environments for the formation of luminous giant ellipticals around which clusters can form. This is a different formation mode to the one in which the BCG forms via mergers of brightest group galaxies during the cluster collapse. Numerical simulations suggest that the isophotes of the BCGs formed in the latter case will be predominantly boxy, characteristic of gas free (``dry'') mergers.  The BGGs and BCGs with non-boxy isophotes in Fig.\\ref{a4} could easily originate as fossil group central galaxies which have been incorporated into larger structures. The difference in photometric structure seen clearly in many BCGs and BGGS, especially the most luminous ones, does not rule out the possibility that they originated in fossil groups. This could still be the case provided that they have undergone later gas-free mergers within the cluster environment -- for example, with the BGGs of infalling galaxy groups. About 40\\% of BCGs contain at least one secondary nucleus\\citep{laine03}, which strongly suggests the action of late mergers within the cluster environment. If fossils are indeed old and undisturbed systems then they should be found to have a very {\\it low} incidence of multiple nuclei.  In support of the above argument we show (Fig \\ref{sigma}) a tight correlation between the luminosity of the central galaxy in fossils and the underlying gravitational mass probed by the group velocity dispersion.  This shows that fossils form a homogeneous population in which the luminosity of the central galaxy is strongly tied to the property of its parent group. This property of the fossils is consistent with their early formation epoch and absence of recent merger.  In contrast the large scatter in the distribution of the non-fossil galaxy groups and clusters, on the same plane, is understood to be merger driven. Absence of a recent major merger helps the dominant fossil galaxy preserve its original structure.  We conclude that there is high chance that disky BCGs are formed in fossil groups. Boxy BCGs, could still result from progenitor fossils, but would need to have undergone a dry merger within the cluster, probably as a result of merger with BGGs of infalling groups.   We would like to thank Doug Burke, for his involvement in the UH 88'' observations, and the INT for the service observations, Ewan O'Sullivan for his involvement in the KPNO observations of  NGC 6482, Graham P. Smith for his valuable input which helped to improve the discussion, and Arif Babul for his useful comments."
    },
    "0608/astro-ph0608693_arXiv.txt": {
        "abstract": "We investigate the X-ray and mass distribution in the merging galaxy cluster 1E 0657-56. We study head-on collisions of two virialized clusters with an NFW density profile in the $\\Lambda$CDM universe using an $N$-body + hydrodynamical code. A clear off-set of an X-ray peak from a mass peak, which is like what is reported in \\citet{Clow04}, is first reproduced in the $N$-body + hydrodynamical simulations. We estimate the ram pressure-stripping conditions of the substructure in mergers of two NFW dark halos using a simple analytical model. We find that the ram pressure dominates the gravity of the substructure when the smaller cluster's mass is less than approximately one tenth of the larger cluster's mass. The characteristic X-ray and mass structures found in 1E 0657-56 suggest that neither the ram pressure nor the gravitational bound force overwhelms the other and that the mass ratio between the progenitors is near the critical value mentioned above. ",
        "introduction": "According to the standard scenario of cosmological structure formation, clusters of galaxies form through mergers of smaller subclusters. In fact, some of them are forming now. Merging clusters are the sites of structure formation in the universe that can be investigated in detail via different types of observations. Mergers play important roles in cluster evolution itself. Cluster mergers cause bulk flow motion, turbulence, shocks, and/or contact discontinuities in the intracluster medium (ICM), which give us clues to investigate various physical processes in the ICM (see \\cite{Sara02} for a review). Turbulence and shocks most likely play crucial roles in particle acceleration in the ICM (e.g. \\cite{TaNa00}; \\cite{Ohno02}; \\cite{Taki02}; \\cite{Fuji03}; \\cite{Buru03} ). Strong bulk flow motion and violent pressure changes in the ICM during mergers may affect star formation activities of the member galaxies (\\cite{Fuji99}; \\cite{Bekk03}). $N$-body + hydrodynamical numerical simulations have been carried out to study cluster merger physics (e.g. \\cite{Schi93}; \\cite{Ishi96}; \\cite{Roet97}; \\cite{Taki99}, \\yearcite{Taki00}; \\cite{Rick01}; \\cite{Ritc02}). Comparison of such numerical simulations with different kinds of observations give us deep insight into the cluster physics.  1E 0657-56 is one of the most well-known examples of a merging cluster. It is the hottest known cluster and has a very powerful radio halos \\citep{Lian00}. Its simple geometry makes this cluster one of the most suitable case to investigate cluster merger physics. There are two peaks in both the X-ray surface brightness distribution \\citep{Mark02} and galaxy distribution \\citep{Barr02}, but their positions do not agree with each other. Observations of the line-of-sight velocities of the member galaxies suggest that its collision axis is almost perpendicular to the line-of-sight \\citep{Barr02}. From X-ray observations, a bow shock and a cold front are found in front of the smaller subcluster, which suggest that its Mach number is $2 - 3$ \\citep{Mark02}.  Recently,  \\citet{Clow04} investigated the mass distribution in 1E 0657-56 through weak gravitational lensing. They show clear offsets of the mass density peaks from the X-ray peaks, and that the mass distribution is quite similar to the galaxy one. The smaller substructure in mass is ahead of the X-ray one. They claim that this structure occurs because ICM experience ram pressure but dark matter (DM) and galaxies do not. Although the above-mentioned naive ram pressure-stripping scenario seems to be correct, such characteristic off-sets of X-ray peaks to mass peaks have never been reported in the past numerical simulations. In this paper, we show the first results that successfully reproduce such characteristic structures in the $N$-body + hydrodynamical simulations, and discuss their implications using a simple analytical model.  The rest of this paper is organized as follows. In \\S 2 we describe the adopted numerical method and initial conditions for our simulations. In \\S 3 we present the simulation results. In \\S 4 we show simple analytical estimation for ram pressure-stripping conditions in mergers of two clusters with an NFW density profile. In \\S 5 we summarize the results. ",
        "conclusions": "Let us discuss the ram pressure-stripping conditions in merger of two clusters with an NFW DM density profile. We consider the merger of two clusters with masses $M_1$ and $M_2$ ($M_1>M_2$), respectively. We concentrate on the physical status of the ICM in the core region of the smaller subcluster. If the gravity on the subcluster's ICM is weaker than the ram pressure force in unit volume, the ICM will be stripped from the substructure potential. This means, \\begin{eqnarray} \\frac{G m_2 \\rho_2}{r_2^2} < A (\\pi r_2^2 \\rho_1 v^2)(\\frac{4}{3} \\pi r_2^3)^{-1}, \\label{eq:condition} \\end{eqnarray} where $G$ is the gravitational constant, and $\\rho_1$ and $\\rho_2$ are the central gas density of the subcluster 1 and 2, respectively. $r_2$ and $m_2$ are the scale radius of the DM profile and the DM mass inside $r_2$ for the cluster 2, respectively. Therefore, the relation between $m_2$ and $M_2$ is \\begin{eqnarray} g(M_2) \\equiv \\frac{m_2}{M_2} = \\frac{\\ln2-1/2}{\\ln(1+c)-c/(1+c)}, \\end{eqnarray} where, $c$ is a concentration parameter of an NFW profile, and weakly depends on the halo mass. $A$ is a fudge factor of an order of unity. It is most likely that $A \\lesssim 1$ because all of the ram pressure force is not effective in stripping the gas from the substructure. Some might be used on the excitation of small-scale eddies through Kelvin-Helmholtz instability, and some on the adiabatic compression and shock heating of the ICM, and so on. The collision velocity $v$ has an order of \\begin{eqnarray} v^2 \\simeq \\frac{2G(M_1+M_2)}{R_1+R_2}, \\label{eq:velocity} \\end{eqnarray} where $R_1$ and $R_2$ are the virial radii for the cluster 1 and 2, respectively. It is convenient to introduce a new parameter $\\alpha \\equiv M_2/M_1$. Then, the stripping condition of inequality (\\ref{eq:condition}) becomes \\begin{eqnarray} F(\\alpha: M_1) \\equiv \\alpha^{2/3 - w} \\frac{1+\\alpha^{1/3}}{1+\\alpha} - \\frac{3 A}{2 g(\\alpha M_1) c(\\alpha M_1)} < 0 \\label{eq:condition2}. \\end{eqnarray} Deriving inequality (\\ref{eq:condition2}), we use the scaling relation that $R_2/R_1 = \\alpha^{1/3}$ and $\\rho_2/\\rho_1 = \\alpha^{-w}$. In the $\\Lambda$CDM universe ($\\Omega_0=0.25$, $\\lambda_0=0.75$), $w \\simeq 1/4$ assuming that $\\rho_1$ and $\\rho_2$ behave like a characteristic density in an NFW profile \\citep{Nava97}.   \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){figure3.eps} \\end{center} \\caption{Function $F(\\alpha: M_1)$ defined by equation (\\ref{eq:condition2}) for $M_1=1.0 \\times 10^{15} M_{\\odot}$ in the $\\Lambda$CDM universe ($\\Omega_0=0.25$, $\\lambda_0=0.75$). The solid, dashed, and dot-dashed lines represent the cases of $A=0.6$, $0.4$, and $0.2$, respectively. $F<0$ means that the ram-pressure dominates the gravity in the substructure.} \\label{fig:F(alpha)} \\end{figure} Figure \\ref{fig:F(alpha)} shows the function $F(\\alpha: M_1)$ for $M_1=1.0 \\times 10^{15} M_{\\odot}$ in the $\\Lambda$CDM universe ($\\Omega_0=0.25$, $\\lambda_0=0.75$). The solid, dashed, and dot-dashed lines represent the cases of $A=0.6$, $0.4$, and $0.2$, respectively. In any case, $\\alpha$ is less than $\\sim 0.1$ when $F(\\alpha)<0$. This means that the ram pressure-stripping is more effective for smaller subclusters. Although our estimation here is rather crude, it is interesting that this criteria of $\\alpha \\sim 0.1$ is close to the mass ratio of our simulations where the clear off-set appears. Obviously, such an off-set does not appear at all if the ram pressure-stripping does not work effectively. The X-ray peak will correspond with the mass peak because the ICM behaves like the DM. On the other hand, if the ram pressure overwhelms the gravity, the ICM in the substructure will not be able to penetrate the larger cluster's center and will be repelled. Furthermore, the larger cluster's ICM is so hot that it cannot be bound by the substructure's gravitational potential. Therefore, we will see a mass peak associated with no X-ray peak. Clear off-sets in the simulation and 1E 0657-56 suggest that the parameter $\\alpha$ is close to the above-mentioned critical value where neither the ram pressure nor the gravity absolutely dominates the other.  In our simulations, we do not take into account a dense cool core in the central region of a cluster. However, this probably affects the ram pressure stripping conditions. The existence of the cool core changes $\\rho_1$ and $\\rho_2$ in inequality (\\ref{eq:condition}). Thus, the scaling relation $\\rho_2/\\rho_1=\\alpha^{-w}$ could be modified and/or have some dispersion. In case of 1E0657-56, the smaller subcluster certainly has a cool core, but the larger one seems to have none. Therefore, $\\rho_2/\\rho_1$ can be larger than when neither subcluster has a cooling core. This makes more difficult to strip the ICM.  Whereas the smaller X-ray peak in 1E0657-56 is triangle, that in our simulation results is rather round. The absence of a cool core might cause this difference. Another physical process that may play a crucial role is the magnetic field in the ICM. The magnetic field along the boundary layer probably works to maintain the smaller subcluster's gas as a distinct structure after it is stripped off the DM potential through both the magnetic tension and suppression of heat conduction \\citep{Asai04}. Dynamical motion of the substructure itself possibly produces this kind of magnetic field configurations \\citep{Vikh01}. Furthermore, temperature gradients in the boundary layer might produce the magnetic field structure through Weibel-type plasma instabilities \\citep{Okab03}. Three-dimensional high-resolution magnetohydrodynamic simulations will be useful in order to investigate detailed evolution.  Please note that an off-set of an X-ray peak to a mass peak is not a structure in dynamical equilibrium but a transient one. A characteristic timescale of the ram pressure-stripping is estimated to be $\\sim R_2/v$, which becomes $0.42$ Gyr for the model presented in \\S 2. In the simulation, it takes $\\sim 0.7$ Gyr to strip the ICM from the substructure potential. Certainly this is roughly equal to the above-mentioned value."
    },
    "0608/astro-ph0608370_arXiv.txt": {
        "abstract": "{} {As part of the DART project we have used the ESO/2.2m Wide Field Imager  in conjunction with the VLT/FLAMES\\thanks{Based on FLAMES observations collected at the European Southern Observatory, proposal 171.B-0588} GIRAFFE spectrograph to study the detailed properties of  the resolved stellar population of the Fornax dwarf spheroidal galaxy out to and beyond its tidal radius. Fornax dSph has had a complicated evolution and contains significant numbers of young, intermediate age and old stars. We investigate the relation between these different components by studying their photometric, kinematic and abundance distributions. } { We re-derived the structural parameters of the Fornax dwarf spheroidal  using our wide field imaging covering the galaxy out to its tidal radius, and  analysed the spatial distribution of the Fornax stars of different ages as selected from Colour-Magnitude Diagram analysis. We have obtained accurate velocities and metallicities from spectra in the Ca~II triplet wavelength region for 562 Red Giant Branch stars which have velocities consistent with membership in Fornax dwarf spheroidal.} {We have found evidence for the presence of at least three distinct stellar components: a young population (few 100 Myr old) concentrated in the centre of the galaxy, visible as a Main Sequence in the Colour-Magnitude Diagram; an intermediate age population (2-8 Gyr old); and an ancient population ($>$ 10Gyr), which are distinguishable from each other kinematically, from the  metallicity distribution and in the spatial distribution of stars found in the Colour-Magnitude  Diagram. }{ From our spectroscopic analysis we find that the ``metal rich'' stars ([Fe/H]$> -1.3$) show a less extended and more concentrated spatial distribution, and display a colder kinematics than the ``metal poor'' stars ([Fe/H$<-1.3$). There is tentative evidence that the ancient stellar population in the centre of Fornax does not exhibit equilibrium kinematics. This could be a sign of a  relatively recent accretion of external material, such as the merger of another galaxy or other means of gas accretion at some point in the fairly recent  past, consistent with other recent evidence of substructure (Coleman et al. 2004, 2005). } {} ",
        "introduction": "The Fornax Dwarf Spheroidal galaxy (dSph) is a relatively distant companion of the Milky Way at high galactic latitude ({\\it b}~=~$-$65.7$^o$).  It is located at a distance of 138 $\\pm 8$~kpc \\citep[e.g.][]{Mateo1998, Bersier2000}, with an heliocentric velocity v$_{\\rm hel} =$ 53 $\\pm 2$ km/s \\citep{Mateo1991}, a luminosity $M_{\\rm V} = -13.0 \\pm 0.3$, and a central surface brightness $\\Sigma_{\\rm 0,V} = 14.4 \\pm 0.3$ mag/arcmin$^2$ \\citep{IH1995}. The Fornax dSph is one of the most massive and luminous of the dwarf galaxy satellites of our Galaxy, second only to Sagittarius, with a total (dynamical) mass in the range $10^8$-$10^9$ \\sm \\citep{Mateo1991, Lokas2002, walker2006, Battaglia2006a}.  Similarly to Sagittarius, Fornax has its own globular cluster population, but contrary to Sagittarius all of the Fornax globular clusters appear to be metal poor \\citep[e.g.][]{strader2003, letarte2006a}. Fornax is known to contain a large number of carbon stars with a wide range of bolometric luminosities indicating significant mass (and hence age) dispersion at intermediate ages \\citep[e.g.][]{Demers1979, aaronson1980, stetson1998, azzopardi1999}. This is supported by a well populated intermediate age sub-giant branch and red clump. Fornax also contains a sizable population of RR Lyrae variable stars with an average metallicity of [Fe/H]$_{\\rm RR} \\sim -1.6 \\pm 0.2$ \\citep{Bersier2002}. From detailed Colour-Magnitude Diagram (CMD) analysis \\citep[e.g.][]{stetson1998, buonanno1999, saviane2000, pont2004} it is clear that, unlike most dSphs in the Local Group, Fornax has had a long history of star formation, which has only ceased few hundreds Myr ago. However, in common with most dSph galaxies, Fornax does not appear to have any HI gas at present \\citep[e.g.][]{Young1999}. It is difficult to understand how Fornax lost all it gas. It is possible that, as consequence of the recent star formation, the gas is fully ionized, however detections of ionized gas are difficult.  A recent proper motion estimate suggests that Fornax has a low eccentricity polar orbit, that crossed the path of the Magellanic Stream about 190~Myr ago \\citep{dinescu2004}, and that the excess of small scale structure in the Magellanic stream, tracing back along the proposed orbit of Fornax comes from gas lost by Fornax during this encounter.  The authors propose this encounter as a viable mechanism to make Fornax lose its gas and stop forming stars. However, previous proper motion measurements do not suggest such an encounter \\citep{piatek2002}.  Fornax has been the centre of some attention recently for the discovery, in photometric surveys, of stellar over-densities in and around the system \\citep{coleman2004, coleman2005, ol2006}. These over-densities have been interpreted by Coleman et al. as {\\it shell structures} caused by the recent capture of a small galaxy by Fornax.  There have been previous spectroscopic studies of individual stars in Fornax. Studies of kinematic properties indicated a (moderate) central mass-to-light ratio $M/L \\gtrsim 5 $($M/L_{\\rm V}$)$_\\odot$, possibly as high as 26 \\citep[e.g.][]{Mateo1991}, and a global $M/L_{\\rm V} \\sim$ 10-40 larger than the $M/L_{\\rm V}$ of the luminous component \\citep[][]{walker2006}. This is much higher than is typical for globular clusters, which have mass-to-light ratios in the range 1-3 \\citep[e.g.][]{mac2005}, but considerably lower than that found for many other dSphs such as Draco and Ursa Minor \\citep[$M/L >$ 100 ($M/L_{\\rm V}$)$_{\\odot}$, e.g.][]{kleyna2003, wilkinson2004}.  There have been two previous Ca~II triplet (CaT) studies of Red Giant Branch (RGB) stars in the central region of Fornax (33 stars, Tolstoy et al. 2001; 117 stars, Pont et al. 2004), both using VLT/FORS1. There are also detailed abundance studies of 3 field stars \\citep{shetrone2003} and 9 globular cluster stars \\citep{letarte2006a} made with VLT/UVES. This  is about to be extended dramatically by VLT/FLAMES \\citep[$\\sim 100$ stars,][]{letarte2006b}. These spectroscopic studies conclude that the bulk of the stellar population is more metal rich than could be inferred from the position of the Red Giant Branch (RGB) in the CMD, with a peak at [Fe/H] $\\sim -0.9$, a metal poor tail extending out to [Fe/H] $= -2$ and a metal rich tail extending beyond [Fe/H] $= -0.4$ with more than half the stars on the RGB having ages $<$~4~Gyr.  Here we present the first results of a study of the Fornax dSph from the DART (Dwarf Abundances and Radial velocities Team) large programme at ESO. The aim of DART is to analyse the chemical and kinematic behaviour of individual stars in a representative sample of dSphs in the Local Group, in order to derive their star formation and chemical enrichment histories, and explore the kinematic status of these objects and their mass distribution. Our targets are Sculptor (Tolstoy et al. 2004), Fornax and Sextans dSphs. From our large program at ESO we obtained for each of these galaxies: extended WFI imaging; intermediate resolution VLT/FLAMES spectra in the CaT region to derive metallicity and velocity measurements for a large sample of RGB stars covering an extented area; and VLT/FLAMES high resolution spectra in the central regions, which give detailed abundances over a large range of elements for $\\sim 100$ RGB stars.  In this work we present our results for both wide field imaging and intermediate resolution FLAMES spectroscopy for 562 RGB kinematic members of Fornax out to the tidal radius. The relatively high signal-to-noise (S/N $\\approx$ 10-20 per \\AA) of our data has enabled us to derive both accurate metallicities ($\\approx$ 0.1 dex from internal errors) and radial velocities (to $\\approx \\pm$2 km/s). We can thus study the different stellar populations in Fornax and how their metallicities and kinematic properties vary spatially with a dramatically larger sample than any previous study. Our study differs from Walker et al. (2006) in that we have both abundance and kinematics and this has been shown to be a crucial piece of information for disentangling surprisingly complicated galaxy properties (Tolstoy et al. 2004).  In future works we will present our CaT calibration (Battaglia et al. 2006b), detailed kinematic modelling of Fornax dSph (Battaglia et al. 2006a) and accurate star formation history from our VLT/FLAMES high resolution abundances of Fornax (Letarte et al. 2006b). ",
        "conclusions": "\\subsection{Age \\& Metallicity Gradients}  One of the main results of our imaging and spectroscopic survey of the Fornax dwarf spheroidal galaxy is the presence of a population gradient. The stellar population shows a clear variation in age and metallicity as a function of radius (e.g. Figs.~6, 14 and 19).  We can associate most of the more metal poor component with an ancient population ($> 10$Gyr), and most of the more metal rich component with an intermediate age population (2-8 Gyr) with overlapping metallicity distributions.  Population gradients have been detected in several other dwarf spheroidal galaxies in the Local Group from imaging (e.g., Harbeck et al. 2001), and spectroscopy (e.g., Tolstoy et al. 2004). This must point to some common process in the formation of stars in these small systems, as this variation most likely reflects the original spatial distribution of gas from which the different stellar populations were formed.  It is unlikely that the spatial distribution of one stellar component can have changed significantly with respect to another over time. Although weak encounters could diffuse an older component, it is unlikely to be an important effect because the relaxation time (Binney \\& Tremaine 1987) of Fornax is $> 1.5 \\times 10^{10}$ years, i.e. longer than the age of the ancient population. Therefore, the present spatial distribution of stars of different ages gives us an indication of how the gas from which they were formed was distributed, and requires an explanation for how gas has apparently been progressively ``removed'' from the outer regions, and more recently, apparently, completely from the system.  Star formation is likely to have a dramatic impact on the interstellar medium of these small galaxies. It is possible that gas is frequently blown away as the result of supernovae type II explosions of massive stars \\citep[e.g.][]{maclow1999, mori2002}. This expelled gas should be able to fall back onto the galaxy, if the galaxy has a large enough potential, and on its return the gas will sink deeper into the central regions, where further generations of star formation can occur, with potentially different spatial, kinematic, metallicity and age characteristics.  This next generation stellar population will likely be a more centrally concentrated, younger, more metal rich component.  In this case we should expect a direct correlation between the mass of the galaxy and the fraction of gas retained by the system, as the deeper potential well of the most massive dSphs is more capable of retaining more of the original gas (e.g., Ferrara \\& Tolstoy 2000). Consistent with this picture, Fornax differs from the less massive Sculptor dSph in having a much more extended star formation history.  However this correlation is not a straight forward prediction of the properties of all the dwarf galaxies in the Local Group. This is because it is not only the mass of galaxies that determines their ability to retain an interstellar medium; dSph typically live in a fairly active environment with competing effects of the tidal field of our much larger Milky Way and several fairly large interacting companions (e.g., the Magellanic Clouds), and the debris left around by these interactions (e.g., the Magellanic Stream).  N-body simulations indicate that tidal stripping combined with ram pressure stripping is a feasible combination to explain the lack of gas at the present time in dSphs \\citep{mayer2005}. The importance of these effects depends strongly on the orbital parameters of the satellite, the apocenter-to-pericenter ratio, and the time at which the satellite enters the Milky Way potential. Whilst the satellite is orbiting the Galaxy, part of its gas may be ionized by the Milky Way corona, perhaps in combination with the cosmic UV background. The gas can be then removed over time, with varying degrees of efficiency, if the apocenter-to-pericenter ratio is large enough.  Recent proper motion measurements for Fornax suggest a pericentric distance of the order of the present Galactocentric distance \\citep{piatek2002, dinescu2004}, and an almost circular orbit with an orbital period of $\\sim$ 4-5 Gyr \\citep{dinescu2004}. Even though proper motion measurements are rather uncertain, this would suggest that we can exclude strong interaction with the Milky Way. Thus, as suggested by \\citet{mayer2005}, an object like Fornax may not be completely stripped of its gas, but would be able to retain it in its central regions and thus maintain an extended star formation history.  N-body models that predict dIrr to be the progenitors of dSphs, transformed into spheroids by tidal stirring from the Milky Way \\citep[][]{mayer2001a, mayer2001b}, predict an increase in the star formation activity of a satellite after pericentric passage.  The pericentric passage can also provoke the formation of a bar that funnels gas into the central region, making the gas distribution more centrally concentrated. These models are highly dependent upon the starting conditions for the satellite. It is found that an apocentre-to-pericentre ratio of $\\sim$ 5 and a pericentre distance of $\\sim$ 50 kpc would be required to transform a dIrr into a dSph. Fornax has never been specifically modeled, however, as it is now thought to be close to its pericentre and the previous passage probably occurred around 4-5 Gyr ago, this would correspond well to the time frame of the last period of intense star formation. This scenario could explain both the variation in the spatial distribution of the Fornax stellar population and the large fraction of intermediate age stars. Specific simulations, emulating the orbital parameters of Fornax, are required for a more detailed comparison between these observations and N-body simulations.  \\subsection{Peculiar Kinematics}  In addition to the different spatial distribution of ancient, intermediate age and young stars, we also find a different kinematic behaviour of metal rich and metal poor stars.  This behaviour appears similar to what is seen, for example, in our Milky Way. The stellar bulge of our Galaxy is much less extended and kinematically colder than the stellar halo.  This distinction may merely reflect the different spatial distribution required by a star moving with a larger velocity to reach larger distances or it may be telling us something more fundamental about the different conditions of the formation of these two different components. To test this we can use the Jeans equation to derive the line of sight velocity dispersion predicted for stars following the different density profiles of old and intermediate age stars \\citep{binney1987}. This calculation shows that it is difficult to interpret the observed differences in velocity dispersion of MR and MP stars as being only due to their different spatial distribution \\citep{Battaglia2006a}, and the same result is found in N-body simulations \\citep[e.g][]{kawata2006}. In a future paper we will present detailed kinematic modelling of these distinct stellar populations, taking into account the possibility of different orbital properties for the two populations \\citep{Battaglia2006a}.  Another aspect that deserves further investigation is the apparently {\\it non-equilibrium} kinematics of the metal poor, presumably the oldest, stellar population in the center of Fornax. A possible interpretation of this is that we are seeing a remnant of accreted material. Given the presence of 5 globular clusters in Fornax, it would not be surprising for this to be a disrupted globular cluster. In this case we would expect the metallicity distribution of the stars forming the velocity feature to be extremely narrow, however the metallicity of the stars in the double peak ranges from $-2\\lesssim$[Fe/H]$\\lesssim -1.4$, showing a much larger spread than can be accounted for by a globular cluster.  An alternative explanation, as proposed by \\citet{coleman2004, coleman2005}, is that Fornax accreted a smaller galaxy at some point in the recent past.  These authors propose the accretion of a gas-rich dwarf $\\sim$ 2 Gyr ago.  In this case the MP stars with unrelaxed kinematics could be the remnant of the disrupted old stellar component of this accreted galaxy, and the asymmetric spatial distribution of the young stars (BL and MS) could be a result of the gas which subsequently formed stars as a consequence of the accretion process. Measuring the metallicities of the young stars in Fornax could help to answer this question by determining whether or not they are consistent with the bulk of the older stellar population in Fornax. Unfortunately we do not have spectroscopic metallicities of these stars, but the indications from the colour of the Blue~Loops in the CMD (see section 3.3.3) suggest that these stars have a metallicity of [Fe/H]$\\sim -0.85 \\pm 0.15$, consistent with the mean of the youngest RGB stars (age $>$1-2 Gyr) in our VLT/FLAMES sample. However, our FLAMES sample also contains a significant number of stars with higher metallicity, up to [Fe/H]$\\sim-0.1$.  Assuming we can directly compare metallicities derived with these two different methods, this could be an indication that Fornax has recently ($<$1-2 Gyr ago) been polluted with external gas of lower metallicity than the gas in Fornax at that time. This would also explain the peculiar kinematics of the MP stars in the centre of Fornax.  However, this conclusion requires further confirmation, such as high resolution follow-up of the metal rich RGB stars in the central region (Letarte et al. 2006b), spectroscopic analysis of the Blue~Loop and/or main sequence stars, and detailed dynamical modelling."
    },
    "0608/astro-ph0608236_arXiv.txt": {
        "abstract": "We present the result of investigations into two theories to explain the star formation rate-density relationship. For regions of high galaxy density, either there are fewer star forming galaxies, or galaxies capable of forming stars are present but some physical process is suppressing their star formation. We use  HI Parkes All Sky Survey's (HIPASS) HI detected galaxies and infrared and radio fluxes to investigate star formation rates and efficiencies with respect to local surface density.  For nearby (vel$<$10000 km s$^{-1}$) HI galaxies we find a strong correlation between HI mass and star formation rate. The number of HI galaxies decreases with increasing local surface density. For HI galaxies (1000$<$vel$<$6000 km s$^{-1}$) there is no significant change in the star formation rate or the efficiency of star formation with respect to local surface density. We conclude the SFR-density relation is due to a decrease in the number of HI star forming galaxies in regions of high galaxy density and not to the suppression of star formation. ",
        "introduction": "\\label{Sec:Paper2Introduction} \\citet{dressler1980b}, finds a well--defined relationship between local galaxy density and morphology for galaxies in clusters, the `morphology--density relation'. Several later studies, e.g. \\citet{hashimoto1998}, \\citet{goto2003}, \\citet{sodre2003} and \\citet{nuijten2005} use large optical surveys and 2 or 3D local galaxy density to investigate the Morphology-density relation. They confirm \\citet{dressler1980b}'s cluster results: as local galaxy density increases, there is a decrease in the relative number of spiral galaxies and an increase in the number of elliptical and SO galaxies.  Two studies investigating the morphology-density relation with respect to the star formation rate (SFR) and local galaxy density (hereafter the SFR-density relation) are \\citet{lewis2002} and \\citet{gomez2003}. These studies are based on optical observations with galaxy SFRs calculated from H$\\alpha$ fluxes. To determine the local galaxy density \\citet{gomez2003} and \\citet{lewis2002} use the local surface density. \\citet{lewis2002} uses the SFR for a given density range for clusters compared to field galaxies.  \\citet{gomez2003} uses two samples, one sample is separated into field and group/clusters galaxies and the other is the Sloan Digital Sky Survey.  Both \\citet{lewis2002} and \\citet{gomez2003} conclude the SFR for a given density range decreases as the local surface density increases. \\citet{balogh2004} investigate the relationship between the distribution of galaxies and local density and find that star-forming and quiescent (elliptical) galaxies vary with local density.  However, they find the `distribution of H$\\alpha$ equivalent widths' of the star-forming population does not depend strongly on environment.  Our work is based on the HI Parkes All Sky Survey Optical Catalogue (HOPCAT), \\citep{doylemt2005mnras}. HOPCAT is, at the present time, the largest optically matched HI radio detected catalogue of the whole southern sky. For the first time we investigate the SFR-density relation using a large sample of galaxies containing the basic fuel for star formation, HI. For regions of high galaxy density we ask, `is the decrease in SFR due to fewer star forming galaxies or are galaxies capable of forming stars present but some physical process is suppressing their star formation?'  When investigating the SFR-density relation, various 3D methods to determine the local galaxy density are used \\citep{allam1999, martinez2002, sodre2003} which require redshifts. At the time of writing, no redshift catalogue of the whole of the southern sky exists, hence we use the local surface density, $\\Sigma_{10}$, similar to \\cite{dressler1980b}, \\cite{lewis2002} and \\cite{gomez2003}.  To calculate the SFR, \\citet{hopkins2001}, \\citet{cardiel2003}, \\citet{lewis2002} and \\citet{gomez2003} use OII, OIII, H$\\alpha$, H$\\beta$, H$\\delta$ and UV detected fluxes, however these bands require correction for dust extinction.  Infrared and radio detections, both indicators of current star formation, are not affected by dust \\citep{hopkins2003, cardiel2003} and studies show a correlation exists between IR and 1.4-GHz SFRs \\citep{hopkins2002, cardiel2003, yun2001}. In our study we use 843-MHz and 1.4-GHz radio, and the 12, 25, 60 and 100 $\\mu$m infrared (IR) detected fluxes to calculate the SFR per HI galaxy.  The star formation efficiency (SFE) is used to determine the efficiency with which the HI present is converted into stars per year. From the HI galaxy SFR, SFE and $\\Sigma_{10}$ we  explain the reason behind the SFR-density relation.  In Section \\ref{Sec:Paper2LocalSurfaceDensity} we describe the process taken to determine the local surface density for each HOPCAT target galaxy.  Section \\ref{Sec:Paper2SFRSFE} details the resources we use to obtain the fluxes and the equations to calculate the SFRs and SFEs. We present our results in Section \\ref{Sec:Paper2Results} and discuss these results in Section \\ref{Sec:Paper2Discussion}.  Section \\ref{Sec:Paper2Summary} summarises our work. We use H$_{\\circ}$=71 km s$^{-1}$ Mpc$^{-1}$ throughout. ",
        "conclusions": "\\label{Sec:Paper2Discussion} In this section we discuss the relationships between M$_{HI}$ and SFR, SFR and $\\Sigma_{10}$, and SFE and $\\Sigma_{10}$. We also compare optical and HI detected galaxy samples. From these results we determine which of our two theories explain the SFR-density relation.  \\subsection{HI mass-SFR Relationship \\label{SubSec:Paper2DiscussMassHISFR}}  We investigate the relationship between the SFR and M$_{HI}$ HOPCAT galaxies with IR and 1.4-GHz SFRs.  We find M$_{HI}$ increases with increasing SFR.  This result is in general agreement with the Schmidt law, a relationship between a galaxy's gas surface density and the disk average SFR as discussed in \\cite{kennicutt1998}.  \\citet{mirabel1988}, investigating 92 luminous infrared galaxies find that 88 galaxies have M$_{HI}$ in the range of $5\\times10^8$ to $3\\times10^{10} M_{\\odot}$. They conclude the FIR luminosities have only a loose correlation with the M$_{HI}$. However, our study looks at HI galaxies and calculates SFR from IR and 1.4-GHz fluxes and finds a strong correlation between the SFR and M$_{HI}$.  \\cite{stevens2005} also investigated the relationship between M$_{HI}$ and SFR using a volume limited subsample of HICAT, containing 85 FIR and 277 radio (NVSS and SUMSS) based SFRs. \\cite{stevens2005} also finds M$_{HI}$ increases with increasing SFR.  Although we do not know where star formation occurs in a galaxy with respect to the HI detection, the correlation between the M$_{HI}$ and SFR shows HI is an indication of star formation.  In fact, as our SFRs are from IR and 1.4-GHz fluxes, which are indicators of current star formation, the M$_{HI}$ and SFR correlation suggests HI indicates `current' star formation in galaxies.  \\subsection{HI radio and Optical galaxy samples \\label{SubSec:Paper2DiscussCompareHiOp}}  Our result from investigating the difference between optical and HI detected galaxy samples show there is a much larger population of elliptical galaxies than HI detected galaxies with increasing local surface density. This result agrees with studies into the morphology-density relation (\\citet{dressler1980b} etc.), which find for an increasing local galaxy density there is a decrease in the number of spiral galaxies and an increase in the number of elliptical galaxies.  It is expected that this decrease in the number of current star forming HI galaxies with increasing $\\Sigma_{10}$ will have a direct effect on the results from any SFRs versus $\\Sigma_{10}$ analysis for optical samples.  \\subsection{SFR and $\\Sigma_{10}$ \\label{SubSec:Paper2DiscussSFRSigma10}}  From our investigation into SFR and $\\Sigma_{10}$ we conclude the SFRs for our HI galaxies are not affected by increasing $\\Sigma_{10}$.  As HOPCAT is based on the Parkes HI All Sky Survey (HIPASS) there may be some concern about the number of low surface brightnes (LSB) galaxies in the sample.  The HIPASS galaxy sample is dominated by late-type galaxies as is expected due to the HI selection (e.g. Helmboldt et al. 2004). As \\citet{helmboldt2004} show, HI-selected galaxies tend to have lower luminosity and surface brightness than a comparison sample of optically-selected galaxies. This difference is not so great for our HOPCAT sample because our mean redshift is much greater than the Helmboldt sample. Our median luminosity is $M_B=-18.8$ mag compared to $M_B=-18.5$ mag in their HI-selected sample and  $M_B=-19.1$ mag in their optical sample. Some 15 per cent of their HI-selected galaxies have LSB compared to the optical sample: the fraction of LSB galaxies in HOPCAT will be smaller because of the higher median luminosity.  If we consider the final HOPCAT sample with reliable densities and star formation rates used in our analysis, only 1.3 per cent of the galaxies are fainter than the M$_{B_j}$=-17 mag faint absolute magnitude limit applied in the density calculations.  In \\citet{gomez2003} and \\citet{lewis2002} the decrease in the average SFR with an increase in local surface density, for a given density range, is explained by their optical sample.  An optical sample contains both spiral and elliptical galaxies. For areas of increasing local galaxy density, such as compact galaxy groups and poor and rich clusters, the number of spiral galaxies decreases as the number of elliptical galaxies increases.  Studies into the morphology-density relation (\\citet{dressler1980b} etc.), show this change in morphology type for increasing local galaxy density. Elliptical galaxies have little or no current star formation, hence the average SFR, as \\citet{gomez2003} and \\citet{lewis2002} use, will decrease in areas of increasing local surface density where elliptical galaxy numbers increase and spiral galaxy numbers decrease.  As our sample does not have elliptical galaxies and the SFRs are for galaxies with current star formation, a more direct picture of SFR with respect to local surface density is seen.  We clearly show the HI galaxies' SFRs are relatively constant regardless of local surface density.  \\citet{gomez2003} and \\citet{lewis2002} describe a characteristic $\\Sigma_{10}$ break at 1 gal Mpc$^{-2}$ where the average SFR decreases sharply. As our scatter plots are a direct comparison of SFR per HI galaxy for each galaxies local surface density, using lines of best fit, no break is seen.  However, at 5 gal Mpc$^{-2}$ (right panel of Fig. \\ref{Fig:Paper26dFK6dfCompareHopcatK6dFSepHists}), the increasing gap indicates a decrease in the number of HI galaxies relative to optical galaxies.  \\citet{lewis2002} and \\citet{gomez2003} use a faint absolute magnitude limit of -19 mag and we use -17 mag which results in our larger $\\Sigma_{10}$ values. Our decrease in HI galaxy numbers at $\\Sigma_{10}$=5 gal Mpc$^{-2}$ may coincide with \\citet{gomez2003} and \\citet{lewis2002} characteristic $\\Sigma_{10}$ break at 1 gal Mpc$^{-2}$.  \\citet{stevens2005}, also using HICAT, finds the SFR for galaxies in groups is smaller than field galaxies. This differs from the results presented in this paper for the same sample. However, \\cite{stevens2005} uses a group finder algorithm to find HI galaxies in groups (143) and compares their SFR to those for field galaxies (17). As the resulting density is 3D and the analysis \\cite{stevens2005} uses is significantly different from ours, no direct comparison is possible.  However, one observation \\cite{stevens2005} makes is there is a larger population of red galaxies in groups than in the field. He concludes this could partially explain the lower SFR in group galaxies, which agree with our results.  It is clear the lower average SFR in regions of high local surface density in \\citet{lewis2002} and \\citet{gomez2003} can be attributed to their larger population of elliptical galaxies.  \\citet{balogh2004} using H$\\alpha$ equivalent widths to investigate the population of star-forming galaxies in SDSS and 2dFGRS and finds `the distribution of line strength' for these star-forming galaxies is independent of environment.  This appears to be in agreement with our results.  Our conclusion still remains, for HI galaxies between 1000 and 6000 km s$^{-1}$ in the southern sky, the SFR does not depend on the local galaxy density.  \\subsection{Fewer Star Forming Galaxies or Suppressed Star Formation?}  In Section \\ref{SubSec:Paper2DiscussMassHISFR} we show a strong correlation exists between M$_{HI}$ and SFR leading to the conclusion our HI galaxies display current star formation. As discussed in Section \\ref{SubSec:Paper2DiscussCompareHiOp}, the number of HI galaxies decreases sharply in regions of high local surface density. The difference between this work and \\citet{gomez2003} and \\citet{lewis2002} is our sample contains current star forming galaxies and theirs are optical galaxy samples which contain galaxies with little or no current star formation.  We also find in Section \\ref{SubSec:Paper2DiscussSFRSigma10}, the SFRs do not depend on $\\Sigma_{10}$. Our results in Section \\ref{SubSec:Paper2ResultsSFESigma10} show the efficiency, with which HI is converted to stars, is not affected by an increase in the local surface density.  We now conclude which theory explains the SFR-density relation.  We find there are, indeed, fewer star forming (HI) galaxies in regions with high local surface density, which affects the average SFR for optical galaxy samples. As the SFR and SFE do not depend on the local surface density, the SFR-density relation is due to the low number of star forming (HI) galaxies, in regions of high local surface density and is not due to the suppression of star formation in these regions.   We use an optically matched HI detected galaxy catalogue, HOPCAT, to investigate SFR, SFE and the local surface density to explain the SFR-density relation. In regions of high local galaxy density either, there are fewer star forming galaxies, or the galaxies capable of forming stars are present but some physical process is suppressing their star formation. The local surface density, $\\Sigma_{10}$, is calculated using HOPCAT with SCosGS and 6dF2MASS as background catalogues. {\\it{IRAS}}, SUMSS and NVSS fluxes are used to calculate the SFRs and SFEs. For the SFR-M$_{HI}$ investigations we use 2 different data sets using 6dF2MASS with IR and 1.4-GHz SFRs, HOPCAT-6dF2MASS-IR (1405 galaxies) and HOPCAT-6dF2MASS-1.4GHz (561 galaxies). For the SFR-SFE-$\\Sigma_{10}$ investigations we have 4 different data sets using SCosGS and 6dF2MASS each with IR and 1.4-GHz SFRs, HOPCAT-SCosGS-IR (923 galaxies), HOPCAT-6dF2MASS-IR (1180 galaxies), HOPCAT-SCosGS-1.4GHz (339 galaxies) and HOPCAT-6dF2MASS-1.4GHz (478 galaxies). We find:  1. A strong correlation exists between M$_{HI}$ and SFR for galaxies out to 10000 km s$^{-1}$ in the southern sky. The SFRs are calculated using IR and 1.4-GHz fluxes, both of which are indicators of current star formation.  We conclude HI galaxies display current star formation and their M$_{HI}$ and SFR are strongly related.  2. The number of HI galaxies decreases with increasing local surface density. This agrees with \\citet{dressler1980b}, \\citet{hashimoto1998}, \\citet{martinez2002}, \\citet{goto2003}, \\citet{sodre2003}, \\citet{nuijten2005} and \\citet{christlein2005}, where optical samples show an increase in the number of non-current star forming elliptical galaxies with a decrease in current star forming galaxy numbers for increasing local galaxy density. This is also seen in \\citet{balogh2004}.  3. For our HI galaxies, the SFR does not depend on local surface density.  Any decrease in the average SFR with increasing local surface density for optical galaxy samples, is explained by the decrease in the number of current star forming galaxies and an increase in elliptical galaxies in regions with increasing local surface density.  4. The efficiency with which the HI present in the HI galaxies is converted into stars is the same regardless of local surface density.  5. The SFR-density relation is due to a decrease in the number of star forming (HI) galaxies in regions of high local (surface) galaxy density and not due to the suppression of star formation in these regions."
    },
    "0608/astro-ph0608599_arXiv.txt": {
        "abstract": "We study the 4--200 keV spectral and temporal behaviour of the low mass X-ray binary \\source\\ with {\\it INTEGRAL\\/} during 2003--2005. This source as been observed in both the soft (banana) and hard (island) spectral states. A high energy tail, above 50 keV, in the hard state has been observed for the first time. This places the source in the category of X-ray bursters showing high-energy emission. The tail can be modeled as a soft power law component, with the photon index of $\\simeq$2.4, on top of thermal Comptonization emission from a plasma with the electron temperature of $kT_{\\rm e}\\simeq 6$ keV and optical depth of $\\tau\\simeq 4$. Alternatively, but at a lower goodness of the fit, the hard-state broad band spectrum can be accounted for by emission from a hybrid, thermal-nonthermal, plasma. During this monitoring the source spent most of the time in the soft state, usual for this source, and the $\\ga$4 keV spectra are represented by thermal Comptonization with $kT_{\\rm e}\\simeq 3$ keV and $\\tau \\simeq 6$--7. ",
        "introduction": "4U 1820--30 is a low mass X-ray binary seen at 0.66$^{\\prime\\prime}$ from the centre of the globular cluster NGC 6624. It was the first identified source of type-I X-ray bursts \\citep{grindlay}. Distance has been estimated as $d\\simeq 5.8$--8.0 kpc (e.g., Shaposhikov \\& Titarchuk 2004; Kuulkers et al.\\ 2003). Kuulkers et al.\\ (2003) found their determination of $d=7.6\\pm 0.4$ kpc compatible with the peak X-ray burst luminosity being a standard candle in their sample of sources, and equal to the Eddington limit for He. The binary system consists of a He white dwarf with mass of 0.06--$0.08 M_{\\sun}$ \\citep{rappa} and a neutron star, with the mass estimated by Shaposhnikov \\& Titarchuk as $\\sim\\! 1.3M_{\\sun}$, orbiting at short period of 11.4 minutes \\citep{Stella}. In X-rays, \\source\\ is classified as an atoll \\citep{hasinger}. However, its flux variation between the soft (banana) to the hard (island) state are quasiperiodic at $\\sim$170 d (Priedhorsky \\& Terrell 1984; {\\v S}imon 2003; Wen et al.\\ 2006), which quasiperiodicity has been proposed to be due to tidal effects of a more remote third star (Chou \\& Grindlay 2001; Zdziarski, Wen \\& Gierli\\'nski 2006). X-ray bursts occur only at low flux levels (Cornelisse et al.\\ 2003) and the kHz QPOs are correlated with the flux (Zhang et al.\\ 1998), proving that the variability is due to an intrinsic change of the source luminosity (and the accretion rate) rather than being due to variable either obscuration or viewing angle of an anisotropic emitter. X-ray spectra of the source from {\\it RXTE\\/} and {\\it BeppoSAX\\/} were fitted by thermal Comptonization with or without an additional blackbody (Bloser et al.\\ 2000; Sidoli et al.\\ 2001; Migliari et al.\\ 2004).  In this paper, we report the first detection of X-ray emission above 50 keV from \\source\\ in a hard spectral state. We studied the source in the 4--200 keV energy range with the JEM-X \\citep{lund} and IBIS \\citep{uber} instruments on board the {\\it INTEGRAL\\/} satellite \\citep{wink}, collected during the Galactic Centre Deep Exposure (GCDE) program \\citep{wink2}. We combined three methods of analysis: {\\it Temporal}, using the light curves in different energy bands; {\\it Photometric}, with the Hardness--Instensity Diagram; and {\\it Spectral}, by spectral fitting. ",
        "conclusions": "We have presented the first detection of hard emission above 50 keV in the island state of the atoll source \\source\\ using IBIS. Up to now, \\source\\ has not been included in the sample of the high-energy emitting bursters \\citep{bazzano} and indeed neither {\\it BeppoSAX}, {\\it CGRO}/BATSE  nor {\\it RXTE\\/} detected this source above 50 keV (Bloser et al.\\ 2000 and references therein). The previous lack of detection might have been due to either the source spending most of the time in the soft state, the weakness of the emission above 50 keV in the island state, or the strong emission at $\\ga 50$ keV appearing in that state only occasionally.  In our island-state observation, the relative flux contribution for energies $>60$ keV is $\\sim$10\\%. The ratio of the 4--20 keV flux to the 20--120 keV ones is  $\\sim$2, whereas it is in the range 20--27 for the soft (banana) state. We point out that we detected emission above 50 keV when the 1--20 keV luminosity was $\\simeq 1.4 \\times 10^{37}(d/5.8\\,{\\rm kpc})^2$ erg s$^{-1}$, i.e., close or somewhat below the critical value at which the X-ray binaries become hard emitters as proposed by Barret et al.\\ (1996, 2000).   While the origin of the thermal Comptonization component in the soft state appears well understood as emission from hot electrons that upscatter soft photons coming from the accretion disk and the neutron star surface, the origin of the high-energy tail is less clear.  The radio emission in the X-ray binaries is suggested to originate from jet-like outflows which may also produce X-rays \\citep{fender04}. In particular, there are indication that both black-hole candidate and neutron stars sytems, when in hard state, produce steady jet (Migliari \\& Fender 2006). Also, in the hard state, positive correlation between radio and X-ray fluxes has been reported (Migliari \\& Fender 2006). \\source\\ is one of the few radio detected atoll source. The only radio emission reported so far was during the soft state \\citep{migliari}, but it is probable that radio emission exist also in the hard state, as e.g. for 4U 1728-34 \\citep{mi03}. Furthermore, the power-law tail detected by us is relatively soft, with $\\Gamma\\simeq 2.4$, and dominates over the Comptonization emission already for $E\\la 4$ keV, as shown in Fig.\\ \\ref{plotfour}. Integration to its minimum energy, $E_{\\rm min}$, yields the luminosity in this component $\\propto E_{\\rm min}^{2-\\Gamma}$. If this component is due to nonthermal synchrotron emission, $E_{\\rm min}$ corresponds to the self-absorption turnover, typically in the IR or optical range, as it is the case in another ultracompact binary, 4U 0614+091 (Migliari et al.\\ 2006). Then the tail luminosity would be $\\sim$10 times that of X-rays (see also discussion in Zdziarski et al.\\ 2003), making the island state more luminous than the banana state, with the maximum of the overall $EF_E$ at $\\sim$1 eV, neither of those seen in X-ray binaries.  Also, the presence of the jets or outflow could be due to magnetic reconnection avents above the accretion disk (Shibata 2005).  A more likely origin of the power law tail appears to be emission of nonthermal electrons in a hybrid (i.e., containing also thermal electrons) plasma, as discussed in Section \\ref{spectral}. Nonthermal tails have been commonly seen in soft states of black-hole binaries (e.g., Cyg X-1, Gierli\\'nski et al.\\ 1999) and neutron-star Z-sources, e.g., GX~17+2 (Di Salvo et al.\\ 2000; Farinelli et al.\\ 2005), Sco~X-1 \\citep{damico}, GX~349+2 \\citep{disalvo_b}, GX~5--1 \\citep{asai}, while they are difficult to detect in the atoll sources, e.g., 4U 0614+091 \\citep{pira}. On the other hand, nonthermal tails have been found in the hard state of some black-hole binaries, e.g. Cyg X-1 (McConnell et al.\\ 2002) and GX 339--4 (Wardzi\\'nski et al. 2002).  The \\source\\ high-energy part of the spectrum (E $\\ga 10$ keV) can be well fitted with purely thermal Comptonization at $kT_{\\rm e}\\simeq 50$ keV, as discussed in Section \\ref{spectral}. This electron temperature is higher than that seen in 4U 1608--52 (Gierli\\'nski \\& Done 2002), but the Comptonization temperature of atolls in the island state has not yet been determined for a sample of sources, and $kT_{\\rm e}\\simeq 50$ keV is similar to that seen in the hard state of black-hole binaries (e.g., Zdziarski \\& Gierli\\'nski 2004). Then, the problem with fitting the spectrum at $<10$ keV with this model may be due to the complexity of the X-ray spectra in the hard state, including contributions from the disk, the corona above the disk, optically-thin and optically-thick parts of the boundary layer and the uncovered part of the neutron star surface.  A different thermal Comptonization model has been considered by Markoff et al.\\ (2005). They proposed thermal synchrotron self-Compton emission in a jet as a model for the hard state of black hole binaries. However, their model requires an extremely high electron temperature, $kT_{\\rm e}\\sim 3$--4 MeV. This would result in the thermal spectral cutoff in black hole binaries at energies $\\ga 3kT_{\\rm e}\\sim 10$ MeV, which has not been seen in {\\it any\\/} object yet. Although we cannot formally rule out a contribution from this process, postulating the presence of $\\sim$3 MeV plasma in 4U 1820--30 appears rather speculative at this time.  In summary, we have discovered strong high-energy emission from 4U 1820--30. The high energy emission is harder than that from some other atolls in the island state. Its most likely interpretation is either emission of some nonthermal electrons or thermal emission from plasmas with a relatively high temperature, $kT_{\\rm e}\\ga 50$ keV."
    },
    "0608/astro-ph0608550_arXiv.txt": {
        "abstract": "Based on a new approach to the detection of radio transients associated with extensive air showers induced by ultra high energy cosmic rays, the experimental apparatus CODALEMA is in operation, measuring about 1 event per day corresponding to an energy threshold $\\sim 5\\times 10^{16}$~eV. Its performance makes possible for the first time the study of radio-signal features on an event-by-event basis. The sampling of the magnitude of the electric field along a 600~meters axis is analyzed. It shows that the electric field lateral spread is around 250~m (FWHM). The possibility to determine with radio both arrival directions and shower core positions is discussed. ",
        "introduction": "Radio emission associated with the development of Extensive Air Showers (EAS) was investigated in the 1960's~\\cite{ask,weekes}. A flurry of experiments provided initial information about signals from $10^{17}$~eV cosmic rays~\\cite{allan}, but plagued by difficulties (poor reproducibility, atmospheric effects, technical limitations) efforts almost ceased in the late 1970's while ground particle~\\cite{agasa} and fluorescence~\\cite{fly} detector work continued.  With the growing interest for ultra high-energy cosmic ray (UHECR) research involving giant surface detectors~\\cite{auger}, radio detection however appears as a promising tool for future apparatus considering its specific advantages: low-cost, high duty cycle and sensitivity to the longitudinal development of the showers.  Although the rebirth of radio pulse investigation~\\cite{casa-mia,cascade-grande,gorham,sub,rice} is recent, first available results~\\cite{nature,nim-ard} demonstrate the feasibility of EAS radio detection.  The CODALEMA (COsmic ray Detection Array with Logarithmic Electro-Magnetic Antennas) experiment, located at the Nan\\c cay radio observatory~\\cite{DAM}, is a part of this new effort. Its originality lies in completely recording the form of the transient radio signals together with sampling the electric field over a few hundred meters range on the ground. This makes possible both the determination of arrival times and directions of radio pulses and the study of their field amplitude impact parameter dependences. The measurement of UHECR around $10^{20}$ eV is the admitted goal of such experimental development, but a proof-of-principle demonstration at such a large energy would suffer from a lack of statistics without an extensive antenna array. This point can be circumvented in a first stage by working around $10^{17}$~eV where a measurable signal amplitude~\\cite{allan,huege} is expected at not too large distances from the shower core. Considering a vertical shower falling upon the detector, the predicted transient should reach 150 $\\mu$V/m with a 10 ns FWHM duration  for $10^{17}$~eV cosmic rays~\\cite{allan,casa-mia,nim-ard}. Transposed in the frequency domain, the corresponding pulse spectrum should extend from 1 to 100 MHz. Thus, with a wide band antenna it should be possible to recover the original pulse shape, allowing for energy determination and providing information on the nature of the primary particle with minimal assumptions concerning its electromagnetic shower- signature. The design of CODALEMA is based on these expectations.  CODALEMA has already provided firm evidence for a radio emission counterpart of EAS with an estimated energy threshold of $5\\times 10^{16}$~eV~\\cite{nim-ard}. In the present paper, we extend the characterization of EAS candidates giving special attention to the electric field pattern observed on an event-by-event basis. The experimental set-up is described in section~2 and general event properties are given in section~3. The selection of EAS radio candidates is discussed in section~4. Section~5 is the central part of the present study and details first observations on EAS electric field lateral distributions from a few illustrative examples. Section~6 discusses frequency dependences. Some conclusions are given in section~7. ",
        "conclusions": "Features of the electric field transients generated by more than a hundred extensive air shower events have been observed with CODALEMA. Through the evolution of this sample of events, the observed characteristics allow the determination of the EAS core location together with the electric field magnitude and spread, on an event-by-event basis using an electric field profile function. Electric field profiles show slopes and amplitudes which are variable between events. These patterns offer possibilities to discriminate between EAS events and radio frequency interferences. Frequency dependence of EAS electric field has been also observed. Such detailed electric field correlations or characteristics have been predicted to be related to important physical quantities such as the EAS energy, the nature of their primaries and various important shower evolution parameters~\\cite{allan,huege}. The present results show for the first time the feasability of such studies on an event-by-event basis using a radio detection set-up in combination with particle detectors. More complete works and higher statistics are clearly necessary to establish the exact nature of the physical correlations detectable by the measurements of the radioelectric field features.  In addition, the various sources that can influence the electric field distribution in a given event are hard to disentangle with the limited set-up described here. An upgrade with new particle detectors is already underway to provide more information on each event. They will give an independent determination of the core position and an estimate of the shower energy, invaluable information at the present stage of development of the radio technique.  It has been emphasized that the observed field pattern on the various antennas constitutes a clear ``radio'' signature. This suggests that it may be possible to discriminate an EAS event from a fortuitous one, using a self triggering array of radio antennas. This is one further step toward a stand-alone system that could be deployed over a large area or added to an existing surface detector such as the Pierre Auger Observatory. The radio signals could then provide complementary information about the longitudinal development of the shower, as well as the ability to lower the energy threshold. More data and technical upgrading are planned in order to examine the contribution that the radio detection method could bring to the determination of the energy and nature of ultra high energy cosmic rays.  The fact that inclined or horizontal radio waves could be efficiently detected with our set-up may be of great interest to EAS detection if the electric field patterns discussed in the present study turn out to allow for a clear discrimination between EAS radio events and other noise transients close to the horizon. Among various particles able to generate air showers, neutrinos are thought to probably be the only ones which could address the nature and source location of ultra energetic cosmic rays. Taking into account the sensitivity of the radio detection method to inclined extensive air showers~\\cite{gousset}, the presented results could considerably enlarge the scope of such studies for cosmology related phenomena.  \\appendix"
    },
    "0608/astro-ph0608285_arXiv.txt": {
        "abstract": "\\Clover\\ is an experiment which aims to detect the signature of gravitational waves from inflation by measuring the B-mode polarization of the cosmic microwave background.  \\Clover\\ consists of three telescopes operating at 97, 150, and 220 GHz.  The 97-GHz telescope has 160 feedhorns in its focal plane while the  150 and 220-GHz telescopes have 256 horns each.  The horns are arranged in a hexagonal array and feed a polarimeter which uses finline-coupled TES bolometers as detectors.  To detect the two polarizations the 97-GHz telescope has 320 detectors while the 150 and 220-GHz telescopes have 512 detectors each. To achieve the target NEPs (1.5, 2.5, and $4.5\\times10^{-17}\\rm\\ W/\\sqrt Hz$) the detectors are cooled to 100~mK for the 97 and 150-GHz polarimeters and 230~mK for the 220-GHz polarimeter.  Each detector is fabricated as a single chip to ensure a 100\\%\\ operational focal plane.  The detectors are contained in linear modules made of copper which form split-block waveguides.  The detector modules contain  16 or 20 detectors each for compatibility with the hexagonal arrays of horns in the telescopes' focal planes.  Each detector module contains a time-division SQUID multiplexer to read out the detectors.  Further amplification of the multiplexed signals is provided by SQUID series arrays.  The first prototype detectors for \\Clover\\ operate with a bath temperature of 230~mK and are used to validate the detector design as well as the polarimeter technology.  We describe the design of the \\Clover\\ detectors, detector blocks, and readout, and present preliminary measurements of the prototype detectors' performance. ",
        "introduction": "\\label{sect:intro}  %  \\subsection{Scientific motivation} \\label{sect:motivation} \\PARstart{T}{homson} scattering of primaeval radiation in the early Universe can lead to linear polarization\\cite{Rees68} in the cosmic microwave background (CMB).  The polarization depends on density fluctuations, and thus carries cosmological information which is complimentary to the well-studied temperature anisotropies of the CMB.  The linear polarization may be decomposed into a curl-free part and a divergence-free part, denoted E- and B-mode respectively, by analogy with the electric field strength $E$ and magnetic induction $B$.  Linear density perturbations do not produce B-mode polarization, while tensor perturbations, as might be produced by gravitational waves, produce E- and B-mode polarization with similar amplitude\\cite{Selj97,Kami97}.  Thus, by measuring the B-mode polarization of the CMB with \\Clover\\ we hope to make an indirect detection of a background of primordial gravitational waves.  \\subsection{Overview of \\Cloverit} \\label{sect:overview} The \\Clover\\ experiment is described in detail elsewhere\\cite{Call06}. \\Clover\\ consists of three telescopes observing at frequencies of 97, 150, and 220~GHz.  The focal plane of each telescope will be populated by feed horns, each connected to a polarimeter.  The polarimeter technology is yet to be determined.  As part of the technology development program for \\Clover\\ we are constructing a Single Pixel Demonstrator.  This instrument will contain six TES detectors, allowing us to evaluate different polarization technologies, as well as validating the detector design. ",
        "conclusions": "We have developed a process that allows us to mass-produce finline-coupled TES bolometers.  This is the first time TES detectors have been mass-produced in the United Kingdom.  We have produced prototype detectors for \\Clover\\ and found that the TES films are of high quality, making for sensitive detectors.  More development of the thermal design is needed to achieve the required power handling.  We expect to measure the R.F. response of these detectors in the near future, demonstrating the operation of a finline transition on a silicon substrate for the first time."
    },
    "0608/astro-ph0608416_arXiv.txt": {
        "abstract": "We present a study of the iron abundance pattern in hot hydrogen-rich (DA) white dwarfs. The study is based on new and archival far ultraviolet  spectroscopy of a sample of white dwarfs in the temperature range $30,000 {\\rm K} \\la T_{\\rm eff} \\la 64,000$K.  The spectra obtained with the {\\it Far Ultraviolet Spectroscopic Explorer} along with spectra obtained with the {\\it Hubble Space Telescope}  Imaging Spectrograph and the {\\it International Ultraviolet Explorer}  sample \\ion{Fe}{3} to \\ion{Fe}{6} absorption lines enabling a detailed iron abundance analysis over a wider range of effective temperatures than previously afforded. The measurements reveal abundance variations in excess of two orders of magnitude between the highest and the lowest temperatures probed, but also show considerable variations (over one order of magnitude) between objects with similar temperatures and surface gravities. Such variations in cooler objects may be imputed to accretion from unseen companions or so-called circumstellar debris although the effect of residual mass-loss and selective radiation pressure in the hottest objects in the sample remain dominant. ",
        "introduction": "The cooling history of white dwarfs involves phases of chemical transformation of their atmospheres. One such transformation occurs early on the hydrogen-rich (DA) cooling sequence. Far ultraviolet (FUV) spectroscopy of hot, hence young DA white dwarfs show that iron is a dominant source of atmospheric opacity \\citep{ven1992, hol1993}. The effect of iron-group opacities has also been observed in extreme ultraviolet (EUV) spectra of the same stars as well as in other objects with similar parameters \\citep{dup1995,wol1998,sch2002}. Before reaching a cooling age of a few million years ($t_{cool} \\leq 2-3\\times 10^6$ years) this source of opacity vanishes in most DA white dwarfs, and the metallicity (i.e., $n(Fe)/n(H)$ by number) drops from nearly solar values \\citep{ven2001} to less than $10^{-2} (Fe/H)_{\\odot}$, i.e., below the detection threshold of high-dispersion {\\it International Ultraviolet Explorer} ({\\it IUE}) or even Space Telescope Imaging Spectrograph (STIS) high-resolution spectra \\citep{bar2003}. However, the absence of an effective iron abundance diagnostics in cooler white dwarfs ($20,000{\\rm K} \\la T_{\\rm eff} \\la 40,000$K, or $t_{cool} = 80-4 \\times10^6$ years), reserved detailed iron abundance analyses to a few young objects ($t_{cool} \\leq 2-3\\times 10^6$ years). A spectral window limited to $\\lambda > 1150$\\AA\\ restricts the iron abundance analysis to  \\ion{Fe}{4}, \\ion{Fe}{5}, and \\ion{Fe}{6} species which are dominant in white dwarfs with  $T_{eff} \\ga 40,000$K, and to \\ion{Fe}{2} represented by a strong doublet at $\\lambda\\lambda$ 1260.533/1267.422\\AA\\ and dominant in cool white dwarfs ($T_{eff} \\la 20,000$K). However, in the temperature range $20,000-40,000$ K the ionization balance shifts toward \\ion{Fe}{3}. Spectral lines of \\ion{Fe}{3} are prominent near $\\lambda \\approx 1125$ \\AA\\  in a spectral window made accessible by the {\\it Far Ultraviolet Spectroscopic Explorer} ({\\it FUSE}).  The new possibilities offered by {\\it FUSE} were demonstrated in a study of the DA white dwarf GD 394 ($T_{eff}\\approx37,000$ K). The presence of trace elements in the otherwise hydrogen-dominated atmosphere of GD 394 had been established, among others, by \\citet{pae1989} in their analysis of EUV photometry of several hot white dwarfs observed with {\\it EXOSAT}. The source of opacity remained uncertain \\citep[see][]{wol1998, dup2000} until \\citet{cha2000} measured a surprisingly high abundance of iron in their analysis of new {\\it FUSE} spectra.  They used several lines of \\ion{Fe}{3} to determine an iron abundance almost solar ($Fe/H \\sim 5\\times 10^{-6}$). Such a high abundance goes against the trend inferred in the study of hotter white dwarfs and indicate a resurgence of heavy elements in the atmosphere of aging DA white dwarfs. Since \\citet{cha1995} demonstrated that the abundance of iron supported by selective radiation pressure decreases well under a solar abundance at $T_{eff} \\la 50,000$ K,  external causes are suspected, in particular accretion from the ISM, a close companion, or from a circumstellar shell.  To investigate the question further, we initiated a study of FUV spectra of a sample of twelve DA white dwarfs  observed with {\\it FUSE},  {\\it IUE}, and STIS. The sample covers a range of effective temperatures from $\\le 64,000$K down to $\\ge 30,000$K. We describe the sample selection process  in \\S 2, with relevant EUV/soft X-ray data listed in Appendix~A, and we present  the FUV observations in \\S 3. The model atmospheres and spectral synthesis used in this investigation are presented in \\S 4, and our analysis is presented in \\S 5. Finally, we summarize and discuss our results in \\S 6. ",
        "conclusions": "We have presented a study of the iron abundance pattern in hot DA white dwarfs. We reported the discovery of \\ion{Fe}{3} spectral lines in a {\\it FUSE} spectrum of the DA white dwarf GD~683 and a corresponding abundance $\\log{\\rm Fe/H}=-6.9$. The discovery of iron in the photosphere of this star confirms our suspicion that heavy elements, and iron in particular, are responsible for lingering EUV/soft X-ray flux deficit in DA white dwarfs with temperatures as low as $\\approx$30,000K.   The detection of iron in the DA white dwarf 1819+580 is particularly intriguing. \\citet{gre2000} uncovered substantial near infrared excess from this star which  they attributed to a dM6 companion, but there is insufficient information to establish its proximity to the white dwarf.  In fact, with the exception of Feige~24 and 1819+580 we do not find  clear evidence of an infrared flux excess in any of the sample stars. We obtained JHK photometry from the 2MASS database available at the Centre de Donnees astronomique de Strasburg and V magnitudes from \\citet{mcc1999}, and we computed $V-J$ indices for all objects in the sample. The index ranges from $V-J=-0.72\\pm0.10$ (GD~683) to $-0.96\\pm0.15$ (2000$-$561) with the exception of 1819$+$580 ($V-J=-0.11\\pm0.10$) and Feige~24 ($V-J=1.15\\pm0.10$). Synthetic $V-J$ indices for DA white dwarfs range from $-0.7$ to $-0.8$ for white dwarf effective temperatures between 30,000K and 50,000K. Allowing for large uncertainties in 2MASS measurements, it appears that only 1819$+$580 and Feige~24 show evidence of a companion in infrared measurements. The estimated absolute J magnitudes are $M_J=11.43$ and 7.06  corresponding to M8.5 and M2.5 spectral types \\citep{kir1994} for 1819$+$580 and Feige~24, respectively. The spectral type of the companion of 1819$+$580 is somewhat later than estimated by \\citet{gre2000} because we adopted a higher gravity, hence lower luminosity, for the white dwarf itself. From our sample, only Feige~24 has a known {\\it close} companion ($P=4.2$ days), but its abundance pattern is indistinguishable  from its sibling and apparently single white dwarf G191-B2B \\citep{ven2001}. In all other cases, predicted $3\\sigma$ upper limits to the companion absolute J magnitude range from $M_J \\la10.3$ or a spectral type earlier than M6 \\citep{kir1994} to $M_J\\la 13.3$ or a  spectral type earlier than L5 \\citep{kir2005}. These hypothetical stellar and  substellar  companions may be responsible for the presence of iron in white dwarfs such as GD~683 and GD~394. Moreover, the variable EUV emission from GD~394 may imply that iron, silicon and other elements are not uniformly distributed over the surface, which may be the result of episodic accretion from a putative companion \\citep{dup2000}.  The iron abundance in DA white dwarfs with temperatures between $\\approx30,000$ and $\\approx 64,000$K does not correlate with effective temperatures in a simple manner.  The sample of DA white dwarfs is part of a EUV/soft X-ray selected population and may be biased against ultra-high metallicity objects, although the largest iron abundance measured is essentially solar. Other non-EUV selected white dwarfs were included in a study by \\citet{bar2003} such as PG~0948$+$534, Ton~21, and PG~1342$+$444 which, although hotter than stars from the present study, may have lower iron abundances. Other EUV-selected DA white dwarfs such as 1RXS~J0619.1$-$0828 \\citep{ven1999} lack FUV data required for an abundance analysis. Hopefully future missions may help enlarge this sample of hot DA white dwarfs.  In summary, our measurements reveal abundance variations in excess of two orders of magnitude between the highest and the lowest temperatures probed, but also show considerable variations (over one order of magnitude) between objects with similar temperatures and surface gravities. Such variations in cooler objects may be imputed to accretion from unseen companions or so-called circumstellar debris \\citep{kil2005} although the effect of residual mass-loss and selective radiation pressure in the hottest objects in the sample may remain dominant. Accretion from a close companion remain a likely source of heavy elements in white dwarf photospheres. A deeper search for infrared excess, revealing the presence of a companion or debris near the white dwarf, may offer clues as to the source of iron in hot DA photospheres.  Guided by present and previous results we will embark on a re-analysis of the complete spectral energy distribution of all hot DA white dwarfs lifting the often convenient but possibly inappropriate assumption of a pure hydrogen composition. In particular, the present analysis of \\ion{Fe}{3} spectral lines, and its sensitive iron abundance diagnostics, will be extended to the complete {\\it FUSE} spectroscopic data set."
    },
    "0608/astro-ph0608351_arXiv.txt": {
        "abstract": "Recently, the Heidelberg-Moscow double beta decay experiment has claimed a detection for a neutrino mass with high significance. Here we consider the impact of this measurement on the determination of the dark energy equation of state. By combining the Heidelberg-Moscow result with the WMAP 3-years data and other cosmological datasets we constrain the equation of state to $-1.67< w <-1.05$ at $95 \\%$ c.l., ruling out a cosmological constant at more than $95 \\%$ c.l.. Interestingly enough, coupled neutrino-dark energy models may be consistent with such equation of state. While future data are certainly needed for a confirmation of the controversial Heildelberg-Moscow claim, our result shows that future laboratory searches for neutrino masses may play a crucial role in the determination of the dark energy properties. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608484_arXiv.txt": {
        "abstract": "In the course of our ongoing multiplicity study of exoplanet host stars we detected a faint companion located at $\\sim$~43\\,arcsec (480\\,AU physical projected separation) north-west of its primary --- the exoplanet host star HD\\,3651 at 11\\,pc. The companion, HD\\,3651\\,B, clearly shares the proper motion of the exoplanet host star in our four images, obtained with ESO/NTT and UKIRT, spanning three years in epoch difference. The magnitude of the companion is $H=16.75\\pm0.16$\\,mag, the faintest co-moving companion of an exoplanet host star imaged directly. HD\\,3651\\,B is not detected in the POSS-II B-, R- and I-band images, indicating that this object is fainter than $\\sim$~20\\,mag in the B- and R-band and fainter than $\\sim$~19\\,mag in the I-band. With the Hipparcos distance of HD\\,3651 of 11\\,pc, the absolute magnitude of HD\\,3651\\,B is about 16.5\\,mag in the H band. Our H-band photometry and the Baraffe et al. (2003) evolutionary models yield a mass of HD\\,3651\\,B to be 20 to 60\\,$M_{Jup}$ for assumed ages between 1 and 10\\,Gyr. The effective temperature ranges between 800 and 900\\,K, consistent with a spectral type of T7 to T8. We conclude that HD\\,3651\\,B is a brown-dwarf companion, the first of its kind directly imaged as a companion of an exoplanet host star, and one of the faintest T dwarfs found in the solar vicinity (within 11\\,pc). ",
        "introduction": "During the last decade high precision radial-velocity studies revealed almost 200 exoplanets around mainly solar-like stars located in the solar neighborhood. The stellar multiplicity of these planet host stars was already investigated by several groups using imaging data from visible and infrared all sky surveys like POSS or 2MASS, see e.g. Bakos et al. (2006) or Raghavan et al (2006). With seeing limited near infrared imaging (see e.g. Mugrauer et al. 2005 and 2006) as well as high contrast diffraction limited AO observations (e.g. Els et al. 2001, Patience et al. 2002, Luhman \\& Jayawardhana 2002, and Chauvin et al. 2006) further companions of exoplanet host stars have been detected during the last few years. Up to now, more than 30 companions of exoplanet host stars are known, suggesting that the multiplicity of these stars is at least 20\\,\\%.  Faint companions of exoplanet host stars were detected close to HD\\,114762 (Patience et al. 2002) and Gl\\,86 (Els et al. 2001) with angular separations smaller than 3\\,arcsec, which corresponds to physical projected separations of 130 and $\\sim$~20\\,AU, respectively. In addition to these faint close objects, Wilson et al. (2001) reported a faint, wide companion to HD\\,89744, detected in the 2MASS point source catalogue (Skrutskie et al. 2006). HD\\,89744\\,B is located 63\\,arcsec ($\\sim$ 2500\\,AU physical projected separation) north-east of HD\\,89744\\,A, and shares the proper motion of this exoplanet host star (Mugrauer et al. 2004). While the faint companions HD\\,114762\\,B ($M_{H}\\sim10.4$\\,mag) and HD\\,89744\\,B ($M_{H}\\sim11.1$\\,mag) are both very low-mass objects with masses of about 0.080\\,$M_{\\sun}$, Gl\\,86\\,B ($M_{H}\\sim14.2$\\,mag) turned out to be a white dwarf (see Mugrauer \\& Neuh\\\"auser 2005 for further details).  So far, substellar companions with masses significantly below the stellar/substellar mass border of 0.075\\,$M_{\\sun}$ have not been directly detected as companions to rad-vel planet host stars, neither with AO nor seeing limited imaging. In contrast, radial-velocity studies have already revealed a number of brown-dwarf companions close to the exoplanet host stars e.g. HD\\,38529 (Fischer et al. 2003a), HD\\,168443 (Marcy et al. 2001), and HD\\,202206 (Correia et al. 2005), with minimum masses of $msin(i)=12.7$, 17.2 and 17.4\\,$M_{Jup}$, respectively. While the brown-dwarf companions of HD\\,38529 and HD\\,168443 both revolve around the planet host star on orbits with semi-major axes which are at least 10 times wider than those of the detected exoplanets, the brown-dwarf companion of HD\\,202206 moves around the planet host star inside the planetary orbit. All of these systems are interesting examples which might point out that planet and brown-dwarf formation processes can occur around the same star at comparable distances (see Correia et. al. 2005 for further discussion).  In this letter we present the first direct imaged brown-dwarf companion of an exoplanet host star which we detected in the course of our multiplicity study. We show the results of our near infrared observations in section\\,2 and discuss the properties of this new very faint companion in section\\,3. ",
        "conclusions": ""
    },
    "0608/astro-ph0608167_arXiv.txt": {
        "abstract": "We present a measurement of the probability distribution function (PDF) of the transmitted flux in the \\op forest from a sample of 3492 quasars included in the SDSS DR3 data release.  Our intention is to investigate the sensitivity of the \\op flux PDF  as measured from low resolution and low signal-to-noise data to a number  of systematic errors such as uncertainties in the mean flux, continuum  and noise estimate. The quasar continuum  is described by the superposition of a power law and emission  lines. We perform a power law continuum fitting on a spectrum-by-spectrum basis, and obtain an average  continuum slope of $\\anu=0.59\\pm 0.36$ in the redshift range  $2.5<z<3.5$. Taking into account the variation in the continuum indices increases the mean flux by 3 and 7 per cent at $z=3$ and 2.4, respectively, as compared to the values inferred with a single (mean) continuum slope.  We compare our measurements to the PDF obtained with mock lognormal spectra,  whose statistical properties have been constrained to match the  observed \\op flux PDF and power spectrum of high resolution data. Using our power law continuum fitting and the SDSS pipeline noise estimate yields a poor agreement between the observed and mock PDFs. Allowing for a break in the continuum slope and, more importantly, for residual scatter in the continuum level substantially improves the agreement. A decrease of $\\sim$10-15 per cent in the mean quasar continuum with a typical rms variance at the 20 per cent level can account for the data, provided that the noise excess correction is no larger than $\\lsim 10$ per cent. ",
        "introduction": "\\label{sec:intro}  The \\op forest seen in quasar spectra probes the intergalactic medium (IGM) and the underlying matter distribution over a wide range of scales ($k\\sim 0.1-10\\hmmpc$) and redshifts ($1\\lsim z\\lsim 6$). Measurements of the mean flux in the \\op  forest shed light  on the reionization history and the physical state of the IGM in  the post-reionization area (Press, \\& Rybicki \\& Schneider 1993, hereafter P93; Rauch \\etal 1997; Bernardi \\etal 2003, hereafter B03; Bolton \\& Haehnelt 2006). Fluctuations in the \\op flux are of great  interest since they provide information on the matter distribution on  scales smaller than those accessible to other observables (e.g. Croft \\etal 1998; 1999; 2002b; Nusser \\& Haehnelt 1999; 2000; Pichon \\etal 2002; Zaldarriaga, Hui \\& Tegmark 2003; McDonald \\etal 2000; McDonald \\etal 2005a). Combined with CMB observations, the power spectrum of the \\op forest can provide stringent constraints  on the shape and amplitude of the primordial power spectrum (Seljak  \\etal 2004; Viel \\& Haehnelt 2005).  The probability distribution function (PDF) of the \\op transmitted flux was first studied by Jenkins \\& Ostriker (1991). It has, however, received less attention in the past years as it is more sensitive to systematics errors such as continuum fitting uncertainties. Yet, the tension between the WMAP and \\op values of the normalisation amplitude $\\sigma_8$ argues in favour of incorporating statistics others than the \\op flux power spectrum. Rauch \\etal 1997 and McDonald \\etal (2000) have computed the PDF of the \\op transmitted flux from  high resolution data and found that $\\Lambda$CDM cosmologies provide a  good fit to the observations. Gazta\\~{n}aga \\& Croft (1999) have provided an analytic description of  the \\op flux PDF based on perturbation theory. Choudhury,  Srianand \\& Padmanabhan (2001), and Desjacques \\& Nusser (2005)  have investigated the  constraints obtained from a joint  fit of the PDF  and power spectrum of the \\op transmitted flux. Becker, Rauch \\& Sargent (2006) have examined the redshift evolution of the flux \\op PDF from a large sample of Keck HIRES data. Note also that Lidz \\etal (2005) have advocated working with the PDF of the fluctuations in the flux about the mean as it is insensitive to the quasar continuum.  The Sloan Digital Sky Survey (SDSS; York \\etal 2000) has greatly increased the statistical power of the \\op forest. Unfortunately, attempts to exploit the data are plagued by a number of poorly constrained parameters that describe the physical state of the IGM, by inaccuracies in the numerical modelling of the \\op forest, and by systematics errors in the measurement such as continuum fitting uncertainties.  The observed flux in the \\op region of a quasar (QSO) spectrum depends both on the quasar continuum (the flux emitted by the quasar, including the emission lines) and on the amount of absorption by intervening  galactic matter (lines and continuum absorptions). The transmitted (or normalized) flux  obtained by continuum fitting of the observed spectrum is related to  the optical depth along the line of sight as \\begin{equation} F(w)\\equiv I_{\\rm obs}(w)/I_{\\rm cont}={\\rm e}^{-\\tau(w)}\\; , \\label{eq:cont} \\end{equation} where $\\tau$ is the optical depth, $w$ is the redshift space coordinate along the line of sight, $I_{\\rm obs}$ is the observed flux and $I_{\\rm cont}$ is the flux emitted from the source (quasar) that would be observed in the absence of any intervening material.  To estimate the unabsorbed continuum, two different approaches have been used. When the signal-to-noise is large, a polynomial continuum  is fitted to regions of the \\op forest free of absorption (e.g. Rauch \\etal 1997; McDonald \\etal 2000). When the signal-to-noise is low, one usually extrapolate the continuum redward of the \\op emission  line, assuming a power law shape. The last method only provides an approximation to the true continuum, and can easily introduce several types of systematic errors (e.g. Steidel \\& Sargent 1987).  The strong degeneracy between the effective optical depth  $\\tef=-\\ln\\la F\\ra$ and the continuum level affects the determination of the clustering amplitude $\\sigma_8$.  Measurements of the effective optical depth from low resolution quasar spectra with comparable signal-to-noise (P93; B03) are systematically higher than those based on high-resolution  spectra by 10-20 per cent. As argued by Seljak  \\etal (2003), Tytler  \\etal (2004), Viel \\etal (2004b), this is most  likely due to systematic  errors in the continuum fitting procedure of the low resolution spectra.  Furthermore, the low signal-to-noise of SDSS spectra in the \\op region ($S/N\\sim 3$ typically) requires an accurate noise estimate. Systematic errors in the noise characterisation will affect the accuracy of the measurements. In this respect, several lines of recent evidence suggest that the SDSS reduction pipeline underestimates the true noise by 5-10 per cent (McDonald \\etal 2006; Burgess 2004).  In this paper, we measure the PDF  of the \\op transmitted flux from a large (public) sample of quasars included in the SDSS DR3 data release (Abazajian \\etal 2005). We assume that the quasar continuum follows the parametric form given in B03. However, unlike B03, we also allow for the variation in the continuum indices of individual spectra and fit a power law continuum for each spectrum separately. We compare our measurements to the probability distribution obtained from lognormal, realistic looking SDSS spectra. The  statistical properties of these mock spectra are beforehand constrained to match the observed \\op flux probability distribution and power  spectrum of high resolution data of the forest. Our intention is to investigate the sensitivity of the \\op flux PDF as measured from low resolution, low signal-to-noise data to a number of systematic errors such as uncertainties in the mean flux, continuum and noise estimate.  The paper is organised as follows. We briefly review the lognormal model of the \\op forest in Section~\\S\\ref{sec:lya}. The constraints on the model parameters are discussed in~\\S\\ref{sec:param}. The continuum fitting procedure and the measurement of the PDF of the  \\op flux are presented in~\\S\\ref{sec:data}. In~\\S\\ref{sec:constrain}, we compare the simulated and observed PDF, and study the effect of a number of systematic errors. We discuss our results  in~\\S\\ref{sec:discussion}. In ~\\S\\ref{sec:conclusion}, we conclude and indicate potential future works. We will present results  for a $\\Lambda$CDM cosmology with normalisation amplitude $\\sigma_8=0.83$, and spectral index $n_s=0.96$. This is consistent with the constraints obtained from the latest CMB and \\op forest data (Spergel \\etal 2006; Viel, Haehnelt \\&  Lewis 2006; Seljak, Slosar \\& McDonald 2006). ",
        "conclusions": "\\label{sec:conclusion}  We have presented measurements of the probability distribution of the \\op transmitted flux in the redshift range $2.5\\lsim z\\lsim 4$, from 3492 quasars included in the SDSS DR3 data release. We have compared the measured PDF to predictions derived from mock spectra, whose statistical properties have been constrained to match those of high resolution data. To proceed, we have generated very long, lognormal spectra of the \\op forest that have been degraded to include the instrumental noise and resolution of real data. The mock spectra provide a good match to the \\op flux PS and PDF measured in McDonald \\etal (2000) in the region $z\\bsim 3$.  We have assumed that the quasar continuum follows the parametric form given in B03. However, unlike B03, we have allowed for the slope of the power-law continuum to vary from object to object. We measure an average continuum slope of $\\anu-0.59\\pm 0.36$ in the range $2.4\\lsim z\\lsim 3.6$, in good agreement with the mean slope reported by Vanden Berk \\etal (2001), $\\anu=-0.44$, and with values found in optically selected samples (e.g. Francis \\etal 1991; Natali \\etal 1998). Accounting for variation in continuum indices has a significant impact on the mean flux $\\bar{F}$. We find that $\\bar{F}$ at redshift $z=3$ and 2.4 is respectively 3 and 7 per cent higher than the mean flux measured for a fixed index $\\anu=-0.44$.  Although the model parameters have been adjusted to reproduce the observed \\op flux PS and PDF of high resolution data, the mock  SDSS spectra predict a probability distribution that is significantly different from the PDF we measure from the SDSS quasar sample. Allowing for a break in the continuum and, more importantly, for residual scatter in the continuum level improve the agreement substantially. We find that the introduction of a  10 per cent scatter in the continuum level of $z\\lsim 3$ mock  spectra  smoothes the flux PDF on a level comparable to a 20-30 per cent  increase in the noise. A combined fit of the SDSS and Keck data indicates that a  decrease of 10-15 per cent in the amplitude of the power law  continuum together with a 20 per cent scatter can account for the data, provided that the noise excess correction is no larger than $\\lsim$10 per cent.  Measuring the probability distribution of the transmitted flux requires  a spectrum-by-spectrum treatment of the quasar continuum, as the latter varies significantly from quasar to quasar. Furthermore, as we have seen,  it is crucial to account for the slow variation of the continuum slope in the \\op region in order to obtain a sensible estimate of the flux probability distribution.  Therefore, it would be desirable to obtain high resolution exposures of a subsample of SDSS quasars so as to quantify the errors introduced by the continuum fitting procedure described in this paper. Alternatively,  Lidz \\etal (2005) have suggested working with the estimator defined by $\\delta_f(\\omega)=\\left(I_{\\rm obs}^r(\\omega)- I_{\\rm obs}^R(\\omega)\\right)/I_{\\rm obs}^R(\\omega)$, where $I_{\\rm obs}^r$ and $I_{\\rm obs}^R$ is the observed flux smoothed with a Gaussian filter on scale $r\\ll R$ and $R\\sim 500\\kms$, respectively. They have demonstrated that the PDF of $\\delta_f$ is insensitive to the shape and  normalisation of the continuum. However, this estimator has an important drawback as it also smoothes out any feature in the redshift evolution of the mean optical  depth, $\\tef(z)$, whose scale is larger than $R$. Their choice of $R$ corresponds to a redshift interval $\\Delta z\\lsim 0.01$ at $z=3$. It is therefore unclear whether the sudden change measured by B03 at $z\\sim 3.2$ can  be detected in the statistics of flux estimators other than the transmitted flux $F$."
    },
    "0608/hep-ph0608334_arXiv.txt": {
        "abstract": " ",
        "introduction": "Various experiments for solar and atmospheric neutrinos provide a range for the values of solar mixing angle $\\theta_\\odot = \\theta_{12}$ (the 1-2 mixing angle) \\cite{sfit} that corresponds to solar neutrino oscillations and also a range for atmospheric mixing angle $\\theta_{\\rm atm} = \\theta_{23}$ (the 2-3 mixing angle) \\cite{sfit} around their best fit values. The tri-bimaximal mixing condition of neutrinos are given by $\\sin \\theta_{12} = \\frac {1} {\\sqrt {3}}$, $\\sin \\theta_{23} = \\frac {1} {\\sqrt {2}}$ and $\\sin \\theta_{13} = 0$ \\cite{tbi}. Possible deviations from tri-bimaximal mixing can be obtained by probing the ranges of $\\theta_{12}$ and $\\theta_{23}$ given by the experiments. Also the exact 13 mixing angle $\\theta_{13}$ is not known except that the CHOOZ \\cite{chooz} gives an upper limit for $\\theta_{13} (< 9^o)$. Probing the deviations of $\\theta_{12}$ and $\\theta_{23}$ for different values for $\\theta_{13}$ is significant not only to understand the neutrino flavour oscillations in general but also for the purpose of model building for neutrino mass matrices.  In this work we explore the possibility for ultra high energy (UHE) neutrinos from distant Gamma Ray Bursts (GRBs) for probing the signatures of these deviations of the values of the mixing angles from tri-bimaximal mixing as discussed above. One such proposition of using UHE neutrinos is described in a recent work by Xing \\cite{xing}. Gamma Ray Bursts are short lived but intense burst of gamma rays. During its occurence it outshines all other luminous objects in the sky. Although the exact machanism of GRBs could not be ascertained so far but the general wisdom is that it is powered by a central engine provided by a failed star or supernova that possibly turned into a black hole, accretes mass at its surroundings. This infalling mass due to gravity bounces back from the surface of black hole much the same way as the supernova explosion mechanism and a shock is generated that flows radially outwards with enormous amount of energies ($\\sim 10^{53}$ ergs). This highly energetic shock wave drives the mass outwards, in the form of a ``fireball\" that carries in it, protons, $\\gamma$ etc. The pions are produced when the accelerated protons inside the fireball interacts with $\\gamma$ through a cosmic beam dump process. UHE neutrinos are produced by the decay of these pions. Thus a generic cosmic accelerator accelerates the protons into very high energies which then beam dump on $\\gamma$ in the ``fireball\" as also at the cosmic microwave background (CMB) and ultra high energy neutrinos are produced.  The GRB neutrinos, due to their origin at astronomical distances from earth, provide a very long baseline for the earth bound detectors for UHE neutrinos such as ICECUBE \\cite{icecube}. The oscillatory part of the neutrino flavour oscillation probabilities ($\\sin^2(\\Delta m^2 [L/4E])$) averages out to 1/2 becuase of this very long baseline $L$ ($\\sim$ hundreds of Mpc) and the $\\Delta m^2$ (mass square difference of two neutrinos) range obtained from solar and atmospheric neutrino experiments are $\\Delta m^2_{21} \\sim 10^{-4}$ eV$^2$ and $\\Delta m^2_{32} \\sim 10^{-3}$ eV$^2$ respectively ($L/\\Delta m^2 >> 1$). Thus for neutrino flavour oscillation, in this case, the effect of $\\Delta m^2$ is washed out and governed only by the three mixing angles namely $\\theta_{12} = \\theta_\\odot$, $\\theta_{23} = \\theta_{\\rm atm}$ and $\\theta_{13}$. The purpose of the present work is to probe whether or not the possible variations of $\\theta_{12}$ and $\\theta_{23}$  from their best fit values can be ascertained by UHE from distant GRBs.  The GRB neutrinos, on arriving the earth, undergo charged current (CC) and neutral current (NC) interactions with the earth rock and the detector material. The CC interactions of $\\nu_\\mu$ produce secondary muons and the same for electrons produce electromagnetic shower ($\\nu_\\mu + N \\rightarrow \\mu + X$ and $\\nu_e + N \\rightarrow e + X$). The former will produce secondary muon tracks and can be detected by track-signal produced by the Cerenkov light emitted by these muons during their passage through a large underground water/ice Cerenkov detectors like ICECUBE. The ICECUBE is a 1km$^3$ detector in south pole ice and can be considered to be immersed in the target material for the UHE neutrinos where the neutrino interactions are initiated. In case of $\\nu_e$, the electrons from the $\\nu_e N$ CC interactions, shower quickly and can also be detected by such ICECUBE detector. The case of $\\nu_\\tau$ is somewhat complicated. The first CC interaction of $\\nu_\\tau$ ($\\nu_\\tau + N \\rightarrow \\tau + X$) produces a shower (``first bang\") alongwith a $\\tau$ track. But the $\\nu_\\tau$ is regenerated (with diminished energy) by the decay of $\\tau$ and in the process produces another hadronic or electromagnetic shower (``second bang\"). The whole process is called double bang event. In case the first bang could not be detected, then by possible detection of second bang (with showers) the $\\tau$ track can be reconstructed or identified and this scenario (the $\\tau$ track and the second bang) is called the lollipop events. An inverted lollipop event is one where only the first bang ($\\nu_\\tau + N \\rightarrow \\tau + X$)  is detected and the subsequent $\\tau$ track is detected or reconstructed. As mentioned in Ref. \\cite {cao}, the detection of $\\nu_\\tau$ from their CC interaction mentioned above is not every efficient by a 1km$^3$ detector since the double bang events can possibly be detected only for the $\\nu_\\tau$ energies between 1 PeV to 20 PeV beyond which the tau decay length is longer than the width of such detector and at still higher energies the flux is too small for such detectors for their detection. Hence, in the present work we do not consider the events initiated by $\\nu_\\tau N$ CC interactions. However, for $\\nu_\\tau$ we consider the process that may yield events higher than the ``double bang\" events. We consider the decay channel of $\\tau$ lepton \\cite{sarada}, obtained from charged current interactions of $\\nu_\\tau$, where muons are produced ($\\nu_\\tau \\rightarrow \\tau \\rightarrow {\\bar \\nu_\\mu} \\mu \\nu_\\tau$) which can then be detected as muon tracks \\cite{nayantau} in ICECUBE detector. The neutral current (NC) interactions of all flavours however will produce the shower events at ICECUBE and they are considered in this investigation.  This paper is organised as follows. In Section 2 we describe the formalism for neutrino fluxes of the three species while reaching the earth. The nature of the GRB flux taken for present calculations is also discussed. The flux suffers flavour oscillations while traversing from GRB site to the earth. The oscillation probabilities are also calculated and the oscillated flux obtained on reaching the earth is determined. They are given in Section 2.1. We also describe in this section the analytical expressions for the yield of secondary muons and shower events at the ice Cerenkov kilometre square detector like ICECUBE. This is given in Section 2.2. The actual calculations and results are discussed in Section 3. Finally, in Section 4, some discussions and summary are given. ",
        "conclusions": "In summary, we investigate the deviation from the well known tri-bimaximal mixing in the case of Ultra High Energy neutrinos from a Gamma Ray Burst detected in a kilometer scale detector such as ICECUBE. We have calculated the ratio ${\\cal R}$ of the muon track events and shower events (electromagnetic shower from charged current interactions of $\\nu_e$ and hadronic showers from neutral current interactions of neutrinos of all flavours) for tri-bimaximal mixing condition given by $\\theta_{12} = 35.2^o$, $\\theta_{23} = 45.0^o$, $\\theta_{13} = 0^o$. We then investigate the possible variation of  ${\\cal R}$ from tri-bimaximal mixing condition by varying $\\theta_{12}$ and $\\theta_{23}$ within their experimentally obtained range for four different values of $\\theta_{13}$ namely $0^o,\\,\\,3^o\\,,\\,6^o$ and $9^o$.  The isotropic flux of GRB neutrinos are obtained following Waxman-Bahcall \\cite{waxman} type parametrization of the flux and summation over the sources. The initial parametrization of neutrino flux can be written as \\begin{equation} \\displaystyle\\frac {dN_\\nu} {dE_\\nu} = \\left \\{ \\begin{array}{c} \\frac {A} {E_\\nu E_\\nu^b}\\,\\, , \\,\\,\\, E_\\nu < E_\\nu^b \\\\ \\frac {A} {E_\\nu^2}\\,\\, , E_\\nu > E_\\nu^b \\end{array} \\right . \\end{equation} where $E_\\nu^b$ is the spectral break energy ($\\sim 10^5$ GeV) and is related to photon spectral break energy, Lorentz factor etc.  The GRB neutrinos after reaching the earth has to pass through the earth rock (for upward going events) to reach the detector to produce muon tracks or shower. In the calculation therefore, the attenuation of neutrinos through the earth (shadow factor) is estimated. The muons produced out of charged current interactions of neutrinos should also survive to enter the detector and produce tracks. Therefore, to estimate the muon track events, the energy loss of muons through the rock is also estimated. The average lepton energy loss rate (with lepton energy $E_\\mu$) due to ionisation and the losses due to Bremsstrahlung, pair-production, hadron production etc. (catastrophic losses) is parametrized as $$ \\displaystyle \\left \\langle \\frac {dE_\\mu} {dX} \\right \\rangle = -\\alpha - \\beta E_\\mu $$ where $\\beta$ describes the catastrophic loss which dominate over the ionisation loss above a certain critical energy $\\zeta = \\alpha/\\beta$. This induces a logarithmic dependence of the lepton energy loss.  The calculated ratio ${\\cal R}$ varies between $\\sim 8\\%$ to $\\sim 18\\%$ for variation from tri-bimaximal mixing scenario and for different values of $\\theta_{13}$. Given the sensitivity of the ICECUBE detector in terms of detecting GRB neutrino flux and considering other uncertainties like that in estimating the flux itself, the atmospheric background, low signal yield and the systematic uncertainties of the detector, it appears that ICECUBE with its present sensitivity will not be able to detect significantly such a small variation due to deviations from tri-bimaximal mixing. Hence to detect such small deviation, very precise mesurement is called for. This requires more data (more years of run) and larger detector size for more statistics. The increase in detector size will not widen the deviation of the ratio significantly as the total area factor of the detector cancels out in the ratio (Eq. 23) although the total number of both muon tracks and total shower yield increase significantly. For the case of shower, the whole detector volume is to be considered and from Eq. (22), there is indeed an $L$ dependence. This makes the deviation of the ratio wider although very marginally as we increase the detector dimension.  It is difficult to predict the detector dimension and/or the time of exposure that will be suitable for such a precision measurement discussed above. Detailed simulation studies taking into account factors like atmospheric neutrino background, photomultiplier tube efficiency and other possible uncertainties like the one carried out in Ref. [16] is required for being able to comment on the detector parameters for such precise measurements.  We also want to mention in passing that we have repeated the same calculation for single GRBs with fixed red shift ($z$) values with similar results."
    },
    "0608/nucl-th0608045_arXiv.txt": {
        "abstract": "We discuss effects of the electron plasma on charged-current neutrino-nucleus reaction, $(\\nu_e,e^-)$ in a core-collapse supernova environment. We first discuss the electron screening effect on the final state interaction between the outgoing electron and the daughter nucleus. To this end, we solve the Dirac equation for the outgoing electron with the screened Coulomb potential obtained with the Thomas-Fermi approximation. In addition to the screening effect, we also discuss the Pauli blocking effect due to the environmental electrons on the spectrum of the outgoing electron. We find that both effects hinder the cross section of the charged-current reaction, especially at low incident energies. ",
        "introduction": "A large number of neutrinos are emitted from a core-collapse supernova. These neutrinos interact with nuclei through the weak interaction. Although their cross sections are small, it is agreed that their contribution to nucleosynthesis (that is, r-, $\\nu$- and p-processes) is not negligible due to the large neutrino luminosity \\cite{Balantekin:2003ip}. Neutrinos may even play a leading role in some cases. For instance, Yoshida {\\it et al.} recently argued that the abundance ratio between $^7$Li and $^{11}$B is sensitive to the $\\nu$-process and thus can be used to extract information on the neutrino mass hierarchy \\cite{Yoshida2006}. Also, the abundance ratio between U and Th elements, which has been used as a cosmochronometer, may be affected by the $\\nu$-process. It is thus important to calculate with high accuracy the cross section of the neutrino-nucleus reactions in a dense star.  In the supernova nucleosynthesis, only the charged current reactions of the electron neutrinos, $\\nu_{e}$, and the electron anti-neutrinos, $\\overline{\\nu}_e$, are relevant, since those of $\\nu_{\\mu}$ and $\\nu_{\\tau}$ (and their antineutrinos) are suppressed due to the threshold effects. These processes are represented by \\begin{eqnarray} \\nu_{e} + (Z,A) &\\rightarrow& e^- + (Z+1,A),\\label{eq:1} \\\\ \\overline{\\nu}_{e} + (Z,A) &\\rightarrow& e^+ + (Z-1,A).\\label{eq:1-2} \\end{eqnarray} The average energies of $\\nu_{e}$ and $\\overline{\\nu}_{e}$ neutrinos in a core collapse supernova are typically 11 MeV and 16 MeV, respectively \\cite{QFMMWW93}. It must be emphasized, however, that the exact hierarchy of neutrino energies depend on the details of the underlying microphysics \\cite{micro}. The outgoing electron and positron feel the Coulomb interaction from the daughter nucleus as they leave. This final state interaction affects the neutrino-nucleus reaction rate \\cite{Morita,Walecka}.  In the supernova environment, the motion of the outgoing electron is further perturbed by other charged particles in the environment besides the daughter nucleus, such as protons, electrons, $\\alpha$ particles and other heavier nuclei. Such effects have been considered in Ref. \\cite{Itoh2002} for electron capture rates in a dense star. Furthermore, in a high electron density, the reaction (\\ref{eq:1}) is suppressed because low energy electron states are Pauli blocked. It is crucial to take into account those two effects in order to accurately estimate the neutrino-nucleus reaction rate for nucleosynthesis.  In this paper, we perform such calculations, taking into consideration both the electron screening and Pauli blocking effects. For this purpose, we use the Thomas-Fermi theory for electron screening, which has often been used in many-body physics. We then calculate the cross section of a neutrino-nucleus reaction using the distorted wave Born approximation (DWBA) method, including the screening and the Pauli blocking effects simultaneously.  The paper is organized as follows. In Sec. II, we detail our model for the charged current neutrino-nucleus reaction in a supernova environment, that uses the Thomas-Fermi theory for electron screening, and the DWBA method with the Pauli blocking effect. In Sec. III, we present the results of our calculation for the $^{56}$Fe($\\nu_{e},e^-)^{56}$Co and $^{208}$Pb($\\nu_{e},e^-)^{208}$Bi reactions, and discuss the effects of the electronic environment. We summarize our paper in Sec. IV. ",
        "conclusions": "The motion of the outgoing electron in the final state of a charged current neutrino-nucleus reaction is affected by surrounding environmental electrons in a supernova. We have discussed the electron screening as well as the Pauli blocking effects due to the environmental electrons on cross sections of the neutrino-nucleus reaction. For this purpose, we used the Thomas-Fermi theory for the screening potential, and the DWBA with the Pauli blocking factor for the cross sections. Our results for the $^{208}$Pb$(\\nu_{e},e^{-})^{208}$Bi and $^{56}$Fe$(\\nu_{e},e^{-})^{56}$Co reactions show that both the effects hinder the cross sections, especially at high electron densities. We have also shown that the Pauli blocking effect is more significant in the former reaction than in the latter reaction due to the larger excitation energy for the Fermi transition. The screening effect is also larger in the former reaction because of the larger atomic number.  As far as we know, the screening and the Pauli blocking effects discussed in this paper have not been taken into account in the calculations for r-process nucleosynthesis. It would be of intriguing to see how these effects influence the r-process nucleosynthesis."
    },
    "0608/astro-ph0608571_arXiv.txt": {
        "abstract": "Stellar oscillations can provide a wealth of information about a star, which can be extracted from observed time series of the star's brightness or radial velocity. In this paper we address the question of how to extract as much information as possible from such a dataset. We have developed a Markov Chain Monte Carlo (MCMC) code that is able to infer the number of oscillation frequencies present in the signal and their values (with corresponding uncertainties), without having to fit the amplitudes and phases. Gaps in the data do not have any serious consequences for this method; in cases where severe aliasing exists, any ambiguity in the frequency determinations will be reflected in the results. It also allows us to infer parameters of the frequency pattern, such as the large separation $\\Delta \\nu$. We have previously applied this method to the star $\\nu$ Indi \\citep{nuindi}, and here we describe the method fully and apply it to simulated datasets, showing that the code is able to give correct results even when some of the model assumptions are violated. In particular, the non-sinusoidal nature of the individual oscillation modes due to stochastic excitation and damping has no major impact on the usefulness of our approach. ",
        "introduction": "Asteroseismology is fast becoming a standard tool in stellar astrophysics \\citep[e.g.][]{2005JApA...26..123K}. The properties of a star's oscillations are determined by the star's structure and composition, so measurements of those oscillations have enabled astronomers to probe the interiors of stars. In particular, the frequencies of the various modes can provide information about the speed of sound at various depths in the star, and hence constrain its composition. Recently, there has been a rapid increase in the amount of data available for studying solar-like stars in this way \\citep[for reviews, see][]{2003PASA...20..203B, 2003Ap&SS.284...21B, 2004soho...14..101K}.  The data analysis procedures used for asteroseismology have tended to be based on Fourier methods. For example, the first step taken by most authors analysing the raw time series data is usually the computation of the power spectrum, or periodogram. However, it has been shown \\citep{bretthorst1} that the periodogram is only an optimal procedure if we are estimating the frequency of the oscillations with a model that assumes that only one frequency is present, or, if there are many frequencies, that they are well separated. In the case of solar-like oscillations, we expect the star to be oscillating in many modes simultaneously, some of which may be closely spaced. Furthermore, the modes are damped and stochastically excited, so each mode cannot be represented by a single sinusoid. There is also a guiding theory telling us that the frequencies of the modes are not arbitrary, but form a well defined pattern. This constitutes additional information that we might wish to take into account when analysing observations.  The desire to extract as much information as possible from an astronomical dataset, and to combine this with prior information from other datasets or theoretical considerations, has led to an explosion in the use of Bayesian methods \\citep[e.g.][]{2005ApJ...631.1198G, 0047, cosmonest}. The basic format of Bayesian inference is as follows. Our state of knowledge about any set of parameters $\\boldtheta$ is described by a probability distribution over the parameter space. With prior information and assumptions $I$, a prior distribution $p(\\boldtheta|I)$ (read as the distribution for $\\boldtheta$ {\\it given} $I$) is assigned to describe what we know about the parameters before taking into account the current dataset. Then, the effect of taking into account data $\\bold{D}$ is to modify our state of knowledge from the prior distribution $p(\\boldtheta|I)$ to the posterior distribution $p(\\boldtheta|\\bold{D},I)$, given by  \\begin{equation}\\label{bayes} p(\\boldtheta|\\bold{D},I) \\propto p(\\boldtheta|I)p(\\bold{D}|\\boldtheta,I). \\end{equation}  The distribution $p(\\bold{D}|\\boldtheta,I)$ is the probability distribution for the data given the parameters. When the data are fixed, this is a function of the parameters and is called the likelihood function. Hence, whether a model (a point in the parameter space) becomes plausible after taking into account data $\\bold{D}$ depends on how plausible the model was before taking into account the data and on how well the model predicts the observed data. This kind of reasoning has led to the development of new methods of data analysis, as well as new interpretations and understanding of familiar methods\\footnote{For example, a least squares fit of a model to some data can be interpreted as finding the most probable set of parameters, with a Gaussian distribution for the data points and no strong prior information.}. An example relevant to spectral analysis is the finding of \\citet{bretthorst1} that the periodogram is proportional to the log of the posterior probability density for the frequency of an oscillating signal, under certain assumptions. For more details on Bayesian data analysis, including its relationship to common Fourier techniques, see the textbook by \\citet{2005blda.book.....G}.  Typically, the datasets that are used for asteroseismology are time series of either the brightness (photometry) or the radial velocity of the stellar surface (from spectroscopy), measured at a discrete set of times $\\{t_1,t_2,...,t_N\\}$, and with each data point subject to an observational error $\\{e_1,e_2,...,e_N\\}$. This results in the observed values for the radial velocity (or magnitude) $y$ at each of the $N$ times: \\begin{equation}\\label{noise} y_i = f(t_i) + e_i \\end{equation} where $f(t)$ is the true form of the signal, and the $\\{e_i\\}$ are the observational errors. If the star is oscillating in $M$ different (pure) modes, with frequencies $\\{\\nu_1, \\nu_2, ..., \\nu_M\\}$, then the functional form of the signal is \\begin{equation}\\label{signal} f(t) = \\sum_{j=1}^M\\left (A_j cos(2\\pi\\nu_j t) + B_j sin(2\\pi\\nu_j t)\\right) \\end{equation} The goal of the data analysis is usually to use the observed data $\\bold{y} = \\{y_1, y_2, ..., y_N\\}$ to infer the frequencies $\\boldnu = \\{\\nu_1, \\nu_2, ..., \\nu_M\\}$, the number of frequencies $M$ and possibly their amplitudes. However, the values of the frequencies are not arbitrary. For low degree p-mode oscillations, the frequencies are given approximately by \\begin{equation}\\label{pattern} \\nu_{n,l} = \\Delta \\nu (n + \\frac{1}{2}l + \\epsilon) - l(l+1)D_0 \\end{equation} The values of the parameters of this pattern, $\\Delta \\nu$, $\\epsilon$ and $D_0$, are also of interest as they are related to the sound speed at various depths in the star \\citep{1994ApJ...427.1013B}. Throughout this paper, we will denote any parameters of the frequency pattern collectively by $\\boldtheta$, and concentrate on finding these parameters and the frequencies. The amplitudes $\\{A_i\\}$, $\\{B_i\\}$ are regarded as nuisance parameters. In a recent paper \\citep{nuindi}, we used a Bayesian method to estimate the large separation $\\Delta \\nu$ of the metal-poor subgiant $\\nu$ Indi. The purpose of this paper is to describe the method in greater detail and demonstrate its effectiveness on simulated data. ",
        "conclusions": "In this paper, we presented a data analysis method for solar-like oscillation data, based on Bayesian inference and Markov Chain Monte Carlo methods. These methods are becoming very popular in astronomy and science in general, as they are often able to extract more information from the dataset than more ad-hoc approaches. In addition, reliable uncertainties in all inferred quantities are easily obtained.  The method we have presented here, which we have already applied to $\\nu$ Ind \\citep{nuindi}, has several advantages over more conventional approaches. For example, the CLEAN algorithm \\citep{CLEAN}, which is based on iterative sine wave fitting, requires that the amplitude and phase be fitted as well. Then, when this fitted curve is subtracted, biases are introduced. This is because the fitted parameters are not exactly correct. Since our code fits all of the frequencies simultaneously, it does not suffer from this problem. Also, our method is able to take into account important prior information that we have about the expected pattern of frequencies, which other methods ignore. It could be argued that we aren't really sure that the prior information is correct, and therefore taking it into account may be giving us overconfident results. However, we have placed safeguards in our method to ensure that this does not happen - this was the purpose of introducing the $\\sigma$ parameter and allowing it to be determined by the data. Importantly, we have also found that when some of the model assumptions are violated (in particular, if the modes are damped and stochastically excited), the method continues to give useful results.  Of course, this method is more computationally intensive than the usual power spectrum based methods\\footnote{For our simulated data, usable output was obtained in a timescale of several hours on a modern PC.}. We would not recommend use of our method on very large datasets, for example those from the sun. In these cases, the power spectrum can be computed swiftly and is very informative. We believe that our method is best suited to those datasets where the data values (and hence the power spectrum) are noisy and incomplete, and aliasing causes difficulties in the interpretation of the power spectrum. As observational techniques improve, the number of target stars will increase and there will always be datasets for which this is the case. We hope that the approach we have presented here will prove useful in this field."
    },
    "0608/astro-ph0608601_arXiv.txt": {
        "abstract": "We have discovered a spectral transition from the low/hard (LH) state to the high/soft (HS) state when Aquila X-1 was {\\em declining} in observations made with the {\\it Rossi X-Ray Timing Explorer (RXTE)}. The 2--200 keV energy flux corresponding to the state transition is $1.1\\times{10}^{-9}~ {\\rm ergs~cm^{-2}~s^{-1}}$, an order of magnitude lower than observed in the past. The 2--200 keV peak flux of the following HS state is $1.6\\times{10}^{-9} ~{\\rm ergs~cm^{-2}~s^{-1}}$. The relation between the luminosity of the hard-to-soft state transition and the peak luminosity of the following HS state confirms the linear relation found previously. This implies that the luminosity of the hard-to-soft state transition is not determined solely by the mass accretion rate, but appears to be determined by the peak luminosity of the soft X-ray outburst. We also found that the LH-to-HS state transition occurred at a luminosity similar to that of the corresponding HS-to-LH state transition, i.e., there is no apparent hysteresis. These results provide additional evidence that the mass in the accretion disk affects the luminosity of the hard-to-soft state transition, and that the accretion flow that powers the LH state is related to the accretion flow that powers the HS state at a later time. ",
        "introduction": "The low/hard (LH) state and the high/soft (HS) state are the two main X-ray spectral states identified in the Galactic black hole binaries (see the review by McClintock \\& Remillard 2006 and references therein). Similar spectral states, namely the island state and the banana state (Hasinger \\& van der Klis 1989), have been seen in the atoll sources in the neutron star low-mass X-ray binaries (LMXBs) (e.g. van der Klis 1994; Barret \\& Vedrenne 1994). The two spectral states, as well as transitions between the two states, are also similar in terms of their X-ray energy spectra and variability properties (e.g., Yu et al. 2003). These observational properties suggest that the same physics is involved in both black holes and neutron stars during the state transitions.  Mass accretion rate has long been thought to determine the spectral state of an accreting black hole or neutron star; thus, the variation of mass accretion rate was thought to cause the transitions  between the states (e.g., Esin et al. 1997). However, recent observations show that mass accretion rate is not the only parameter that determines the spectral states (Homan et al. 2001; Smith et al. 2002; Yu et al. 2004,hereafter YKF04). Suggestions for possible other parameters include (1) coronal size (Homan et al. 2001), (2) the size of the accretion disk (Smith et al. 2002), (3) the past history of the location of the inner edge of the accretion disk (Zdziarski et al. 2004), and (4) the mass in the accretion disk (YKF04).  The suggestion that the mass in the accretion disk affects state transitions is based on the correlation between the luminosity of the hard-to-soft state transition and the peak luminosity of the following HS state in the observations of two outbursts in each of the transients \\mbox{Aql X-1}, \\mbox{XTE J1550$-$564}, and a flaring neutron star LMXB \\mbox{4U 1705$-$44} (YKF04). The correlation indicates that the higher the luminosity of the hard-to-soft state transition, the brighter the following HS state will be. Because a brighter outburst or flare normally has a duration longer than or similar to a dimmer outburst or flare, a brighter outburst is associated with the accretion of more mass onto the compact object. There is additional evidence that the hard, power-law spectral component is associated with the mass in the disk. In the black hole transient \\mbox{GX 339$-$4}, the peak fluxes of the LH states in the seven outbursts seen by BATSE and HEXTE in the past 15 years are nearly linearly related to the outburst waiting time (Yu et al. 2007). The correlation also suggests that the earlier accretion flow that powers the LH state and the later accretion flow that powers the HS state are somehow related, suggesting that the outer disk plays a role in the generation of the hard spectral component that dominates the X-ray emission in the LH state (YKF04; Yu et al. 2007). A similar correlation in the flaring neutron star LMXB \\mbox{4U 1705$-$44} (YKF04; W. Yu 2007, in preparation) indicates that the phenomena are associated with the variation of the mass accretion rate in the accretion flow, rather than the outburst mechanism of soft X-ray transients (e.g., Cannizzo et al. 1995).  In this paper, we report our study of four outbursts of \\mbox{Aql X-1} in which the flux peaks of the LH states during the outburst rise and the following HS states were covered by {\\it RXTE} observations. It is especially noteworthy that we have identified a hard-to-soft state transition during a {\\it luminosity decline} of the LH state in a very faint outburst in 2001, in which the following soft-to-hard state transition occurred at a similar luminosity level. This outburst is actually fainter than the 2005 outburst, which was a hard state outburst candidate, observed by {\\it RXTE} and {\\it INTEGRAL} (Rodriguez et al. 2006). We also show that these observations establish a nearly linear relation between the luminosity of the hard-to-soft state transition and the peak luminosity of the following HS state over a luminosity range of an order of magnitude. The results suggest that the hysteresis of state transitions (see Miyamoto et al. 1995) is not always present in the same LMXB. ",
        "conclusions": "We have studied four outbursts of \\mbox{Aql X-1} in which the hard-to-soft state transitions were covered by {\\it RXTE} observations. In one of the four outbursts, we have discovered a hard-to-soft state transition which occurred during a luminosity decline and at a luminosity similar to that of the following soft-to-hard state transition. We have also found that the peak luminosities of the HS states follow the X-ray luminosities of the hard-to-soft state transitions nearly linearly. The study confirms the luminosity correlation previously found by Yu et al. (YKF04) and extends it to much lower luminosities. Using the best-fit linear relation, we can predict the peak luminosity of an outburst in \\mbox{Aql X-1} at the very beginning of the hard-to-soft state transition.  The observation of the hard-to-soft state transition during a luminosity decline demonstrates that the hard-to-soft state transition is not associated with the flux peak of the LH state before the state transition. In all other outbursts of \\mbox{Aql X-1}, as well as in the outbursts of several other transient sources such as \\mbox{GX 339$-$4} (Yu et al. 2007), the state transitions were associated with the luminosity peaks of the LH states (Yu et al. 2003, 2007; YKF04). Thus, it is the luminosity of the hard-to-soft state transition, instead of the peak luminosity of the LH state before the state transition, that is correlated with the peak luminosity of the following HS state.  In the most popular picture of spectral states in X-ray binaries, e.g., the model proposed by Esin et al. (1997), spectral states are determined by the mass accretion rate. A source in the LH state can only transit to the HS state when the mass accretion rate increases. Therefore, the hard-to-soft state transition only occurs during a luminosity increase. The observation of a hard-to-soft state transition during a luminosity decrease in \\mbox{Aql X-1} contradicts this idea, implying that other parameters in addition to the mass accretion rate also determine the spectral transition. Since the luminosity of the hard-to-soft state transition can vary by an order of magnitude and that of the soft-to-hard state transition is usually at a much lower luminosity and remains nearly constant (Maccarone 2003), the hysteresis of the state transitions (e.g., Miyamoto et al. 1995) likely originates from the large variation of the luminosity of the hard-to-soft state transition. Therefore, the physics of the hard-to-soft state transition should be the focus.  The nearly linear relation between the luminosity of the hard-to-soft state transition and the peak luminosity of the following HS state indicates that the luminosity at which the hard-to-soft state transition occurs appears correlated with the peak luminosity of the following HS state, and is thus related to the properties of the accretion flow which powers the soft X-ray outburst. This observation challenges the interpretation that the variation in the luminosity of state transition is due to different coronal sizes (e.g. Homan et al. 2001). Since there is a tight correlation between the luminosity of the hard-to-soft state transition and the peak luminosity of the following HS state, the cause of the variation in the luminosity of the hard-to-soft state transition is also the cause of the variation of the peak luminosity in the HS state. In the HS state, the thermal spectral component dominates, and neither a static spherical corona nor a static disk corona could power the source or collapse to form a disk flow to power the source in the HS state. Therefore, coronal properties, such as coronal size or height, can not be the additional that parameter determines the spectral state transitions. The observation is also inconsistent with the interpretation invoking the history of the location of the inner-disk radius (Zdziarski et al. 2004). According to this interpretation, the smaller the initial inner disk radius, the faster a transient source reaches a hard-to-soft state transition, and thus yields a lower transition luminosity. There is no correlation between the luminosity of the hard-to-soft state transition and the peak luminosity of the HS state.  The idea that disk size affects state transitions originated from the proposed geometry of two independent flows, which is based on the long-term correlation between power-law indices and fluxes (Smith, Heindl \\& Swank 2002). According to this interpretation, the mass accretion rate in the disk flow is suspected to be an integration of that of a halo flow in the recent past which powers the LH state, because there is a {\\it significant} difference between the viscous timescales in the disk flow and the halo flow.  Although our suggestion is similar to the two-flow idea, there are significant differences. Our suggestion is based on the observations of individual state transitions and the correlation between the fluxes, which is shown in Figure 3. If the two-flow geometry is correct during the state transitions, the two flows have to be related in instantaneous mass accretion rate, instead of being related in a way involving integration of mass accretion rate over time (see also YFK04). A reduced integration effect is also suggested by the timescale of the rise of the LH state and the rise of the HS state during the start of an outburst, which differ by only a factor of 2. Therefore, the halo flow, if there is any, should be a sub-Keplerian flow which is very close to a Keplerian flow (see also YFK04). It is worth noting that instead of a halo inflow, we also regard an outflow as a possible alternative that would power the source in the LH state.  It is also known that \\mbox{Aql X-1} and other sources show ``parallel tracks\" in the color-count rate or QPO frequency-count rate plot (e.g., Mend\\'ez et al. 1998) which would be explained with a two-flow geometry (van der Klis 2003). The hard-to-soft state transitions occurring at different luminosities are actually extreme cases of the parallel tracks. Therefore, the mechanism for the variation in  the state transition luminosity is likely the same one that causes the parallel tracks.  Our current analysis determines that it is the luminosity of the hard-to-soft state transition, instead of the peak luminosity of the LH state, that is correlated with the peak luminosity of the soft X-ray outburst. Therefore, there is only an underlying positive correlation between the instantaneous mass accretion rate at the transition and that at the luminosity peak of the HS state. The relation suggests that the X-ray radiation in the LH state is related to the optically thick disk flow further out, which, at a later time, would approach the innermost region and power a HS state. Because of the positive correlation between the luminosities, the optically thick disk in the outer region is probably the ultimate source which powers the hard spectral component originating from the Comptonization of hot electrons. Our study suggests that the disk flow may split part of its mass flow to contribute to a halo flow or an outflow on its way towards the compact star. Thus, at a given time, the flux of the LH state is an integration of the contribution from the disk over a range of radii which would vary in different outbursts. We speculate that the competition between the inner disk, which radiates soft photons and cools hot electrons, and the outer disk, which probably heats hot electrons, determines the spectral transition through the process of Comptonization. Assuming the mass accretion rate is a function of time and radius, the hard-to-soft state transition probably occurs when the inner disk dominates as a result of the mass accretion rate at the outer disk declining faster than that of the inner disk or increasing slower than that of the inner disk.  The soft-to-hard state transition occurs when the outer disk dominates as a result of the mass accretion rate through the inner disk declining faster than that of the outer disk or increasing  slower than that of the outer disk. This speculation is consistent with the idea that the mass in the disk somehow determines the state transitions, since at the beginning or at the end of an outburst, most of the mass is in the disk that is farther out. However, it is worth noting that only when the disk mass accretion rate reaches a certain threshold can the disk flow survive diminishment due to a halo flow or outflow, reach the inner most region, and form a HS spectral state."
    },
    "0608/astro-ph0608437_arXiv.txt": {
        "abstract": "High resolution spectroscopy of GRB~021004 revealed a wealth of absorption lines from several intermediate ionization species. The velocity structure of the absorber is complex and material with velocity up to $\\gsim3000$~km~s$^{-1}$ is observed.  Since only the blueshifted component is observed, the absorber is very likely to be material closely surrounding the gamma-ray burst.  We use a time-dependent photoionization code to track the abundance of the ions over time. Thanks to the presence of absorption from intermediate ionization states at long times, we can estimate the location and mass of the components of the absorber.  We interpret those constraints within the hypernova scenario showing that the mass loss rate of the progenitor must have been $\\sim10^{-4} \\,M_\\odot$ per year, suggestive of a very massive star. In addition, the wind termination shock must lie at a distance of at least $100$~pc, implying a low density environment. The velocity structure of the absorber also requires clumping of the wind at those large distances. ",
        "introduction": "The high luminosity of gamma-ray burst afterglows makes them ideal sources to probe their surrounding environment through absorption lines in their powerlaw, generally featureless spectra.  In particular, time variability (or lack thereof) in the equivalent width of absorption lines can be used to constrain the size and density of the absorbing region surrounding the burst (Perna \\& Loeb 1998).  The association of GRBs with the collapse of massive stars (Stanek et al. 2003; Hjorth et al. 2003) makes absorption studies particularly useful as a new way to probe star-forming regions at intermediate and high redshifts and/or the last hundreds of years of the progenitor evolution. Lines associated with the material expelled by the star allow one to probe the velocity structure and metal content of the ejecta. The high quality and well sampled temporal evolution of the spectroscopic data gathered for GRB~021004 (Moeller et al. 2002; Matheson et al. 2003; Mirabal et al. 2003; Schaefer et al. 2003; Fiore et al. 2005; Starling et al. 2005) allows us for the first time to perform a detailed analysis.  GRB~021004 was detected by HETE II and an optical afterglow was detected within 10 minutes with an R-band magnitude of 15.34 (Fox et al. 2002).  Superimposed on the standard power-law decay, the optical light curve shows several bumps (Holland et al. 2003), which can be interpreted either as density fluctuations (Lazzati et al. 2002; Heyl \\& Perna 2003), continuous activity of the inner engine or inhomogeneities in the fireball energy content (Nakar, Piran \\& Granot 2003). High resolution spectroscopy of the optical afterglow was reported by Fiore et al. (2005).  Low-to-intermediate resolution spectroscopy was reported by Moeller et al. (2003), Mirabal et al. (2003), Matheson et al. (2003), Schaefer et al. (2003) and Starling et al. (2005).  These observations revealed a large set of absorption lines spanning a velocity range of about $3000$~km~s$^{-1}$ toward the observer.  This is the only burst to date for which such high-velocity lines have been detected.  The possibility of observing such an absorption system due to chance alignment of different absorbers was discussed and ruled out in previous work (see Schaefer et al. 2003; Mirabal et al. 2003; Fiore et al. 2005). Interpreted as absorption from the local GRB environment, the results of the spectroscopic follow-up studies are controversial.  Starling et al. (2005) conclude that the lines must be coming from a fossil stellar wind with hydrogen enrichment from a companion.  Mirabal et al. (2003) and Schaefer et al. (2003) attribute the lines to shells of material present around the progenitor. They also allow for the possibility of clumping.  Neither of those studies took adequately into account the effect of the flash ionization of the burst photons on the environment.  High resolution spectroscopy ($R\\sim50000$) was obtained by Fiore et al. (2005).  Several narrow components were resolved down to a width of a few tens of km~s$^{-1}$, showing that some of the components observed at low resolution are in fact due to blending of finer velocity structures.  Fiore et al. (2005) derived physical parameters for the absorbing medium modeling the ionization ratios with the code CLOUDY, which assumes ionization equilibrium. They obtained fairly constant ionization parameters for the different absorbers. They concluded that the absorber is stratified as a wind, since a $n\\propto{}r^{-2}$ profile gives a constant ionization parameter over a broad range of distances.  However, due to the highly time-dependent nature of the ionizing flux and the long recombination times with respect to the ionization times, their conclusions should be taken with caution.  In this paper we combine all the available spectroscopic data and adopt a time-dependent photoionization code (Perna \\& Lazzati 2002) to determine the physical conditions of the absorber.  This paper is organized as follows: in $\\S$2 we discuss the spectra of GRB~021004 and show that there is moderate evidence of time-variability in the equivalent width (EW) the observed CIV line. In $\\S$3 we present our model for the light curve of GRB~021004 while the code and the data modeling technique are discussed in $\\S$4 together with the results of the fit.  We finally discuss our findings in $\\S$5. ",
        "conclusions": "We have collected all the available spectroscopic data of GRB~021004 from the literature and compiled the time resolved evolution of the EW of CIV and SiIV absorption lines for the high velocity systems with $v\\sim3000$~km~s$^{-1}$. We have then used our time dependent photoionization and dust-destruction code (Perna \\& Lazzati 2002) to model the evolution of the column densities of these ions under the intense radiation field of the GRB and its early afterglow. Previous work on subsets of these data (Mirabal et al. 2003; Schaefer et al. 2003; Starling et al. 2005) concluded that the observed absorption can be produced by material expelled by the GRB progenitor star under the form of a smooth wind. Wolf Rayet stars are the principal candidate due to the large speed and mass loss rate of their winds ((Nugis et al.~1998; Willis 1999 ; Nugis \\& Lamers 2000).  Wolf-Rayet stars are the final evolutionary stage of massive O-type stars, with initial main-sequence masses above about 35 M$_{\\odot}$ (for solar metallicities).  These stars eject material continuously in the form of line-driven winds.  The observed mass loss rates of WC stars (those WR stars that show strong carbon lines) are $\\sim 1-6 \\times 10^{-5}M_{\\odot} {\\rm yr}^{-1}$ and they produce winds with terminal velocities 800 - 3500 km s$^{-1}$ (Nugis et al.~1998; Willis 1999 ; Nugis \\& Lamers 2000).  In the case of GRB progenitors there are not direct observations of the wind properties.  Rykoff et al. (2004) argue that their mass loss may be substantially larger. Fitting the afterglow photometric data of GRB~030418 they find that a mass loss of $\\sim10^{-3}\\,M_\\odot$~yr$^{-1}$ is required.  WR wind terminal velocities can explain the observed blueshifted velocities we see within the high velocity system without invoking any further radiative acceleration mechanisms due to the GRB itself.  In the conventional model of a wind-blown bubble formed by mass-loss from a massive star, the fast wind from a massive star expands into the slower ambient medium, which could be a fossil wind from a previous epoch, or even a molecular cloud. The fast wind sweeps up the medium into a thin, dense, swept-up shell. The high pressure behind this shell drives a reverse, or wind termination shock, into the freely expanding wind (Weaver et al.~1977; Dwarkadas 2005). Although WR stars form a late evolutionary stage of a massive star, which may undergo a large amount of mass-loss prior to the WR stage, this general picture provides a useful guide.  In this work we have shown how, by considering the flash ionization of the burst and early afterglow photons, we were able to further constrain the wind properties. In fact, the intense radiation flux fully photoionizes the gas and destroys dust particles out to a large radius. We have found that, in order for the remaining portion of the wind material at large radii to provide the required column density, the termination shock of the wind must lie at at least $\\sim100$~pc from the star and the mass loss rate from the progenitor star must have been large, of the order of $10^{-4}\\,M_\\odot$ per year.  These constraints are much more restrictive than what concluded before and allow us to infer some important properties of the progenitor star. Once again, we emphasize how this conclusion does not depend on the possible decrease of the EW of the CIV line with time, but on the simple and robust observation of the presence of high velocity CIV several days after the burst explosion.  The value of the termination shock that we have determined is in the upper range of what generally discussed theoretically (Dwarkadas 2005; Fryer, Rockefeller \\& Young 2006). However, this large value can be achieved by a very massive progenitor evolving in a low-density environment. An illustrative example is shown in Figure~\\ref{fig:starmod}, where we used a stellar evolution model provided by Meynet \\& Maeder (2003) for a 120 solar mass star. The model incorporates mass-loss from the star. We used the standard assumption that the velocity of the emitted wind is a given multiple of the escape velocity from the star at each time-step (Kudritzki \\& Puls 2000). The wind velocity and mass-loss rate are then used as input parameters to compute the evolution of the surrounding medium by means of the VH-1 code (Blondin \\& Lundqvist 1993; Dwarkadas 2005). As the star evolves, first as a main-sequence O star and then as a Wolf-Rayet star, the wind mass-loss parameters gradually change. The wind velocity, initially about 3500 km/s, drops to a low of about 1100 km/s and then increases again in the final stages to around 3000 km/s. The mass-loss rate increases from about $10^{-5} M_\\odot$ yr$^{-1}$ to about twice that value. The ambient density was chosen to be about 10$^{-3}$ particles~cm$^{-3}$, in order to have the wind termination shock form at a large radius (about 215 pc in this particular case).  Our inference of a low-density environment for the progenitor of GRB~021004 could be interpreted within the context of the recent observations that some GRB progenitors have been expelled from their parent molecular clouds (Hammer et al. 2006). High velocity in OB stars are indeed not uncommon (Mdzinarishvili \\& Chargeishvili 2005). Another possibility for the creation of a low-density environment is that the massive star progenitor of the GRB was in a star cluster which created a superbubble; these structures are commonly observed in our own Galaxy as well as in nearby ones (e.g. Heiles 1978).  \\begin{figure} \\psfig{file=lazzati_f9.ps,width=\\columnwidth} \\caption{{Characteristics of the bubble formed by the wind of a star of 120 $M_\\odot$ expanding in a medium of density $10^{-3}$ cm$^{-3}$. The solid line shows the density (left y axis) while the dashed line shows the pressure (right y axis).} \\label{fig:starmod}} \\end{figure}   Finally, we have shown that a smooth wind is insufficient to fully account for the observations (Mirabal et al. 2003; Schaefer et al. 2003; Fiore et al. 2005; Starling et al. 2005).  The structure in the velocity of the absorber suggests the presence of clumping, at least at a moderate level.  To check if extreme clumping could explain the observations, we modeled the absorber as material concentrated in a shell of width $\\Delta R=0.1R$, where $R$ is the radius of the shell. This configuration is relevant for the case of clumps with an angular size larger than the photon source (total covering\\footnote{A partial covering geometry is not allowed by the high-resolution data that show line cores with no transparency.}). We found that the fits are statistically equivalent, but the physical constraints on the mass loss rate of the progenitor star are tighter ($\\dot{M}\\gsim10^{-3}\\,M_{\\odot}$ per year). This is due to the fact that the material is more concentrated at a single radius with respect to a smooth wind.  The high resolution spectra of Fiore et al. (2005) allow us to set some constraints on the characteristics of the clumps.  Since the line opacity is complete, the clumps must be larger than the source size. In addition, since we see two clumps, the surface filling factor must be close to unity.  At the time of the observations, the transversal radius of the observable portion of the fireball is given by: \\begin{equation} R_{\\rm{fb}}(t) = c\\,t\\,\\Gamma = 3\\times10^{16} \\frac{t}{2{\\rm{days}}} \\frac{\\Gamma}{5} \\qquad {\\rm{cm}} \\end{equation} This is likely to be smaller than the dimension of clumps in the wind at a distance of $\\sim100$ parsec from the star.  It is possible that our estimate of the mass loss rate is moderately affected by the wind inhomogeneity. Consider a clumpy wind with a surface filling factor of $0.1$. In $90\\%$ of the cases no high velocity absorption would be seen.  Given that some absorption was detected, in $10\\%$ of the cases velocity structure would be present. If this were the case, our $\\dot{M}$ would be overestimated by a factor $\\sim10$.  These considerations on the presence of inhomogeneities in the wind of the GRB~021004  progenitor star should not come as a surprise. It has indeed been suggested for many years that WR winds are clumped (Moffat et al. 1988; Li et al. 2000).  Clumping in the wind has been proposed as one scenario to solve the so-called ``wind-momentum'' problem present in these outflows.  It has been noted by several authors that the momentum in the wind ${\\dot{M}v_{\\inf}}$ is several times the momentum flux (L/c) from the star (Barlow et al. 1981). Although this was initially explained away by scattering the radiation from the star more than once, it is unclear if there is sufficient opacity in the wind to accomplish this.  Clumps solve this problem in two ways 1) by increasing the density and therefore increasing the luminosity by the $n^2$ factor in the emissivity (but see Brown et al. 2004) and 2) the consequence of clumping is to reduce empirically derived mass-loss rates of hot stars by about a factor of 3, thus reducing the wind momentum itself.  In summary, the complex spectra of GRB~021004 can be explained as a result of the absorption due to a very extended wind surrounding the GRB progenitor star. It also requires a sizably large mass loss rate of $\\sim10^{-4}\\,M_\\odot$ per year. The large radial development of the wind is likely due to a combination of the high mass loss, the wind speed and of a low density environment. Some degree of clumping is also required by the structure of the absorption in the velocity space. This study shows how time resolved medium-high resolution spectroscopy of the early GRB afterglow is a powerful tool to investigate the properties of the GRB environment and constrain the characteristics of the progenitor."
    },
    "0608/astro-ph0608423_arXiv.txt": {
        "abstract": "\\noindent  We present an analysis of 20 galaxy clusters observed with the \\Chandra\\ X-ray satellite, focussing on the temperature structure of the intracluster medium and the cooling time of the gas. Our sample is drawn from a flux-limited catalogue but excludes the Fornax, Coma and Centaurus clusters, owing to their large angular size compared to the \\Chandra\\ field-of-view. We describe a quantitative measure of the impact of central cooling, and find that the sample comprises 9 clusters possessing cool cores and 11 without. The properties of these two types differ markedly, but there is a high degree of uniformity amongst the cool core clusters, which obey a nearly universal radial scaling in temperature of the form $T\\propto r^{\\sim0.4}$, within the core. This uniformity persists in the gas cooling time, which varies more strongly with radius in cool core clusters ($\\tcool \\propto r^{\\sim1.3}$), reaching $\\tcool <1\\,$Gyr in all cases, although surprisingly low central cooling times ($<5\\,$Gyr) are found in many of the non-cool core systems. The scatter between the cooling time profiles of all the clusters is found to be remarkably small, implying a universal form for the cooling time of gas at a given physical radius in virialized systems, in agreement with recent previous work. Our results favour cluster merging as the primary factor in preventing the formation of cool cores. ",
        "introduction": "\\label{sec:intro}  Clusters of galaxies are uniquely suitable environments for studying galaxy formation and evolution, as well as excellent cosmological tools for measuring the fundamental parameters of the Universe. A study of the scaling properties of clusters, which span roughly two decades in mass, is one of the best means of tackling both these objectives, and many analyses have been conducted on this basis, establishing broad trends and quantifying departures from simple self-similarity \\citep[e.g.][]{edge91,mar98b,hel00,fukazawa04}.  Increasingly, however, attention is shifting towards a more complete understanding of cluster physics, focussing in particular on the scatter in these scaling relations. With current generation telescopes, we are reaching the point where the precision of scaling relations is limited by systematic effects rather than measurement errors. Further progress requires a better knowledge of cluster physics to understand the sources of intrinsic scatter in these relations. Not only does this hold great potential for understanding galaxy evolution, it also promises to hone cluster observables as proxies for mass, which will be vital in extracting precise cosmological constraints from forthcoming large cluster surveys.  Since the cooling time of gas in the centres of most clusters is short compared to their age, radiative cooling would be expected to play an important role in cluster physics. However, many clusters do not show strong signs of cooling, possibly as a result of disruption caused by merging and associated shock heating of the gas, or else non-gravitational energy input (e.g. from AGN), or even heat redistribution via conduction or mixing. Moreover, even where cooling dominates, its progress is apparently impeded \\citep[see the recent review by][and references therein]{peterson06}, by processes which may well relate to these same phenomena. Whatever the explanation, it is clear that clusters can be separated into two distinct categories, according to the presence or absence of a `cool core' \\citep{peres98,bauer05}, where radiative heat losses have significantly lowered the central temperature of the gas. It is the combination of cool core and non cool core systems in the cluster population which accounts for much of the intrinsic scatter in scaling relations.  While cool core clusters have tended to be relatively well-studied, non-cool core clusters have received less attention in scaling relation analyses. Consequently there is great potential for exploring cluster physics by comparing the properties of the two types of clusters. This aim of this work is to investigate precisely this comparison, by focussing on a systematically selected sample of clusters, using data from the \\Chandra\\ satellite, to allow a high-resolution study of the core properties of the intracluster medium (ICM).  Previous detailed studies of the temperature structure of the hot gas in clusters have generally concentrated on the most relaxed (i.e. morphologically regular) clusters \\citep[e.g.][]{san03,piffaretti05,vikhlinin05}, which frequently possess cool cores. In this work, we have compiled a statistically selected sample of 20 clusters, to provide a more rigorous selection process. This naturally includes a larger number of non-cool core and perhaps less-relaxed clusters, which are more representative of the cluster population as a whole. We have subjected these 20 systems to a non-parametric analysis, so as to avoid any model-dependent biases in the comparison of cool core (CC) and non-CC systems. Furthermore, in restricting the analysis to X-ray data taken with the \\Chandra\\ satellite, we are able to exploit the unrivalled spatial resolution of this telescope, to permit a detailed examination of the inner regions of the ICM, where the impact of radiative cooling is most pronounced.  Throughout this paper we adopt the following cosmological parameters; $H_{0}=70$\\kmpspMpc, $\\rmsub{\\Omega}{m}=0.3$ and $\\Omega_{\\Lambda}=0.7$. Throughout our spectral analysis we have used \\XSPEC\\ 11.3.1, incorporating the solar abundance table of \\citet{grevesse98}, which is different from the default abundance table. Typically this results in larger Fe abundances, by a factor of $\\sim$1.4. All errors are 1$\\sigma$, unless otherwise stated. ",
        "conclusions": ""
    },
    "0608/astro-ph0608565_arXiv.txt": {
        "abstract": "Owing to effects arising from quantum electrodynamics (QED), magnetohydrodynamical fast modes of sufficient strength will break down to form electron-positron pairs while traversing the magnetospheres of strongly magnetised neutron stars. The bulk of the energy of the fast mode fuels the development of an electron-positron fireball. However, a small, but potentially observable, fraction of the energy ($\\sim 10^{33}$ ergs) can generate a non-thermal distribution of electrons and positrons far from the star. This paper examines the cooling and radiative output of these particles. Small-scale waves may produce only the non-thermal emission. The properties of this non-thermal emission in the absence of a fireball match those of the quiescent, non-thermal radiation recently observed non-thermal emission from several anomalous X-ray pulsars and soft-gamma repeaters.  Initial estimates of the emission as a function of angle indicate that the non-thermal emission should be beamed and therefore one would expect this emission to be pulsed as well.  According to this model the pulsation of the non-thermal emission should be between 90 and 180 degrees out of phase from the thermal emission from the stellar surface. ",
        "introduction": "\\label{sec:introduction}  Simply put magnetars are neutron stars whose magnetic fields dominate their emission, evolution and manifestations.  In the late 1970s and early 1980s,  a fleet of sensitive detectors of high-energy radiation uncovered two new phenomena, the soft-gamma repeater and the anomalous x-ray pulsar.  Strongly magnetized neutron stars provide the most compelling model for both types of object, and observations over the past few years indicate that these phenomena are two manifestations of the same type of object.   Soft-gamma repeaters exhibit quiescent emission similar to that of anomalous x-ray pulsars \\cite[e.g][]{1994Natur.368..432R,1994Natur.368..127M,1996ApJ...463L..13H,1999ApJ...510L.111H}, and anomalous x-ray pulsars sometimes burst \\cite[][]{2002Natur.419..142G,2003ApJ...588L..93K}. What makes magnetars a hot topic of research is the rich variety of physical phenomena that strong magnetic fields exhibit.  This article will focus on the quiescent emission from these interesting objects rather than the bursts themselves (The reader may wish to refer to the seminal work of \\cite{Thom95} for details of the burst but may also want to look at \\cite{Heyl03sgr} for an alternative).  Furthermore, the article will concentrate on a possible model for the recently detected non-thermal emission from these objects.  In earlier work, Lars Hernquist and I analysed wave propagation through fields exceeding the quantum critical value $B_\\rmscr{QED} \\equiv m^2c^3/e\\hbar \\approx 4.4\\times 10^{13}$ G, and demonstrated circumstances under which electromagnetic \\citep{Heyl98shocks} and some MHD waves, particularly fast modes \\citep{Heyl98mhd} evolve in a non-linear manner and eventually exhibit discontinuities similar to hydrodynamic shocks, owing to vacuum polarisation from quantum electrodynamics (QED).  In Paper I \\cite[][]{Heyl03sgr}, we developed a theory to account for bursts from SGRs and AXPs based on ``fast-mode breakdown,'' in which the wave energy is dissipated into electron-positron pairs when the scale of these discontinuities becomes comparable to an electron Compton wavelength.  We showed that, under appropriate conditions, an extended, optically thick pair-plasma fireball would result, radiating primarily in hard X-rays and soft $\\gamma$-rays.  In Paper II \\cite[][]{Heyl05sgr} we developed the theory of the non-thermal emission from the fast-mode cascade under a series of assumptions and approximations.  In particular, MHD fast modes of insufficient amplitude to generate an optically thick fireball will still dissipate through pair-production, seeding non-thermal emission. \\cite{Heyl03sgr} estimates the Thomson optical depth through the pair plasma created by a passing fast mode (their Fig.~3).  If this optical depth exceeds unity, one would expect a fireball to form.  In what follows, I extend our previous investigation of fast-mode breakdown to estimate the spectrum of non-thermal emission expected outside the region containing an optically thick fireball with particular emphasis on the assumptions made in Paper II and on the angular dependence of this emission. ",
        "conclusions": "\\label{sec:discussion}  This paper has examined the angular dependence of the non-thermal emission from AXPs and SGRs in the weak-field regime.  It is a natural extension of the model for bursts and non-thermal emission presented in Papers I and II.  When a dislocation of the surface of the neutron star is sufficiently large, the resulting fast modes will produce sufficient pairs to make the inner magnetosphere of the neutron star opaque to x-rays, generating a fireball.  Some small but observable fraction of the energy initially in the fast modes is dissipated outside the opaque region yielding a characteristic fast-mode breakdown spectrum.  A second possibility is that the crust of the neutron star is constantly shifting over small scales, generating fast modes whose breakdown is insufficient to produce a fireball.  In this case we would associate the non-thermal radiation with the quiescent thermal radiation from the surface of the star.  Both are produced by the quasi-continuous decay of the star's magnetic field.  Observations of the AXPs and SGRs continue to surprise, as do theoretical investigations of ultramagnetised neutron stars.  This and the preceding papers have presented a unified model for the thermal burst emission and non-thermal emission from ultramagnetised neutron stars.  The model has few underlying assumptions: magnetars produce fast modes sufficient to power the non-thermal emission and, more rarely, the bursts, the magnetic field far from from the star is approximately dipolar and quantum electrodynamics can account for the dynamics of pairs and photons in strong magnetic fields.  The model for the non-thermal emission from a particular fast mode depends only on the product of the strength of the wave and its wavenumber, the angle the wave propagates relative to the magnetic field direction and the strength of the dipole component of the field.  The intensity of the emission depends strongly on the direction of the initial wave. Waves that travel sufficiently close to the direction of magnetic field do not produce non-thermal emission by this mechanism at all. The emission is therefore beamed perpendicular to the magnetic field lines, so the non-thermal emission should be between ninety and one hundred eighty degrees (depending on the viewing geometry) out of phase relative to the thermal emission which is expected to be strongest along the magnetic field lines.  Further calculations of the fate of the pairs produced in the strong field regions ($\\xi \\gtrsim 0.1$) is needed to understand the spectra produced by fast-mode breakdown in general; however, the beaming of the radiation does appear robust.  Further observations can easily verify or falsify this model and potentially provide direct evidence for the ultramagnetised neutron stars that power AXPs and SGRs and the macroscopic manifestations of QED processes that account for their unique attributes."
    },
    "0608/astro-ph0608086_arXiv.txt": {
        "abstract": "We present new {\\it Hubble}, {\\it Gemini-S}, and {\\sl Chandra} observations of the radio galaxy 0313-192, which hosts a 350-kpc double source and jets, even though previous data have suggested that it is a spiral galaxy. We measure the bulge scale and luminosity, radial and vertical profiles of disk starlight, and consider the distributions of H II regions and absorbing dust. In each case, the HST data confirm its classification as an edge-on spiral galaxy, the only such system known to produce such an extended radio source of this kind. The {\\it Gemini} near-IR images and {\\sl Chandra} spectral fit reveal a strongly obscured central AGN, seen through the entire ISM path length of the disk and showing X-ray evidence of additional absorption from warm or dense material close to the central object. We consider several possible mechanisms for producing such a rare combination of AGN and host properties, some combination of which may be at work. These include an unusually luminous bulge (suggesting a black hole of mass $5-9 \\times 10^8$ $M_\\sun$), orientation of the jets near the pole of the gas-rich disk, and some evidence of a weak gravitational interaction which has warped the disk and could have enhanced fuelling of the central engine. We detect an X-ray counterpart of the kiloparsec-scale radio jet emerging to the south; jet/counterjet limits on both radio and X-ray regimes allow them to be symmetric if seen more than $15^\\circ$ from the plane of the sky, still consistent with the jet axes being within $\\sim 30^\\circ$ of the poles of the gas-rich galaxy disk. A linear or disklike emission-line structure is seen around the nucleus, inclined by $\\sim 20^\\circ$ to the stellar disk but nearly perpendicular to the jets; this may represent the aftermath of a galaxy encounter, in which gas is photoionized by a direct view of the nuclear continuum. ",
        "introduction": "The relationships between active galactic nuclei (AGN) and their host galaxies may be important not only for the evolution of the AGN, but for the host galaxies themselves. Since dynamical evidence supports the existence of massive black holes in numerous galactic nuclei, and very general arguments suggest that the active growth of such black holes by accretion is what we see as AGN, the current demographics of such black holes must result from the whole history of AGN episodes. Furthermore, some aspects of galaxy formation become easier to understand if AGN episodes regulate the properties of gas in galaxies via such feedback mechanisms as heating or wind sweeping. These factors have renewed interest in the connections between properties of AGN and their host systems.  Among the most durable systematic patterns between the occurrence of AGN and the structure of surrounding galaxies has been the finding that large double radio sources occur only around elliptical galaxies, or systems that plausibly would have been called ellipticals until a recent external disturbance (in studies ranging from Matthews, Morgan, \\& Schmidt 1964 to Heckman et al.\\ 1986). This correlation has figured in theoretical explanations of large-scale jets, and in understanding why such jets are associated with extensive synchrotron lobes. There are a handful of nearby radio galaxies showing evidence for gas-rich disks, generally in host galaxies so disturbed that these are well explained as captured material from major mergers -- 3C 120 (Arp 1975; Heckman \\& Balick 1979; Garc{\\'{\\i}}a-Lorenzo et al. 2005), 3C 293 (van Bregel et al. 1984; Floyd et al. 2006), 3C 305 (Heckman et al. 1982; Jackson et al. 2003), and 3C 285 (Roche \\& Eales 2000) are well-studied examples. In each of these, the system is disturbed enough to leave its original galaxy type ambiguous. While {\\it normal} spiral galaxies with AGN often have kiloparsec-scale radio jets; no previously reported case of a large-scale double radio source associated with a spiral has withstood detailed examination as regards both the host Hubble type and its identification with the radio source. In contrast, the edge-on galaxy associated with the radio source 0313-192 in the cluster Abell 428 obviously has a significant stellar disk (Ledlow, Owen, \\& Keel 1998), and a prominent dust lane of the kind usually associated with spiral galaxies. We report here {\\it Hubble} ACS, {\\it Gemini-S}, and {\\sl Chandra} ACIS images of this galaxy, resulting in quantitative photometric evidence that the galaxy is in fact a spiral, and pointing to peculiar features of its gas distribution which may shed light on how this rare combination of galaxy and nuclear activity happened.  As shown by Ledlow et al.\\ (2001), the radio structure of 0313-192 includes a nuclear source, a nuclear jet toward the southern lobe of about the same extent as that in M87 but several times less powerful, larger twin jets directed at projected angles of about 70$^\\circ$ to the galaxy disk plane, and edge-darkened lobes spanning roughly 280$\\arcsec$ (about 360 kpc at $z=0.067$ for ${\\rm H}_0$=70 km s$^{-1}$ Mpc$^{-1}$). This structure would typically fall into Fanaroff-Riley type I (FR I), a classification in which it would lie near the middle of the power distribution ($\\approx 10^{24}$ W Hz$^{-1}$ at 1400 MHz). Strong 21-cm H I absorption is detected against the central source at a column density ${\\rm N}_H > 10^{22}$ cm$^{-2}$, with the limit set by its large optical depth and the resolution of the data (Ledlow et al.\\ 2001). The H I absorption is also narrow in velocity span (34 km s$^{-1}$ FWHM), more characteristic of a line of sight cutting radially through a large disk than one dominated by gas closer to the central mass. This large column density is comparable to what we see toward Sgr A, supporting the idea that we are seeing a galaxy similar to our own viewed exactly edge-on. ",
        "conclusions": "These results argue convincingly that 0313-192 is in fact a spiral galaxy (albeit one which we see at a challenging orientation for morphological classification). This prompts us to consider what factors might allow such a demonstrably rare occurrence -- not only a kpc-scale radio jet, but an extensive and powerful FR I double radio source -- beginning with the salient galaxy properties.  The host galaxy of the radio source 0313-192 is a large, bright spiral, roughly Hubble type Sb. It has a luminous bulge, slightly more luminous than that of M104. The extensive disk is twice as bright as the bulge in the optical, and we can trace luminous H II regions out to at least 17 kpc from the nucleus. The dust lane is thick, richly structured including features reaching several kpc from the midplane, and is seen to nearly 20 kpc from the core. The dust lane shows a warp of several degrees with respect to the stellar disk.  We see the nucleus through an amount of foreground material typical of edge-on, gas-rich spiral disks, as seen both from soft X-ray and 21-cm H I absorption. Likewise, the nucleus is completely obscured in the optical, shortward of about 1.6 $\\mu$m. The emission lines seen in optical spectroscopy arise mostly in a roughly linear structure seen emerging behind the dust lane on both sides at an angle of about $20^\\circ$; this structure is seen nearly perpendicular to the radio jets.  If a tidal interaction has taken place, as might be suggested by the disk warp and more strongly inclined inner ionized-gas features, it must have been at large mass ratio (such as acquisition of a gas-rich former companion). The only potential companions close to 0313-192 are dwarf systems. The apparently rich environment, in projection, is part of Abell 428; this catalogued cluster shows such a large velocity dispersion that it may be largely a large-scale filament seen end-on rather than a traditional galaxy cluster. The redshift data presented by Ledlow et al.\\ (2001) indicate that this filament is broken up into a number of small groups, which could well have associated intragroup hot gas. The environment is important, since production of large-scale radio sources is fostered by interaction with an external medium. However, there is no specific evidence that the external gaseous environment around 0313-192 is unusually dense for a spiral, since a large fraction of spirals are found in groups.  We can picture several circumstances which may contribute to the rarity of extensive double sources in spiral galaxies, and to the ability of 0313-192 to produce such a source. The galaxy is large and luminous, with a bulge luminosity suggesting a central black-hole mass of order $8 \\times 10^8$ $M_\\sun$. This fits with statistical evidence that radio-loud active nuclei have an occurrence which is a strong function of the host galaxy's stellar mass (Best et al. 2005). Furthermore, the jet axis is projected roughly along the minor axis of the galaxy (at an angle of $22^\\circ$). A polar orientation would make jets more likely to escape the dense interstellar medium of a spiral rather than having their energy dumped into the ISM in a disordered way. Nearby Seyfert galaxies furnish several examples of jets on small scales which lose their collimation rather quickly through interaction with the interstellar medium. In Mkn 78, two small-scale jets are deflected or decollimated during encounters which produce emission-line features (Whittle \\& Wilson 2004). The anomalous ``radio arms\" of NGC 4258 are well described as stalled, precessing jets, whose interactions with disk gas can be observed in some detail (Cecil et al. 2000), providing a striking example of how a normal galaxy's interstellar medium can dictate the fate of jets as they propagate from the nucleus. Since the black-hole masses inferred for most Seyferts lie below the mass we suggest from the bulge properties of 0313-192, its jets may well emerge with a higher velocity than the ``weak, slow, and heavy\" jets (Whittle et al. 2004) seen in many Seyferts.  A requirement for the formation of extensive radio lobes is the escape of jets to eventually encounter an intergalactic medium, which would occur only if they avoided encountering high column densities of cool interstellar gas, possible at high angles to the plane in those spirals with bulges and nuclei that are relatively free from such gas (a condition which probably excludes the many Seyfert galaxies with surrounding star-forming regions, for example).  Aside from 0313-192, the best-attested {\\it disk} galaxy hosting luminous radio jets and lobes is the S0 galaxy NGC 612 (V{\\'e}ron-Cetty \\& V{\\'e}ron 2001). Their analysis confirmed that it is a bulge-dominated S0 system with stellar disk and dust annulus, confirming the photographic results of Ekers et al.\\ (1978). From their photometric decomposition, the bulge-to-disk ratio in NGC 612 is 2.2 in the $i$ band, much smaller than the value of 0.4 we find from modelling the images of 0313-192. In this comparison, the bulge in NGC 612 remains over twice as luminous as we find for 0313-192.   Our original interest in studying 0313-192 was to understand the circumstances which allow a spiral galaxy to produce a large double radio source with attendant jets. This detailed analysis leads us to several clues. The galaxy is large and luminous, with a particularly luminous bulge. This in turn suggests that the central black hole is unusually massive for a spiral, so that jets might emerge from its vicinity faster than found in many Seyfert galaxies. The jets emerge roughly along the poles of the disk; if this is needed for them to escape, a further geometric factor limits the number of spirals which could host such radio sources even during episodes of accretion of the black hole. A final contributing factor in this case may be a recent minor merger or weak tidal encounter, which would account for the warped disk and perhaps for more strongly inclined gaseous structure which we see within a few kpc of the nucleus."
    },
    "0608/astro-ph0608279_arXiv.txt": {
        "abstract": "We re-analyze the precision radial velocity (RV) observations of  \\hd{} (\\muarae{})  by the Anglo-Australian Planet Search Team. The star is supposed to host two Jovian companions (\\hd{}b, \\hd{}c) in long-period orbits ($\\sim 630$~days and  $\\sim 2500$~days, respectively) and a  hot-Neptune (\\hd{}d) in $\\sim 9$~days orbit. We perform a global search for the best fits in the orbital parameters space with a hybrid code employing the genetic algorithm and simplex method. The stability of  Keplerian fits is verified with the $N$-body model of the RV signal that takes into  account the dynamical constraints (so called GAMP method). Our analysis reveals a signature of the fourth, yet  unconfirmed, Jupiter-like planet \\hd{}e in $\\sim 307$~days orbit. Overall, the global architecture of four-planet  configuration recalls the Solar system. All companions of \\muarae{} move in quasi-circular orbits.  The orbits of two inner Jovian planets are close to the 2:1~mean motion resonance. The alternative three-planet system involves two Jovian planets in eccentric orbits ($e\\sim 0.3$), close to the 4:1~MMR, but it yields a significantly worse fit to the data. We also verify a hypothesis of the 1:1~MMR in the subsystem of two inner Jovian planets in the four-planet model. ",
        "introduction": "The star \\hd{} (\\muarae{}) is a Sun-like main-sequence dwarf monitored by the long-term, precision radial velocity (RV) surveys.  It has been observed over more than 7~years by the Anglo-Australian Telescope (AAT)  Planet Search Team and by the Geneva Planet Search Team (CORALIE and HARPS spectrometers).  The work of the AAT team lead to the discovery of a Jupiter like companion \\hd{}b in about 630~days orbit \\citep{Butler2001}. One year later, \\cite{Jones2002a} confirmed the Jovian planet and discovered a linear trend in the RV data revealing  a signature of the second, more distant body. In the next paper, \\cite{McCarthy2004} published a new  orbital solution with the orbital period of the long-period planet \\hd{}c about 3000~days and the eccentricity $e_{\\idm{c}} \\sim 0.57$. The same year, \\cite{Santos2004}, using observations done with the ultra precise HARPS spectrometer,  announced $\\sim 14$ Earth-mass planet \\hd{}d in $\\sim 9$~days orbit.  That discovery is a breakthrough in the field as the long-term precision of spectrometers approaches $1$~m/s. Actually, the instrumental errors are much smaller than the RV  variability (stellar jitter) induced by the Sun-like stars themselves. Recently, \\cite{Butler2006} published a new updated set of 108 observations of \\muarae{}, spanning 2551~days (about 7.5~yr). Thanks to the  updates in the UCLES instrument installed at the AAT and the software pipeline \\citep{Butler2006}, the long-term precision of the measurements is amazing. It also reaches  $1$~m/s at the end of the observational window and is kept at the mean level of $\\sim 2.8$~m/s over its whole length.  This paper is devoted to the analysis of the RV observations of \\muarae{} using the so called GAMP approach (an acronym of genetic algorithm with MEGNO penalty) that makes explicit the use of Newtonian, self-consistent $N$-body model of the stars' reflex motion, dynamical properties of the planetary system and the Copernican Principle \\citep{Gozdziewski2003e,Gozdziewski2005a}. This algorithm makes it possible to  derive meaningful bounds on the elements of the outermost planet in spite of the fact that the data cover only a part of the  longest orbital period. Without the dynamical constrains, the kinematical as well as pure Newtonian best fits to the \\muarae{} RV tend to show large eccentricity of the outermost companion and then the system becomes catastrophically  unstable.  We verify the results of the previous papers based on a much smaller number of relatively less accurate RV observations.  The results of the analysis of the new RV data set greatly improved and extended over time by \\cite{Butler2006} lead us to a conclusion that the \\muarae{} may host {\\em four} planets in quasi-circular orbits, including a new, Jupiter-like object in $\\sim 307$~days orbit (\\hd{}e)\\footnote{In this paper we use a widely adopted (but unofficial) convention for naming the extrasolar planets with lower-case Roman  letters starting from ``b'', in the order of their announcement.}. The new best fit solution describes the orbital architecture of this extrasolar system as very different from the previous ones.  Shortly after submitting the manuscript, we learned about an independent work by \\cite{Pepe2006} who announced  a very similar orbital solution and  also the fourth planet in the \\muarae{} system. These authors study a  different set of the RV measurements but the conclusions of the two papers are in an excellent agreement. Still, in this paper we focus our attention on the analysis of the AAT data along  the lines of the originally submitted manuscript. However, the last Section~5 is  devoted to a preliminary study of the full available data set, including the new  measurements published by \\cite{Pepe2006}. In particular, we verify a hypothesis  of the 1:1~MMR in the subsystem of two Jovian planets. % ",
        "conclusions": "The long-term radial velocity surveys of the Sun-like stars constantly reveal more and more exciting features of the planetary systems.    The \\muarae{} system may be the second known four-planet configuration, after 55~Cnc \\citep{McArthur2005}.  Remarkably, in such a multi-planet  system the orbits are close to circular ones, similarly to three-planet systems of HD~37124 \\citep{Vogt2005} and HD~68930 \\citep{Lovis2006} resembling the Solar system architecture. The alternative best-fit three-planet configurations  may contain two Jupiter-like planets in the 4:1~MMR.  In that case, the eccentric orbits ($e_{\\idm{b,c}}\\sim 0.3)$ would be localized in a dynamically active region of the phase space; in fact, on the edge of dynamical stability. Besides the worse fit quality, it could be a heuristic argument against the three-planet system. Obviously, the key for the proper understanding of the orbital architecture is the improved precision of observations and their extended time span. Curiously, the new data reveal a Jovian planet  that has the orbital period two times shorter than the companion already detected  a few years ago. It remains an open question whether the two inner companions (the planets b and~e) in the four-planet configuration are locked in the 2:1~MMR. Most likely, the presence of the massive companion~c prevents the creation of the 2:1~MMRs island known in the other four extrasolar systems  (Gliese~876, HD~82943, HD~128311, and HD~73526) presumably locked in this resonance. In any case, even the proximity of their orbits to such particular dynamical state  can be counted as the fifth occurrence of the 2:1~MMR among $\\sim 20$ multi-planet systems known up to date. It may indicate a universal dynamical mechanism governing the creation and orbital evolution of extrasolar planetary systems.  We also found stable configurations related to  the 1:1~MMR of the inner Jovian planets in eccentric orbits. However, such fits are significantly worse than those derived for the system with quasi-circular orbits. The results of our analysis permit us to conclude that a few years of observations are still required to constrain the outermost orbit without a doubt.   Looking at the orbits of the \\muarae{} planets  and recalling the results of \\cite{Laskar2000}, we see a free space in the range of $\\sim (0.2,0.8)$~AU in which  new planetary objects may yet exist. Extensive numerical simulations concerning the less constrained parameters of the outermost planet would be necessary to answer this question."
    },
    "0608/hep-th0608152_arXiv.txt": {
        "abstract": "We generalize the calculation of cosmic superstring reconnection probability to non-trivial backgrounds.  This is done by modeling cosmic strings as wound tachyon modes in the 0B theory, and the spacetime effective action is then used to couple this to background fields.  Simple examples are given including trivial and warped compactifications. Generalization to $(p,q)$ strings is discussed. ",
        "introduction": "There has been a recent surge of interest in the possibility that superstrings may be observed as cosmic strings stretched across the sky \\cite{Witten:1985fp} \\cite{Polchinski:2004ia} \\cite{Copeland:2003bj} \\cite{Becker:2005pv}.  An important parameter in their observation is the intercommutation probability $P$, the likelihood that two cosmic strings approaching each other will reconnect to form a kinked shape.  This quantity determines  how many strings will be present in the present universe, and may also serve to differentiate what type of cosmic string it is \\cite{Jackson:2004zg}.  Additionally, since superstrings are highly sensitive to any background present, they may also provide information about e.g. extra dimensions.  While order-of-magnitude estimates of how backgrounds would influence the intercommutation probability have been made, no rigorous technique has existed.  The purpose of the present article is to develop a simple method for calculating tree-level reconnection probability in slowly-varying backgrounds.  The driving motivation behind this study is that by examining the relationship between background and intercommutation probability, measurement of the latter could yield information about the former. ",
        "conclusions": "In this article we have generalized the computation of basic cosmic superstring reconnection probabilities for a general background.  We have considered only very simple examples but it would be instructive to study more sophisticated backgrounds and to higher order in $\\alpha'$.  It would be especially interesting to consider nonfactored amplitudes, as this would allow a sensitive probe of the extra dimensions at energies inaccessible to accelerators.  Ideally we would be able to use the same technique to calculate the effect of the background for all types of cosmic superstrings, fundamental as well as $(p,q)$ \\cite{Schwarz:1995dk}; as it stands, the methods are completely different \\cite{Jackson:2004zg} \\cite{Hanany:2005bc}.  Though we have used the effective action of the perturbative 0B string theory, it is known that IIB can accomodate the nonperturbative $(p,q)$ strings, so we conjecture that the 0B action should be capable of this in our approach.  The main difference would seem to be that the tachyon would couple to both the NSNS and RR two-forms, though it is unclear precisely how this would work since there are now two sets of RR fields.  In this case the wavefunction for the $(p,q)$ string is to leading order in $g_s$ simply $\\psi(y) = \\delta(y-y_0)$, where $y_0$ will depend on the charges $p$ and $q$ coupled to the two-form background fields.  Finally the effect of the background may also prove important in string and brane gas studies \\cite{Easther:2002mi} \\cite{Easther:2004sd}, where the decompactification rate of the universe depends on winding mode interaction rates."
    },
    "0608/astro-ph0608035_arXiv.txt": {
        "abstract": "We report the detection of a Jupiter--mass planet in a \\period~d orbit around the \\mstar~\\msun\\ subgiant \\star. The eccentricity of \\star b ($e = \\ecc$) is unusually large compared to other planets within 0.1~AU of their stars. Photometric observations demonstrate that the star is constant to $\\pm$0.0001 mag on the radial velocity period, strengthening our interpretation of a planetary companion. This planet was detected as part of our radial velocity survey of evolved stars located on the subgiant branch of the H--R diagram---also known as the Hertzsprung Gap. These stars, which have masses between 1.2 and 2.5~\\msun, play an important role in the investigation of the frequency of extrasolar planets as a function of stellar mass. ",
        "introduction": "The assemblage of $\\sim 185$ known extrasolar planets\\footnote{For the updated catalog of extrasolar planet and their parameters see http://exoplanets.org.} has revealed important relationships between the physical characteristics of stars and the likelihood that they harbor planets \\citep{udry03, butler06}. One example is the strong correlation between planet occurrence and stellar metallicity \\citep{gonzalez97, santos04, fischer05b}. The planet--metallicity relationship can be understood in the context of the core accretion model of planet formation, in which Jovian planets begin as large rocky cores that then accrete gas from their surrounding protoplanetary disks once they exceed a critical core mass \\citep[e.g.][]{pollack96}. The growth of these embryonic cores is enhanced by increasing the surface density of solid particles in the disk, which is related to the metallicity of the star/disk system \\citep{alibert05, ida05a}.  The growth rate of rocky cores in protoplanetary disks is also related to the total disk mass. Assuming that the disk mass increases with the mass of the central star, the planet occurrence rate should therefore also correlate with stellar mass. \\citet{laughlin04} showed that the lower surface densities of M~dwarf protoplanetary disks impede the growth of Jupiter--mass planets. The relationship between stellar mass and planet occurrence was studied in further detail by \\citet{ida05b} for a larger range of stellar masses. Based on their Monte Carlo simulations, they predict a positive correlation between the number of detectable planets and stellar mass up to about 1~\\msun. However, there is some debate about whether surface density of solid material in protoplanetary disks is proportional to the mass of the central star. By assuming that the initial conditions of the disk are independent from the mass of star, \\citet{kornet06} find that the occurrence rate of planets is inversely related to stellar mass.  The prediction that low--mass stars should have a lower frequency of Jovian planets is in accordance with observations; only one M~dwarf, GL~876, is known to harbor Jupiter--mass planets \\citep{marcy01}. From their survey of 90 M~dwarfs, \\citet{endl06} estimate that less than 1.27\\% of such low--mass stars harbor Jovian--mass planets within 1~AU. \\citet{laws03} studied the planet occurrence rate in light of the current observational data using the larger set of FGK stars surveyed as part of the California \\& Carnegie Planet Search (CCPS, the full, updated target list can be found in \\citet{wright04b})). They found evidence that the planet rate decreases for lower stellar masses and peaks near 1.0~\\msun. However, the planet rate is poorly constrained for stellar masses greater than about 1.2~\\msun\\ due to the small number of intermediate--mass stars ($1.3 \\lesssim M_{star} \\lesssim 3$~\\msun) in the CCPS sample. This dearth of massive stars is due to an observational bias since main--sequence stars with spectral types earlier than F8 tend to be fast rotators \\citep{donascimento03}, have fewer spectral lines, and display a large amount of  chromospheric activity \\citep{saar98}. These features result in a decrease in the radial velocity precision attainable from the spectra of main--sequence stars more massive than $\\sim 1.2$~\\msun \\citep{galland05}. Thus, early--type dwarfs are not typically monitored as part of most radial velocity surveys.  One method to circumvent these difficulties is to observe intermediate--mass stars after they evolve into the region of the H--R diagram between the main--sequence and red giant branch---also known as the Hertzsprung Gap (HG). After stars have expended their core hydrogen fuel sources their radii expand, their photospheres cool, and convection sets in below their photospheres. Convective motion and stellar rotation drive magnetic fields that couple with an expanding stellar wind and act as a rotational brake \\citep{gray85, schrijver93, donascimento00}. The cooler atmospheres and slower rotational velocities of evolved stars lead to an increased number of narrow absorption lines in their spectra, making HG stars better suited for precise radial velocity measurements than their main--sequence progenitors.  We are conducting a radial velocity survey of 159 HG stars to search for planets orbiting intermediate--mass stars. We describe the selection criteria of our sample of stars in \\S~\\ref{sample}. We present here the first planet detection from our sample of HG stars: an eccentric hot jupiter orbiting the \\mstar~\\msun\\ subgiant, \\star. The properties of the host star are presented in \\S~\\ref{properties}, and we describe our observations and orbital solution in \\S~\\ref{orbit}. We present the results of our photometric monitoring in \\S~\\ref{photometry} and conclude with a discussion in \\S~\\ref{discussion}. ",
        "conclusions": "\\label{discussion}  We are monitoring the radial velocities of a sample of 159 intermediate--mass Hertzsprung Gap (HG) stars in order to study the relationship between stellar mass and planet occurrence rate. We present the detection of a \\msini~\\mjup\\ planet in an eccentric, \\period~d orbit around the \\mstar~\\msun\\ subgiant \\star.  Compared to other planets with orbital separations $a<0.1$~AU (also known as ``hot jupiters''), \\star b has an unusually large eccentricity. Figure~\\ref{ecc_plot} shows the distribution of eccentricities for the 33 planets having $a < 0.1$~AU \\citep{butler06}. With $e = \\ecc$, \\star\\ b stands out from the distribution as one of only three hot jupiters with eccentricities in excess of 0.2. The other two planets are HD~162020\\ b \\citep[$e = 0.277$;][]{udry02} and HD~118203\\ b \\citep[$e = 0.309$;][]{dasilva06}.   \\begin{figure}[t!] \\epsscale{1} \\plotone{f4.eps} \\figcaption{The eccentricity--period distribution of hot jupiters listed in \\citet{butler06}. \\star\\ is shown as a filled pentagram. The inset shows the histogram of eccentricities for the same sample of stars. The arrow denotes the position of \\star. \\label{ecc_plot}} \\end{figure}  The origin of eccentricities amongst single exoplanets is not well understood, but the gravitational influence of additional companions can drive the eccentricities of planets in multicomponent systems. An as--yet undetected second planet is likely the cause of the large eccentricity of HD~118203b---\\citet{dasilva06} report a linear trend with a 49.7~\\ms~yr$^{-1}$ slope. However, we do not see evidence of a linear trend in our observations of \\star, nor significant periodicities in the Keplerian fit residuals. We are continuing to monitor this star at Keck and Lick Observatories to search for additional low--mass companions.  It is thought that the nearly circular orbits of most hot jupiters is a result of tidal circularization. This leaves open the possibility that the tidal circularization timescale at the orbital distance of \\star b is longer than the age of the star. Indeed, theoretical predictions of the circularization timescale have a strong dependence with semi--major axis, typically $t_{circ} \\propto a^{13/3}$ \\citep{terquem98}. Eight of the 10 planets with $a > 0.055$~AU have eccentricities greater than 0.1, which may indicate that the timescale for tidal circularization at distances beyond $\\sim 0.055$~AU is longer than the age of most FGK stars. The detection of additional hot jupiters by programs such as the N2K Consortium \\citep{fischer05a} will help shed light on the relationship between the eccentricity and orbital separation of short--period planets.  Other searches for planets around intermediate--mass stars have so far focused on the red giant branch \\citep{frink02, hatzes05} and clump regions of the H--R diagram \\citep{sato03, setiawan03}. These programs have to date discovered a total of 6 sub-stellar objects orbiting giant stars, proving that planets do form and can be detected around evolved stars more massive than $\\sim 1.5$~\\msun. However, stars in the clump and red giant branches follow closely--spaced and often overlapping evolutionary tracks, making precise stellar mass estimations from isochrone interpolation difficult (see Figure~\\ref{sg_hr}). By contrast, HG stars have mass tracks that are nearly parallel and widely--spaced in $M_V$, enabling precise mass determinations. HG stars also exhibit lower velocity jitter and have smaller radii than red giants \\citep{frink01, hekker06}, which enables the detection of a wider range of planet masses and orbital separations. We will present the stellar characteristics and velocity behavior of our sample in a forthcoming paper."
    },
    "0608/astro-ph0608203_arXiv.txt": {
        "abstract": "{} {We present two-dimensional spectroscopy of the extremely metal-deficient blue compact dwarf (BCD) galaxy SBS 0335--052E aiming to study physical conditions, element abundances and kinematical properties of the ionised gas in this galaxy.} {Observations were obtained in the spectral range $\\lambda$3620 -- 9400\\AA\\  with the imaging spectrograph GIRAFFE installed on the UT2 of the Very Large Telescope (VLT). These observations are the first ones carried out so far with GIRAFFE in the ARGUS mode which allows to obtain simultaneously 308 spectra covering a 11\\farcs4$\\times$7\\farcs3 region.} {We produced images of SBS 0335--052E in the continuum and in emission lines of different stages of excitation. We find that while the maximum of emission in the majority of lines, including the strong lines H$\\beta$ 4861\\AA, H$\\alpha$ 6563\\AA, [O {\\sc iii}] 4363,5007\\AA, [O {\\sc ii}] 3726,3729\\AA, coincides with the youngest south-eastern star clusters 1 and 2, the emission of He II 4686\\AA\\ line is offset to the more evolved north-west clusters 4, 5. This suggests that hard ionising radiation responsible for the He II $\\lambda$4686\\AA\\ emission is not related to the most massive youngest stars, but rather is connected with fast radiative shocks. This conclusion is supported by the kinematical properties of the ionised gas from the different emission lines as the velocity dispersion in the He II $\\lambda$4686\\AA\\ line is systematically higher, by $\\sim$ 50\\% -- 100\\%, than that in other lines. The variations of the emission line profiles suggest the presence of an ionised gas outflow in the direction perpendicular to the galaxy disk. We find a relatively high electron number density $N_e$ of several hundred cm$^{-3}$ in the brightest part of SBS 0335--052E. There is a small gradient of the electron temperature $T_e$ and oxygen abundance from the East to the West with systematically higher $T_e$ and lower 12+log O/H in the western part of the galaxy. The oxygen abundances for the whole H II region and its brightest part are 12 + log O/H = 7.29 $\\pm$ 0.02 and 7.31 $\\pm$ 0.01, respectively. We derive the He mass fraction taking into account all systematic effects. The He mass fraction $Y$ = 0.2463 $\\pm$ 0.0030, derived from the emission of the whole H II region, is consistent with the primordial value predicted by the standard big bang nucleosynthesis model. We confirm the presence of Wolf-Rayet stars in the cluster 3.} {} ",
        "introduction": "The blue compact dwarf (BCD) galaxy SBS 0335--052E is an excellent nearby laboratory for studying star formation in low-metallicity environments. Since its discovery as one of the most metal-deficient star-forming galaxies known \\citep{I90}, with oxygen abundance 12 + log O/H $\\sim$ 7.30 \\citep{M92,I97b,I99,TI05}, SBS 0335--052E has often been proposed as a nearby young dwarf galaxy \\citep{I90,I97b,T97,P98,P04}. \\citet{T97} and \\citet{P98}, using the same Hubble Space Telescope (HST) images, have found several luminous clusters. The brightest clusters are labelled in Fig. \\ref{fig1} which represents the highest spatial resolution archival UV HST/Advanced Camera for Surveys (ACS) image of SBS 0335--052E obtained by \\citet{K03}. Some of the clusters are very young and produce extended regions of ionised gas \\citep{M92,I97b,P98,P04}. In particular, \\citet{I01b}, using deep long-slit spectra of SBS 0335--052E, have shown that extended H$\\alpha$ emission is detected over $\\sim$ 6--8 kpc, suggesting that hot ionised gas is spread out far away from the central part of the galaxy.   \\begin{figure}% \\hspace*{1.0cm}\\psfig{figure=s0335UV_1.ps,angle=0,width=6.0cm}% \\caption{Archival HST/ACS UV image of SBS 0335--052E with the labelled compact clusters. North is up and East is to the left. } \\label{fig1} \\end{figure}    \\begin{table*}[t] \\caption{Journal of observations \\label{tab1}} \\begin{tabular}{lccrccc} \\hline Date       & Setup & Wavelength& Resolving &Exposure,     & Airmass$^{\\rm a}$& Seeing$^{\\rm a}$   \\\\ &       & range, \\AA& power $R$ &   s        &        &    \\\\ \\hline 19 Nov 2003& LR1 & 3620--4081       &12800&2$\\times$1200 &1.066, 1.065    &1\\farcs3, 0\\farcs9  \\\\ 20 Nov 2003& LR2 & 3964--4567       &10200&2$\\times$1200 &1.082, 1.062    &0.8, 0.8      \\\\ 20 Nov 2003& LR3 & 4501--5078       &12000&2$\\times$1350 &1.061, 1.080    &1.0, 0.7      \\\\ 21 Nov 2003& LR4 & 5010--5831       & 9600&2$\\times$1200 &1.295, 1.170    &0.8, 0.8      \\\\ 21 Nov 2003& LR5 & 5741--6524       &11800&2$\\times$1500 &1.070, 1.066    &1.3, 0.7      \\\\ 21 Nov 2003& LR6 & 6438--7184       &13700&2$\\times$1200 &1.067, 1.103    &1.0, 0.7      \\\\ 21 Nov 2003& LR7 & 7102--8343       & 8900&2$\\times$1200 &1.286, 1.486    &0.8, 0.9      \\\\ 22 Nov 2003& LR8 & 8206--9400       &10400&2$\\times$1200 &1.076, 1.061    &0.7, 0.8      \\\\ \\hline \\end{tabular}  $^{\\rm a}$The first value is at start of exposure, the second value is at end of exposure.  \\end{table*}   \\citet{TI97} using HST/GHRS UV spectrum of SBS 0335--052E have discovered a very broad Ly$\\alpha$ line in absorption suggesting that this galaxy is embedded in a large envelope of neutral gas. The column density of $N$(H {\\sc i}) = 7$\\times$10$^{21}$ cm$^{-2}$ in SBS 0335--052E derived by \\citet{TI97} is the largest one known for the BCDs. Later, \\citet{P01} using Very Large Array (VLA) observations in the line H {\\sc i} $\\lambda$21 cm have detected a large neutral gas cloud around SBS 0335--052E with a size 66 by 22 kpc elongated in the east-west direction and with two maxima separated by 22 kpc. The first maximum in H {\\sc i} distribution is connected to SBS 0335--052E, and the second one to the companion dwarf galaxy SBS 0335--052W discovered by \\citet{P97}. The latter galaxy is shown by \\citet{I05a} to be the lowest metallicity emission-line galaxy known with 12 + log O/H = 7.12 $\\pm$ 0.03. \\citet{T05} using Far Ultraviolet Spectroscopic Explorer (FUSE) observations have found that the oxygen abundance of the neutral gas around SBS 0335--052E of 12 + log O/H $\\sim$ 7.0 is only slightly lower than that of the ionised gas, implying that this galaxy was not formed from pristine gas.  Despite the low metallicity of SBS 0335--052E, an appreciable amount of dust has been detected in it. \\citet{I97b} and \\citet{T97} have found variations of extinction in this galaxy from the optical spectroscopic and photometric observations. Later, \\citet{T99} and \\citet{H04} using Infrared Space Observatory (ISO) and Spitzer mid-infrared observations have found an emission of an appreciable amount of warm dust with a characteristic temperature of $\\sim$ 100K. Even hotter dust with a temperature of several hundred degrees is expected to be present in SBS 0335--052E which is indicated in the infrared spectra at shorter wavelengths, namely 2--4 $\\mu$m \\citep{V00,H01}.  Since the discovery paper by \\citet{I90} it is known that the ionising radiation in SBS 0335--052E is hard. The presence of the high-ionisation He {\\sc ii} $\\lambda$4686\\AA\\ emission line \\citep[e.g. ][]{I90,M92,I97b} and of the [Fe {\\sc v}] $\\lambda$4227\\AA\\ emission line \\citep{I01b,F01} suggests that radiation with photon energies greater than 4 Rydberg is intense and could not be explained by stellar emission. Furthermore, \\citet{I01b} and \\citet{TI05} have found [Fe {\\sc vii}] and [Ne {\\sc v}] emission lines which require intense radiation with photon energies above 7 Rydberg. SBS 0335--052E has also been detected in the X-ray range \\citep{T04}. The origin of the hard ionising radiation remains unclear. Several mechanisms such as radiation from most massive main-sequence stars, Wolf-Rayet stars, high-mass X-ray binaries and radiative shocks have been discussed by e.g. \\citet{G91}, \\citet{S96}, \\citet{I01b}, \\citet{I04} and \\citet{TI05}. The most recent investigations have shown that although the stellar origin of hard radiation is not completely excluded, the most likely source of hard radiation is fast radiative shocks.  Despite the efforts of different groups in studying the properties of SBS 0335--052E and its evolutionary status many problems remain unsolved. Since this galaxy is possibly a young galaxy it could be considered as a local counterpart of the high-redshift young dwarf galaxies. Therefore the continuation of its studies is important for cosmological applications. In this paper we present a two-dimensional spectroscopic study of SBS 0335--052E with the VLT/GIRAFFE. These are the first observations carried out so far in the ARGUS mode. Two new features are characteristic for these new observations which were not present in all previous spectroscopic studies of this galaxy. First, new observations allow to map the whole galaxy in different emission lines and in the continuum. This gives integrated characteristics of different emission lines for the whole H {\\sc ii} region, such as the line luminosities, which are necessary input parameters for building up the model of the H {\\sc ii} region. Second, the spectral resolution of new observations is by one order of magnitude better than in all previous spectroscopic observations of SBS 0335--052E and it is enough to obtain the intrinsic profiles of emission lines. This allows to study the kinematics of the H {\\sc ii} region, and to make a comparison of kinematic characteristics in regions of different ionisation stages. In particular, it is important to compare the He {\\sc ii} $\\lambda$4686\\AA\\ line profiles with the profiles of other emission lines and, hence, to make conclusions concerning the origin of hard radiation in SBS 0335--052E.  In \\S\\ref{sec:obs} we describe the observations and data reduction. Morphology of SBS 0335--052E in different emission lines and continuum is considered in \\S\\ref{sec:morph}. Kinematic properties are discussed in \\S\\ref{sec:vel}. Heavy element abundances and helium abundance are derived in \\S\\ref{sec:abund}. The Wolf-Rayet stellar population in SBS 0335--052E is discussed in \\S\\ref{sec:wr}. Our conclusions are summarised in \\S\\ref{sec:concl}. ",
        "conclusions": "}  In this paper we present panoramic spectroscopic observations with the VLT/GIRAFFE in the spectral range $\\lambda$3620--9400\\AA\\ of one of the most metal-deficient blue compact dwarf (BCD) galaxies, SBS 0335--052E. Our findings can be summarized as follows:  1. The morphology of the galaxy in different lines is very similar, except for the He {\\sc ii} $\\lambda$4686\\AA\\ line, suggesting that the main ionising source in this galaxy is the compact super-star clusters (SSC) 1+2 in the SE part of the galaxy, most likely the cluster 1. The equivalent number of O7{\\sc v} ionising stars in this cluster is $\\sim$13000, making it the most powerful SSC known so far.  2. The maximum emission of the He {\\sc ii} $\\lambda$4686\\AA\\ line is offset to the older compact clusters 4 and 5. Additionally, its width is 1.5 -- 2 times greater than the width of other lines. Furthermore, the width of this and other emission lines is much higher in clusters 4,5 than in cluster 1. These facts imply that the hard ionising radiation responsible for He {\\sc ii} emission is most likely produced by fast radiative shocks. However, some contribution from WR stars located in cluster 3 cannot be excluded.  3. The analysis of kinematical properties of the ionised gas in SBS 0335--052E suggests the presence of gas outflow in the direction perpendicular to the galaxy disk.  4. We derive physical parameters in the H {\\sc ii} region of SBS 0335--052E, suggesting that the ionised gas is hot ($\\sim$20000K) and dense ($\\ga$ several hundred particles per cm$^{-3}$). The oxygen abundances 12 + log O/H in the brightest region and in the integrated spectrum of SBS 0335--052E are respectively 7.31$\\pm$0.01 and 7.29$\\pm$0.02, and they are consistent with previous determinations. The other heavy element-to-oxygen abundance ratios are consistent with the average values derived for other most metal-deficient galaxies. There is a slight decrease of oxygen abundance from the East to the West suggesting some self-enrichment of the ionised gas by heavy elements.  5. We derive the He mass fraction $Y$ from the integrated spectrum taking into account all possible systematic effects in the He abundance determination. Our value $Y$ = 0.2463 $\\pm$ 0.0030 is slightly lower but consistent within the errors with the primordial He mass fraction determined from the 3 year WMAP data and from primordial deuterium measurements.  6. We confirm the presence of Wolf-Rayet stars in the cluster 3 found previously by \\citet{P06}. The lower limits of the WNL and WC4 star numbers are $\\sim$ 35$\\pm$8 and $\\sim$ 3$\\pm$1, respectively if standard WR line luminosities are assumed."
    },
    "0608/astro-ph0608459_arXiv.txt": {
        "abstract": "{ During an analysis of optical spectra of 80 young star clusters in several nearby spiral galaxies, $[$\\ion{O}{iii}$]$ and $[$\\ion{N}{ii}$]$ emission lines were noted in some cases.  Three of these emission line sources are associated with clusters older than 30 Myrs, and are identified as likely planetary nebula (PN) candidates. } { These objects may represent a rare opportunity to study PNe whose progenitor stars are known to be of intermediate masses, although detailed analysis is challenging because of the underlying strong continuum from the cluster stars. This paper presents and discusses basic properties of the PN candidates and their host clusters. } { Based on the observed emission line fluxes, the excitation parameters and luminosities of the nebulae are derived. This allows a crude placement of the central stars in two of the objects on the H-R diagram. Cluster ages and masses are estimated from broad-band colours and by fitting model SSP spectra to the observed spectra. } { The two PN candidates where central star luminosities and temperatures can be estimated are found to be consistent with post-AGB model tracks for a central star mass of about 0.60 $M_{\\odot}$. One of the host clusters has an age of 32--65 Myrs, corresponding to a main sequence turn-off mass of $M_{\\rm TO}$ = 6.6--9.0 M$_\\odot$. For the other cluster the age is 282--407 Myrs, corresponding to M$_{\\rm TO}$ = 3.2--3.6 M$_\\odot$. By estimating the number of stars evolving off the main sequence per year, a total of 6 PNe are expected in the full sample of 80 clusters for a PN lifetime of $10^4$ years. The factor of two disagreement with the actual observed number may be due, among other things, to uncertainties in PN lifetimes. It is interesting to note that all three PN candidates are associated with clusters which are more diffuse than average. } { While PNe have previously been found in some old globular clusters, the candidates identified here are among the first identified in \\emph{young} star clusters. } ",
        "introduction": "Star clusters and planetary nebulae (PNe) can both be observed at distances well beyond the Local Group, and are thus valuable extragalactic population tracers. However, much remains to be understood about the post-main sequence evolution of low- and intermediate stars leading up to the PN stage, partly because of the difficulty in determining the distances and progenitor stellar masses of planetary nebulae in the Milky Way. For PNe associated with stellar clusters the initial mass of the progenitor star can be assumed to be that corresponding to the main sequence turn-off (unless the progenitor star is a close binary in which mass exchange has occurred), but the number of PNe known to be associated with stellar clusters remains small. Less than a handful of such objects are known to exist in Milky Way open clusters (Pedreros \\cite{ped87,ped89}; O'Dell \\cite{odell63}) and four PNe have been found in Galactic globular clusters (M15, M22, Pal 6, NGC~6441; Jacoby et al.\\ \\cite{jac97}).  The distance problem can be alleviated by studying PNe in external galaxies with well determined distances. Many such searches have been carried out or are currently under way (Kwok \\cite{kwok00}; Feldmeier \\cite{feld06}; Magrini \\cite{mag06}), but the detected candidates are typically field objects and the initial mass thus remains unknown. PNe have been identified in a few GCs in the nearby giant elliptical NGC~5128 (Minniti \\& Rejkuba \\cite{mr02}; Rejkuba et al.\\ \\cite{rej03}), and the significant number of on-going large spectroscopic surveys of extragalactic globular cluster systems are likely to reveal many more candidates. However, by studying PNe in ancient globular clusters one remains fundamentally limited to low-mass stars. One way to sample a larger range of stellar masses would be to collect a sample of PN candidates in younger stellar clusters. So far, the number of such objects are limited to the small number of PNe in Milky Way open clusters.  We have recently obtained optical spectra for a number of young star clusters in several nearby galaxies. The original aim of this programme was to carry out a detailed test of simple stellar population (SSP) models, but during the analysis we noted emission lines in some of the spectra. As discussed below, the clusters in question span a range of ages from a few tens of Myrs to several hundreds of Myrs, corresponding to main sequence turn-off masses between 3 and 9 M$_{\\odot}$.  These objects potentially represent a valuable addition to the very few cases where PNe have been associated with young star clusters, and could provide a unique opportunity to study PNe where the distances and initial masses are known with some accuracy.  The full spectroscopic dataset and analysis will be presented in a forthcoming paper, and here we concentrate on the few spectra where PN candidates were identified. ",
        "conclusions": "In a sample of optical spectra of 80 young extragalactic star clusters, 3 emission line objects were detected which we identify as likely PN candidates. Two of these have radial velocities consistent with those of their host clusters, while the third displays a shift of about 150 km/s and may be a field object falling on the slit by chance. The analysis of the candidate PN spectra is complicated by the relatively low spectral resolution and the strong underlying stellar spectra. This is a problem especially for the Balmer lines, which are very strong in absorption in the cluster spectra and thus mask the comparatively weak line emission from the PNe.  However, it was found that the line emission can be reasonably well separated by subtracting out synthetic SSP model spectra from the library of Gonz{\\'a}lez-Delgado (\\cite{gon05}).  All three objects have [\\ion{O}{iii}] emission line fluxes consistent with those expected from PNe. For the two objects where the association with clusters is most robust (both in NGC~5236), the excitation parameters and total luminosities were estimated and the central stars were placed on the H-R diagram. Comparison with post-AGB evolutionary tracks from Vassiliadis \\& Wood (\\cite{vw94}) indicates central star masses of about 0.60 $M_\\odot$.  Based on stellar evolutionary timescales, we estimate that a total of about 6 PNe are expected in the 80 star clusters for a PN lifetime of $10^4$ years. Instead, 2--3 are detected here, which we consider reasonable agreement. Scaling the empirical estimate of the PN specific frequency from Magrini et al.\\ (\\cite{mag03}) predicts a total of 4 PNe in our sample, although the age distribution of our PN host clusters is likely different from that of the Local Group stellar populations on which the Magrini et al.\\ estimate is based.  A systematic search for PNe in a larger sample of extragalactic star clusters might put useful constraints on the PN lifetimes and the range of stellar masses that produce PNe. High-dispersion spectroscopy would allow a better subtraction of the underlying cluster spectra and thus more accurate measurements of the emission lines, and would provide constraints on the kinematic properties (expansion velocities) of the nebulae. Further observations would also provide constraints on time variability, and thus help distinguish between PNe and other possibilities (e.g. symbiotic stars)."
    },
    "0608/astro-ph0608173_arXiv.txt": {
        "abstract": "The Chandra X-ray Observatory is providing fascinating new views of massive star-forming regions, revealing all stages in the life cycles of massive stars and their effects on their surroundings.  I present a {\\em Chandra} tour of some of the most famous of these regions: M17, NGC~3576, W3, Tr14 in Carina, and 30~Doradus.  {\\em Chandra} highlights the physical processes that characterize the lives of these clusters, from the ionizing sources of ultracompact HII regions (W3) to superbubbles so large that they shape our views of galaxies (30~Dor).  X-ray observations usually reveal hundreds of pre-main sequence (lower-mass) stars accompanying the OB stars that power these great HII region complexes, although in one case (W3 North) this population is mysteriously absent.  The most massive stars themselves are often anomalously hard X-ray emitters; this may be a new indicator of close binarity.  These complexes are sometimes suffused by soft diffuse X-rays (M17, NGC~3576), signatures of multi-million-degree plasmas created by fast O-star winds.  In older regions we see the X-ray remains of the deaths of massive stars that stayed close to their birthplaces (Tr14, 30~Dor), exploding as cavity supernovae within the superbubbles that these clusters created. ",
        "introduction": "High-resolution X-ray images from the Chandra X-ray Observatory and XMM-Newton elucidate all stages in the life cycles of massive stars -- from ultracompact HII (UCHII) regions to supernova remnants -- and the effects that those massive stars have on their surroundings. X-ray studies of massive star-forming regions (MSFRs) thus give insight into the massive stars themselves, the accompanying lower-mass cluster population, new generations of stars that may be triggered by the massive cluster, interactions of massive star winds with themselves and with the surrounding neutral medium, and the fate of massive stars that die as cavity supernovae inside the wind-blown bubbles they created.  In the era of {\\em Spitzer} and excellent ground-based near-infrared (IR) data, why are X-ray studies of MSFRs important?  {\\em Chandra} routinely penetrates $A_V > 100$~mag of extinction with little confusion or contamination, revealing young stellar populations in a manner that is unbiased by the presence of inner disks around these stars.  These vast pre-main sequence (pre-MS) populations are easily seen by {\\em Chandra} and the $17^\\prime \\times 17^\\prime$ imaging array of its Advanced CCD Imaging Spectrometer (ACIS-I); a typical ACIS-I observation of a nearby ($D<3$~kpc) MSFR, lasting 40--100 ksec, finds 500--1500 young stars, tracing the initial mass function (IMF) down to $\\sim 1$M$_\\odot$. These observations can increase the number of known cluster members by as much as a factor of 50 in poorly-studied regions \\citep[e.g.\\ NGC~6357,][]{Wang06}.  In the first 7 months of 2006, at least 18 refereed papers on X-ray observations of MSFRs (using {\\em Chandra} or {\\em XMM}) were published or submitted.  These include studies of 30~Doradus \\citep{Townsley06a,Townsley06b}, NGC~6334 \\citep{Ezoe06}, $\\sigma$ Orionis \\citep{Franciosini06}, Westerlund~1 \\citep{Skinner06,Muno06}, RCW~38 \\citep{Wolk06}, Cepheus~B/OB3b \\citep{Getman06}, NGC~6357 \\citep{Wang06}, M17 \\citep{Broos06}, NGC~2362 \\citep{Delgado06}, NGC~2264 \\citep{Flaccomio06,Rebull06}, Trumpler 16 in Carina \\citep{Sanchawala06}, NGC~6231 \\citep{Sana06}, the Orion Nebula Cluster \\citep{Stassun06}, the Arches and Quintuplet clusters \\citep{DWang06}, and W49A \\citep{Tsujimoto06}.  Early {\\em Chandra} and {\\em XMM} studies were summarized in \\citet{Townsley03}; more recent work is described in \\citet{Feigelson06}.  Other MSFRs currently under study at Penn State include RCW~49 \\citep{Townsley04}, NGC~7538 \\citep{Tsujimoto05}, W51A \\citep{Townsley04,TownsleyIAUS227}, and the Rosette Nebula (Wang et al.\\ 2006, in preparation).  This contribution begins with recent X-ray results on the young MSFRs M17 and NGC~3576, two regions at different evolutionary stages but with notably similar outflows of hot, X-ray-emitting gas.  Next I introduce a new ACIS mosaic of the Westerhout 3 complex, showing how triggered star formation in the same molecular cloud can result in very different stellar clusters.  Then I address the rich and complicated point source and diffuse X-ray emission seen towards Trumpler~14 in Carina.  Finally I provide a first look at a new {\\em Chandra} observation of 30 Doradus obtained in early 2006. ",
        "conclusions": "High-resolution X-ray studies of MSFRs reveal magnetically-active pre-MS stars, massive star microshocks, and colliding-wind binaries.  {\\em Chandra} provides a disk-unbiased stellar sample for IR study.  Diffuse X-ray emission is also seen, caused by O star winds interacting with themselves and with the surrounding media. Supernovae occurring in the cavities blown by OB clusters produce brighter soft X-rays, dominating the X-ray emission in large multi-cluster and multi-generation star-forming complexes.  Several {\\em Chandra} discoveries are noted in this short overview. For all regions, early O stars often show anomalously hard spectra, perhaps indicating close binarity.  For individual regions, discoveries include:\\\\ $\\bullet$ M17:\tFirst clear detection of an X-ray outflow.\\\\ $\\bullet$ NGC 3576:  Bipolar bubble seen in {\\em MSX} data is likely the wind-blown bubble from an embedded OB cluster; this field shows the second clear example of X-ray outflowing gas; hard X-rays reveal cluster members not seen in the IR -- these are likely the missing ionizing massive stars.\\\\ $\\bullet$ W3:  W3 Main is huge compared to other W3 young clusters; the W3 Main IRS5 and W3(OH) massive protostars are very hard X-ray sources; W3 North is missing its cluster.\\\\ $\\bullet$ Tr14:  Some early O stars are likely binary or magnetically active, some are not; soft diffuse emission appears between Tr14 and Tr16, while separate, unabsorbed soft diffuse emission located far from massive stars may indicate a cavity supernova remnant inside the Carina complex's young superbubble.\\\\ $\\bullet$ 30 Dor:  More spatial structure and point sources emerge in the longer observation; no clear sign of hard X-rays from a cluster wind in R136 are seen."
    },
    "0608/astro-ph0608490_arXiv.txt": {
        "abstract": "{This manuscript reviews recent progress in our understanding of the nucleosynthesis of medium and heavy elements in supernovae. Recent hydrodynamical models of core-collapse supernovae show that a large amount of proton rich matter is ejected under strong neutrino fluxes. This matter constitutes the site of the $\\nu p$-process where antineutrino absorption reactions catalyze the nucleosynthesis of nuclei with $A>64$.  Supernovae are also associated with the r-process responsible for the synthesis of the heaviest elements in nature. Fission during the r-process can play a major role in determining the final abundance patter and in explaining the almost universal features seen in metal-poor r-process-rich stars.}  \\FullConference{International Symposium on Nuclear Astrophysics - Nuclei in the Cosmos - IX\\\\ 25-30 June 2006\\\\ CERN}  \\begin{document} ",
        "introduction": "\\label{sec:intro}  In recent years, observations of metal-poor stars have contributed to increase our understanding of the nucleosynthesis of medium and heavy nuclei and its evolution during the history of the galaxy.  Metal-poor stars with large enhancements of r-process elements (the abundance of Eu is typically considered to represent the presence of heavy r-process nuclei) with respect to iron show a variation of two to three orders of magnitude in the absolute amount of r-process elements present for stars of similar metallicities~\\cite{Cowan.Sneden:2006}. However the relative abundance of elements heavier that $Z>56$ (but not including the radioactive actinides) shows a striking consistency with the observed solar abundances of these elements~\\cite{Cowan.Sneden:2006}.  This consistency does not extend to elements lighter than $Z=56$ where some variations are observed.  In most of the cases elements lighter than $Z<56$ are underabundant when compared with a scaled solar r-process abundance distribution that matches the observed heavy element abundances~\\cite{Cowan.Sneden:2006}.  However, recent observations of the metal-poor star HD 221170~\\cite{Ivans.Simmerer.ea:2006} show that in some cases the agreement between the scaled solar r-process abundance pattern and the observed abundances of elements can be extended to elements heavier than $Z>37$. All these observations indicate that the astrophysical sites for the synthesis of light and heavy neutron capture elements are different~\\cite{Wasserburg.Busso.Gallino:1996,Qian.Wasserburg:2001} suggesting two disting r-processes.  Possible sites are supernovae and neutron-star mergers. The exact site and operation for both types of r-process is not known, however, there are clear indications that while the process responsible for the production of heavy elements is universal~\\cite{Wanajo.Ishimaru:2006} the production of lighter elements (in particular Sr, Y and Zr) has a much more complex Galactic history~\\cite{Travaglio.Gallino.ea:2004}.  Even if the astrophysical site of the r-process(es) is (are) unknown, it is clear that the process is of primary nature. This means that the site has to produce both the neutrons and seeds necessary for the occurrence of a phase with fast neutron captures that characterizes the r-process~\\cite{Cowan.Thielemann.Truran:1991}.  Moreover, in order to explain the observed abundances of U and Th the neutron-to-seed ratio needs to be larger than $\\sim 100$. Under these conditions fission of r-process nuclei beyond U and Th can play a mallor role in explaining the universality of the heavy r-process pattern in metal-poor stars. This issue will be discussed in section~\\ref{sec:fission}.  In section~\\ref{sec:nupprocess} we present a new nucleosynthesis process that we denote the $\\nu p$-process which occurs in proton-rich matter ejected under explosive conditions and in the presence of strong neutrino fluxes. This process seems necessary to explain the observed abundances of light p-nuclei, including $^{92,94}$Mo and $^{96,98}$Ru. ",
        "conclusions": "\\label{sec:conclusions}  The study of the nucleosynthesis processes responsible for the production of medium and intermediate elements and their relationship to supernovae constitutes a challenge to astronomers, astrophysicists and nuclear physicists. Our current understanding is driven by high-resolution spectroscopic observations of metal-poor stars that aim to probe individual nucleosynthesis events. At the same time progress in the modeling of core-collapse supernovae has improved our knowledge of explosive nucleosynthesis in supernovae. In particular the presence of proton-rich ejecta has open the way to find a solution to the long-standing problem of the origin of light p-nuclei. Further progress will come from advances in the modeling of the supernovae explosion mechanism and from improved knowledge of the properties of the involved nuclei to be studied at future radioactive-ion beam facilities. These facilities will also open the door to the study many of the nuclei involved in r-process nucleosynthesis, in particular the nuclei located near the $N=82$ waiting point that are important in determining the role of fission in the r-process.  They will also provide valuable data needed to constrain theoretical models to allow for more reliable extrapolations to the region of the nuclear chart where fission takes place during r-process nucleosynthesis."
    },
    "0608/gr-qc0608134_arXiv.txt": {
        "abstract": "We study the properties of a future singularity encountered by a perfect fluid observer in tilting spatially homogeneous Bianchi cosmologies. We derive the boost formulae for the Weyl tensor to establish that, for two observers that are asymptotically null with respect to each other, their respective Weyl parameters generally both tend to zero, constant, or infinity together. We examine three classes of typical examples and one exceptional class. Given the behaviour of the Weyl parameter, we can predict that the singularity encountered is a Weyl singularity or a kinematic singularity. The analysis suggests that the kinematic variables are also useful in indicating a singularity in these models.   \\ PACS numbers: 98.80.Jk, 04.20.-q ",
        "introduction": "A cosmological model ($\\mathcal{M},\\mathbf{g},\\mathbf{u}$) is defined by specifying the spacetime geometry, determined by a Lorentzian metric $\\mathbf{g}$ defined on the manifold $\\mathcal{M}$, and a family of fundamental observers, whose congruence% \\footnote{A congruence is a family of curves such that through each point there passes precisely one curve in this family.} of world-lines is represented by the 4-velocity field $\\mathbf{u}$. The covariant derivative $u_{a;b}$ of the 4-velocity field is decomposed into \\emph{kinematic} variables according to \\be u_{a;b} = \\sigma_{ab} + \\omega_{ab} + H (g_{ab} + u_a u_b) - \\dot{u}_a u_b, \\ee where $\\sigma_{ab}$ is the rate of shear tensor, $\\omega_{ab}$ is the vorticity tensor, $H$ is the Hubble scalar, and $\\dot{u}_a$ is the acceleration vector. Sometimes there is another preferred family of fundamental observers, $\\hat{\\mathbf{u}}$, whose covariant derivative defines a corresponding set of kinematic variables $(\\hat{\\sigma}_{ab}, \\hat{\\omega}_{ab}, \\hat{H}, \\hat{\\dot{u}}_a)$, which are different from the first set.  The study of spatially homogeneous (SH) Bianchi cosmologies forms part of the effort to understand the underlying dynamics of anisotropic and inhomogeneous universes. In recent papers \\cite{CHL06,CHL06b} we examined the future asymptotic dynamics as experienced by the perfect fluid observer in Bianchi cosmologies with a tilted perfect fluid with linear equation of state $p=(\\gamma-1)\\mu$, focussing on physical properties of the models as seen by the fluid observer. We found that given a Bianchi type, for a large enough value of the equation of state parameter, $\\gamma$, the perfect fluid observer can encounter a singularity, in spite of the fact that the SH observers, moving geodesically and orthogonally to the spatially homogeneous hypersurfaces, do not encounter any singularity into the future \\cite{Rendall}. Although there is no universally accepted definition of a singularity~\\cite[Section 9.1]{Wald}, we shall make use of the congruence in defining a singularity within the context of a cosmological model: a singularity is encountered by a congruence if the proper time is finite and inextendible. This definition is consistent with the general definition adopted in the early work on Bianchi cosmologies by Ellis \\& King, which states that a singularity is recognized by the existence of an inextendible curve which has a finite length as measured by a generalized affine parameter \\cite[page 121]{EK74}. \\footnote{Many also use the stricter definition of a singularity as geodesic incompleteness. However, in our case this is not satisfactory since the fluid does not follow geodesics; thus we will not require the congruence to be geodesic.} The encounter of a singularity is accompanied by an extreme tilt limit (i.e., the fluid observer become asymptotically null with respect to the SH observer). The singularity encountered by the fluid observer has the following properties:% \\footnote{Here, we refer to the properties of variables associated with the orthonormal frame of the fluid observer.} \\begin{itemize} \\item\tThe proper time needed to reach the singularity is finite, \\item\tThe Hubble scalar $H$ and some other kinematic quantities diverge, \\item\tThe matter density tends to zero (i.e., the singularity is not a matter singularity).% \\footnote{Since the source is a perfect fluid with a linear equation of state, this means that all components of the Ricci tensor also tend to zero.} \\end{itemize} All components of the Weyl curvature tensor converge for some cases and at least one component diverges to infinity for other cases. In this paper we shall further examine the behaviour of the components of the Weyl tensor (according to the fluid observer), with the aim to predict whether they converge or diverge. In order to do so, we recall the following concept regarding the Weyl curvature tensor.  \\subsection*{The Weyl parameter}  The Weyl tensor is defined as the trace-free part of the Riemann curvature tensor and, consequently, is not directly involved in the Einstein field equations. In some sense, the Weyl tensor describes the `free gravitational field' and therefore is of particular interest. Penrose \\cite{penrose} suggested that the Weyl tensor is related to a measure of a `gravitational entropy' and therefore could be used to shed light on the initial state of the Universe. In this context, the Weyl tensor and its Weyl scalars have been used to characterise different behaviours of non-tilted Bianchi models \\cite{wa,gw,BHWeyl}.  Wainwright, Hancock and Uggla \\cite{WHU99} studied the late-time dynamics of Bianchi type VII$_0$ cosmologies with a non-tilted perfect fluid. They introduced a quantity $\\mathcal{W}$ called the `Weyl parameter' (see \\cite{WHU99}, p 2580): \\be \\label{Weyl_param} \\mathcal{W} = \\frac{W}{H^2},\\quad \\text{where} \\quad W^2 = \\frac16(E_{ab}E^{ab} + H_{ab}H^{ab}). \\ee They coined the term% \\footnote{In \\cite{BHWeyl} the term `extreme Weyl dominance' was used for the case $\\mathcal{W} \\rightarrow \\infty$ while `Weyl dominance' was used for the case when the Weyl curvature invariant dominated the Ricci invariant but $\\mathcal{W}$ was bounded.}  `Weyl curvature dominance' to describe the phenomenon in which $\\mathcal{W} \\rightarrow \\infty$. This means that the ratio of some components of the Weyl curvature tensor and the square of the Hubble scalar tends to infinity.% \\footnote{It is worth noting that both $W$ and $H$ tend to zero in \\cite{WHU99}.} This phenomenon also occurs in Bianchi type VII$_0$ and VIII cosmologies with a tilted perfect fluid \\cite{VII,VIII}.  In \\cite{WHU99,VII,VIII}, the Weyl parameter is defined with respect to the SH observer. However, the Weyl parameter can also be defined with respect to different observers, and particularly the fluid observer in tilted Bianchi cosmologies.% \\footnote{The Weyl parameter, unlike the invariant Weyl scalars, is an observer-dependent quantity.}  Although the limits of the Weyl parameter along two congruences of worldlines are in principle different, we conjecture that in general they both tend to zero, constant, or infinity together. We shall refer to this as `the Weyl parameters have the same convergence/divergence property'. The argument is considered in two cases: whether or not the two congruences are asymptotically null with respect to each other ($v^2 \\rightarrow 1$).  If the two congruences are not asymptotically null with respect to each other ($v^2 \\not\\rightarrow 1$), then $\\G = (1-v^2)^{-1/2}$ is bounded. Then, from the boost formulae (\\ref{boost_Eab})--(\\ref{boost_Hab}) for the Weyl components, in general the Weyl components $\\hat{C}_{abcd}$ and $C_{abcd}$ relative to two different congruences have the same convergence/divergence property. Also recall from the boost formula for the Hubble scalar $H$ (\\cite{CHL06}, equation (B.15), reproduced below) \\be \\label{boost_H} \\hat{H} = \\frac13 \\left[ \\ehat_\\mu(\\Gv^\\mu) - \\frac{\\G}{\\G+1} v^\\mu \\ehat_\\mu(\\G) +3\\G H  - 2 \\G a_\\mu v^\\mu + \\G \\udot_\\mu v^\\mu \\right] \\ee that $\\hat{H}$ and $H$ also have the same convergence/divergence property. It then follows from (\\ref{Weyl_param}) that generally the Weyl parameters $\\hat{\\mathcal{W}}$ and $\\mathcal{W}$ have the same convergence/divergence property.  In the more interesting case for this analysis, when the two congruences are asymptotically null with respect to each other ($v^2 \\rightarrow 1$), then $\\G \\rightarrow \\infty$ and, in general, $\\hat{C}_{abcd}$ becomes of order $\\G^2 C_{abcd}$. From (\\ref{boost_H}) we see that $\\hat{H}$ is of order $\\G H$ as $\\G \\rightarrow \\infty$. As a result, generally we have \\be \\label{CC} \\frac{\\hat{C}_{abcd}}{\\hat{H}^2} \\quad \\text{is of order} \\quad \\frac{C_{abcd}}{H^2}; \\ee i.e., the Weyl parameter $\\hat{\\mathcal{W}}$ of a second congruence is generally of the same order as the Weyl parameter $\\mathcal{W}$ of the first congruence, and they therefore have the same convergence/divergence property. Exceptions are possible when cancellation occurs in the leading order terms in all the Weyl components, as we will show occurs in the LRS Bianchi type V cosmologies. Indeed, we will show that, when the argument works, we can use it to predict the type of singularity encountered by the fluid observer. ",
        "conclusions": "\\label{sec:discussion}  The future asymptotic dynamics in tilting SH Bianchi cosmologies has recently been studied, with an emphasis on the physical properties of the models as experienced by the fluid observer \\cite{CHL06,CHL06b}. In this paper we have been primarily interested in the properties of the singularity that the perfect fluid observer can encounter, with particular emphasis on the behaviour of the Weyl tensor (according to the fluid observer). As we have seen in the examples, Weyl blow-up is not a good indicator of a singularity. A better measure is the blow-up of kinematic variables.  We now discuss the work by Collins and Ellis \\cite{CE79} and Ellis and King \\cite{EK74} on Weyl blow-up at the singularity. In Collins and Ellis \\cite{CE79}, it is claimed that the future singularity in the LRS Bianchi type V example is a Weyl singularity (i.e., a `conformal singularity', see page 97), in that some components of Weyl tensor diverge. This is not the case, for the only component of the Weyl tensor is $E_+$, and it tends to zero. All components of the Riemann tensor tend to zero at the singularity. Some kinematic variables diverge, and are responsible for the singular behaviour.    Ellis and King \\cite{EK74} classify singularities as `curvature singularities', `locally extendible singularities' or `intermediate singularities'. Curvature singularities correspond to our Weyl singularities with the blow-up of at least one of the Weyl scalars; locally extendible singularities correspond to our kinematic singularities; and intermediate singularities correspond to our Weyl singularities with bounded Weyl scalars. According to this classification, all of our examples of Weyl singularities are intermediate singularities. We prefer the terminology `kinematic singularity' over `locally extendible singularity' since in our examples the fluid worldlines are certainly {\\em not} extendible at the singularity, due to the blow up of the kinematic variables; indeed, our examples are not `locally extendible' in the sense of Clarke \\cite{C73}.    \\subsection*{Conclusion}  In discussions regarding singularities, usually attention is focussed on the Weyl tensor components and the matter density. We have argued that in a cosmological context more emphasis should be put on the kinematic variables.  In this paper we have derived the boost formulae for the components of the Weyl tensor. We have applied the boost formulae to establish that, for two observers that are asymptotically null with respect to each other, their respective Weyl parameter generally both tend to zero, constant, or infinity together. We examined three classes of typical examples and one exceptional class. Given the behaviour of the Weyl parameter, we can predict that the singularity encountered is a Weyl singularity or a kinematic singularity. The examples suggest that the kinematic variables are more useful than the Weyl tensor components in indicating a singularity."
    },
    "0608/astro-ph0608345_arXiv.txt": {
        "abstract": "The existence of deconfined quark matter in the superdense interior of neutron stars is a key question that has drawn considerable attention over the past few decades. Quark matter can comprise an arbitrary fraction of the star, from 0 for a pure neutron star to 1 for a pure quark star, depending on the equation of state of matter at high density. From an astrophysical viewpoint, these two extreme cases are generally expected to manifest different observational signatures. An intermediate fraction implies a hybrid star, where the interior consists of mixed or homogeneous phases of quark and nuclear matter, depending on surface and Coulomb energy costs, as well as other finite size and screening effects. In this brief review article, we discuss what we can deduce about quark matter in neutron stars in light of recent exciting developments in neutron star observations. We state the theoretical ideas underlying the equation of state of dense quark matter, including color superconducting quark matter. We also highlight recent advances stemming from re-examination of an old paradigm for the surface structure of quark stars and discuss possible evolutionary scenarios from neutron stars to quark stars, with emphasis on astrophysical observations. ",
        "introduction": "In the aftermath of a core-collapse supernova of massive stars 8 solar masses (8~\\Mo) and above, the iron core of the progenitor star implodes from a radius of $\\sim 1000$~km to a compact object of radius $\\sim 15$~km that is neutron-rich and bound by gravitational forces. A neutron star (NS) is thus born. In 1934, Baade and Zwicky\\cite{Baade} hypothesized the supernova-neutron star connection, but did not specify any observational signals associated to the formation of a neutron star. Thirty years later, Bell and Hewish\\cite{Bell} made the serendipitous discovery of the first radio pulsar, interpreted shortly thereafter by Gold\\cite{Gold} as a rotating neutron star. Since then, over 1500 similar ``rotation-powered'' neutron stars have been identified\\cite{Manchester}, with spin periods ranging from few milliseconds to several seconds. In addition, less common categories such as the anomalous X-ray pulsars (AXPs), soft gamma-repeaters (SGRs) and X-ray dim isolated neutron stars (XDINS) now exist and their relation to the ``garden-variety'' pulsars is not well-understood. In some cases, by monitoring the pulsed and thermal radiation from these diverse sources, important physical quantities such as mass, radius, spin-period, magnetic fields, age and surface black-body temperature for the neutron star can be inferred.  The main aim of neutron star observations is to understand the underlying structure of the neutron star and its observed thermal and magnetospheric emissions. From a theoretical standpoint, the study of neutron star structure involves understanding the nature of matter over an enormous density range from $\\rho\\sim 10^{2}-10^{15}$ g/cm$^3$. The neutron star {\\it atmosphere} and {\\it envelope}, having density $\\rho\\leq 10^4$ g/cm$^3$ constitute a tiny fraction of the star's volume, but can cause significant deviation from the emergent Planckian photon flux depending on the atmospheric composition and magnetic field strength. The {\\it outer crust} of the neutron star, from $\\rho\\sim 10^{4}-10^{11}$g/cm$^3$, is composed chiefly of increasingly neutron-rich nuclei starting with $^{56}$Fe, that are embedded in an ordered lattice structure (at zero temperature) that minimizes the Coulomb interaction energy. The presence of free electrons that form a degenerate Fermi sea ensures charge neutrality and stability against $\\beta$-decay. In the {\\it inner crust}, somewhat above $\\rho\\sim 10^{11}$g/cm$^3$, neutrons drip out of nuclei, leading to a mixture of the ordered phase with a degenerate gas of neutrons. As densities approach the saturation density ($\\rho_0\\sim 2.6\\times 10^{14}$g/cm$^3$) of nuclear matter, the fraction of free neutrons increases while the nuclei may form extended non-spherical shapes leading to rod/slab-like ``pasta'' phases (2-D voids) and ``Swiss cheese'' phases (3-D voids). The {\\it outer core} of the neutron star, from $\\rho\\sim 0.5-2\\rho_0$, consists of an admixture of neutrons, protons and electrons in accordance with charge neutrality and $\\beta$-equilibirum. Other negatively charged leptonic or hadronic species can appear when their ground state energy becomes lower than their respective chemical potential dictated by $\\beta$-equilibrium. In the {\\it inner core} ($\\rho\\geq 2\\rho_0$), with increasing density, negatively charged bosons can form Bose-condensates and quarks may be deconfined, resulting in a free quark-gluon phase where the quark pairs may also condense. These ideas are summarized in several existing reviews\\cite{Heiselberg,Page06,Weber06}.  Currently, our constraints in performing ab initio calculations of strongly interacting dense matter severely limits accurate formulations of the equation of state (EoS) at higher densities, leading to theoretical uncertainties that preclude a definitive statement about the phase of matter in the interior of neutron stars. However, models of the EoS that address pure or mixed phases can be constructed and its predictions compared to astrophysical observations to judge the viability of the model. This brief review is aimed at assessing the evidence for the existence of quark matter inside neutron stars, given the current status of neutron star observations. We set the stage by recounting recent progress in theoretical studies of dense quark matter. ",
        "conclusions": "It is apparent that recent neutron star observations disfavor the appearance of soft exotic (non-hadronic) matter at high density. They do not as yet completely rule out all quark matter equations of state. Hybrid stars may exist only as a small population among neutron stars and should not be dismissed until long-term evolution of neutron stars and transient phenomena occuring at neutron star surfaces are understood within a consistent neutron star picture. In particular, it is premature to rule out quark matter in neutron stars on the basis of mass-radius measurements alone."
    },
    "0608/astro-ph0608549_arXiv.txt": {
        "abstract": "The Atacama Cosmology Telescope is a project to map the microwave background radiation at arcminute angular resolution and high sensitivity in three frequency bands over substantial sky areas. Cosmological signals driving such an experiment are reviewed, and current progress in hardware construction is summarized. Complementary astronomical observations in other wavebands are also discussed. ",
        "introduction": "The full-sky microwave maps from the WMAP satellite have provided a definitive measurement of the microwave background temperature on angular scales down to around 15 arcminutes, in five frequency bands from 23 to 90 GHz \\cite{ben03}. These measurements determine the primordial power spectrum of fluctuations in the baryon-photon plasma of the early universe, resulting in tight constraints on the basic cosmological model of the universe \\cite{wmap_params}  and bringing to fruition a decade of anticipation \\cite{jun96}. While additional measurements will provide refined constraints on some cosmological parameters, they are unlikely to alter qualitatively our fundamental picture of the universe which has emerged.  Significant experimental attention has now turned to microwave background temperature measurements at smaller angular scales. At scales smaller than 10 arcminutes, the power in primary fluctuations produced at the surface of last scattering begins to drop precipitously, due to diffusion damping on length scales smaller than the thickness of the last scattering surface. But secondary fluctuations in microwave temperature at arcminute angular scales are induced by interaction of the microwaves with large-scale structure at redshifts between zero and a few. In constrast with large-angle temperature measurements, which image fluctuations in the early universe, high-resolution temperature measurements mainly probe the subsequent development of structure in the universe. The Atacama Cosmology Telescope (ACT) project \\cite{kos03,fow05,act_site} is one effort to map the microwave background at angular resolutions better than two arcminutes, at sensitivities sufficient to detect all major effects contributing to the temperature fluctuations at these scales. The rest of this introductory section will review the specific signals of interest; then an overview of the ACT experiment and its current status will be given, along with plans for complementary astronomical observations in other wave bands.  \\subsection{Power Spectrum}  The power spectrum of primordial perturbations, usually characterized by the power law index $n$, has been measured by WMAP over a range of multipole moments from $l=2$ to $l=700$ \\cite{wmap_params}.  Further constraints on $n$ come from combining WMAP data with other measures of power at smaller scales (e.g., \\cite{sel05}), but these heterogeneous approaches are susceptible to systematic errors; the microwave background provides the cleanest way to measure the power spectrum. ACT will extend the range of precision multipole measurements out to the point where the primary fluctuations cease being the dominant contribution to the temperature fluctuations, around $l=3500$.  In principle, ACT will be able to measure $n$ over nearly an additional decade in scale, significantly increasing the precision with which $n$ can be constrained. Measurements over a larger range of angular scales also increase sensitivity to any departure from a perfect power law, often characterized by the ``running'' of the spectral index $dn/d\\ln k$ \\cite{kt95}.  Determination of $n$ and its running are the most powerful current sources of information  about inflation or other fundamental processes producing fluctuations in the very early universe. In practice, the precision in $n$ will be determined by the accuracy of the relative calibration of ACT maps at 145 GHz with WMAP maps at 90 GHz and much lower resolution. We are currently studying the issue of relative calibration using sky and instrument simulations outlined below.  \\subsection{Gravitational Lensing}  Mass density variations along the path of microwave photon propagation result in the photon experiencing a varying gravitational potential, and thus deflections in the direction of the potential gradient. This gravitational lensing effect is always small in cosmological contexts; the propagation direction of an average microwave photon will be deflected by an angle of tens of arcseconds over its journey of 13 billion light years. These deflections are also coherent over larger scales up to a degree, so that detecting them requires more work than simply having angular resolution better than the detection angle. The original temperature fluctuations are a gaussian random field to a good approximation; gravitational lensing will induce a particular additional four-point correlation function into the temperature map \\cite{ber97}. Detecting these characteristic correlations requires comparing temperature fluctuations in the maps at small scales of a few arcminutes with larger scales around a degree \\cite{oka02,kes03,hir03}. The overall amplitude of the temperature signal is on the order of the mean temperature difference on the deflection scale, which is on the order of five $\\mu$K.  In addition to lensing from large-scale structure distributed evenly across the sky, galaxy clusters also give localized gravitational lensing signals. If the cluster is in front of a background microwave temperature gradient, the lensing produces a dipole-like pattern aligned with the temperature gradient \\cite{sel00}. The separation of the hot and cold peaks is on the order of one arcminute and the peak-to-peak amplitude can be of order 10 $\\mu$K for a large cluster in front of a substantial gradient \\cite{hol04,dod04,val04}. Such signals are in principle observable with ACT's sensitivity and resolution, but other signals associated with clusters will make reliable lensing mass estimates difficult.   \\subsection{Thermal Sunyaev-Zeldovich Effect} The thermal Sunyaev-Zeldovich (SZ) Effect describes the energy boost received by low-energy microwave background photons Compton scattering from high-energy electrons in hot, ionized gases \\cite{zel69}. The SZ effect has so far been detected towards a number of galaxy clusters through pointed observations in the direction of known clusters (see \\cite{car02} for a review); ACT and similar experiments aim to detect thousands of clusters via blind surveys of large sky areas. The effective temperature shift of the microwave background blackbody spectrum has the characteristic form of a Compton y-distortion; below the ``null\" at a frequency of around 217 GHz the radiation flux is decreased, while above 217 GHz the flux is increased, and the amplitude is proportional to the pressure of the ionized gas integrated along the path of photon propagation. The largest galaxy clusters have SZ effective temperature distortions as large as 100 $\\mu$K, making them the largest amplitude signals expected at ACT frequencies.  Unlike clusters selected via X-ray or optical observations, the cluster SZ signal is essentially independent of cluster redshift for a cluster of a given mass and gas distribution. The SZ signal, which depends linearly on the electron density, is less sensitive to internal gas structure than the X-ray emission, which is quadratic in the electron density. Selecting clusters via their SZ distortion is also physically much simpler than selecting clusters optically via the light from their resident galaxies, since the SZ signal is directly related to the majority of the baryonic cluster mass, while the optical light comes from galaxies which represent only a small fraction of the baryonic mass, and whose light emission is due to complex, highly nonlinear processes which are not entirely understood. These considerations make thermal SZ distortions an ideal method for selecting galaxy cluster samples of cosmological interest. The overall number density of clusters in the universe at a given redshift and mass limit is a sensitive function of cosmological expansion history \\cite{eke96,bar96}, making this a potentially powerful probe of dark energy \\cite{wel02}; combining an SZ catalog with optical cluster redshifts provides a probe of cluster growth through time. The difficulty in this approach is that cluster number counts also depend sensitively on systematic errors, particularly in understanding the cluster mass selection function \\cite{maj03,fra05}. Over the coming few years, significant numbers of clusters with SZ, X-ray, and optical observations, including weak-lensing mass estimates, will reveal whether clusters are regular enough to understand systematics in their selection functions and exploit them as probes of cosmology.  Outside of clusters, small diffuse SZ signals may also come from from hot gas in numerous individual galaxies heated by star formation or active galactic nuclei \\cite{pla02} or from the first generation of stars \\cite{oh03}. These signals will be at amplitudes of a few $\\mu$K, and the extent to which they can be exploited for useful scientific information, as opposed to serving only as a noise source, is not yet clear.  \\subsection{Kinematic Sunyaev-Zeldovich Effect} The kinematic Sunyaev-Zeldovich Effect also arises from photon-electron scattering; it is essentially a Doppler distortion coming from bulk velocity of the electrons along the line of sight, and is proportional to the cluster ionized gas mass times its radial peculiar velocity \\cite{sun80}. For large galaxy clusters with characteristic radial velocities of a few hundred km/s, the kSZ blackbody distortion will be between a few and 10 $\\mu$K. Simulated clusters show that the internal gas motions in a cluster give an intrinsic scatter of around 50 km/s to cluster velocities inferred from kSZ observations \\cite{nag03}. In principle, galaxy cluster radial peculiar velocities can be measured directly from kSZ observations, if cluster velocities and gas masses can be disentangled using thermal SZ or X-ray observations \\cite{seh05}. The cluster velocity field then provides a direct and clean probe of the cosmic velocity field and an independent test of the hierarchical structure formation picture. Such velocity catalogs would be a significant improvement over current velocity catalogs based on galaxy velocities, which are inferred from their redshifts combined with a distance estimate.  \\subsection{Other Sky Signals} These sky signals -- thermal and kinematic SZ effect, gravitational lensing, and primary microwave background blackbody fluctuations, can be separated through their frequency and spatial dependences. Detecting them will be complicated, however, by other astrophysical ``noise\" sources (which, however, are interesting in their own right!). Galactic dust can contribute substantial signal at these wavelengths; to minimize this signal, the cosmologically useful portions of the sky are likely to be well away from the galactic plane. Infrared point sources, which are thought to be high-redshift galaxies with significant dust emission from massive star formation, will likely give a confusion noise limit in the range of a few $\\mu$K at ACT frequencies \\cite{kno04,whi04}; removing the brightest infrared point sources through complementary observations at higher redshift will certainly be helpful. Likewise, bright radio point sources, generally emission from active galactic nuclei, can contribute, though these will probably have signficantly less overall signal than the infrared sources, judging from naive extrapolations of microwave and SCUBA observations \\cite{kno04}, unless the radio sources are strongly correlated with clusters . ",
        "conclusions": ""
    },
    "0608/astro-ph0608580_arXiv.txt": {
        "abstract": "A special purpose N-body simulation has been built to understand the tidal heating of the smallest dark matter substructures ($10^{-6}\\msun$ and 0.01pc) from the grainy potential of the Milky Way due to individual stars in the disk and the bulge.  To test the method we first run simulations of single encounters of microhalos with an isolated star, and compare with analytical predictions of the dark particle bound fraction as a function of impact parameter. We then follow the orbits of a set of microhalos in a realistic flattened Milky Way potential.  We concentrate on (detectable) microhalos passing near the Sun with a range of pericenter and apocenter. Stellar perturbers  near the orbital path of a microhalo would exert stochachstic impulses, which we apply in a Monte Carlo fashion according to the Besancon model for the distribution of stars of different masses and ages in our Galaxy. Also incorporated are the usual pericenter tidal heating and disk-shocking heating.  We give a detailed diagnosis of typical microhalos and find microhalos with internal tangential anisotropy are slightly more robust than the ones with radial anisotropy. In addition, the dark particles generally go through of a random walk in velocity space and diffuse out of the microhalos. We show that the typical destruction time scales are strongly correlated with the stellar density averaged along a microhalo's orbit over the age of the stellar disk. We also present the morphology of a microhalo at several epochs which may hold the key to dark matter detections. ",
        "introduction": "Cold Dark Matter (CDM) has enjoyed a great deal of success in explaining the fluctuations of the cosmic microwave background and large scale structure in general (Spergel et al. 2006). However, its predictions on small scales have not been fully understood or confirmed by observations yet.  This is partly limited by our computing power in resolving the smallest dark matter structures and detailed gas and stellar dynamics in structure formation.  Nevertheless in a pioneering work Diemand, Moore \\& Stadel (2005) were able to resolve a galaxy-size dark halo down to the free-streaming scale.  The smallest structures (called microhalos) have a minimum mass ($m_0\\sim 10^{-6}M_{\\sun}$), which were followed in simulations from redshift $z=64$ to $z=26$.  As much as 0.1 percent of the dark matter in the universe are in these substructures.  At later redshift larger substructures collapse out of perturbations on larger scales, and hierarchical halo formation (Lacey \\& Cole 1994) predicts that the microhalos also fall into these larger structures and are expected to survive inside given their high internal density.  Extrapolations show the Milky Way may contain around $10^{15}$ such microhalos with the nearest one being on the average about 0.1 pc distant from the Earth, and with typical size 0.01 pc.  The numerical simulations of Diemand et al (2005, 2006) also confirm the analytical predictions of Berezinsky, Dokuchaev \\& Eroshenko (2005).   There has been much intrigue over tidal mass-loss from satellite galaxies, for instance the tidal streams of the Sagittarius dwarf galaxy (e.g. Savage, Freese \\& Gondolo 2006) and the Magellanic clouds (e.g. Johnston, Spergel \\& Haydn 2002). However, details of the mass-loss from the smallest scale dark matter halos are still sketchy and although some work has been done to develop semi-analytical models of microhalo destruction (Zhao et al. 2006, Berezinski et al. 2005, Green \\& Goodwin 2006), virtually no rigorous modelling using N-body codes has been published. It has been noted (Green \\& Goodwin 2006) that making analyses of disruption rates from seperate consideration of the tides, stellar encounters and disk shocking is insufficient and one consistent approach is needed. Here we performed N-body simulations in a respectable star distribution (Robin, Reyle \\& Picaud 2003) with microhalos on orbits in a disk type galaxy potential with dark halo, disk and bulge components, adding contributions from star encounters in each time step over the age of the stellar disk.  The paper is laid out as follows; in \\S \\ref{sec:incon} we give an analytical prescription of a typical microhalo using the NFW profile (Navarro et al. 1997). We have two sets of initial conditions for the internal particles of the microhalo. The first is for randomly inclined circular orbits and the second is with a radially anisotropic dispersion.  Then we set up the geometry of single encounters with stars and compare the impulse approximation with numerical simulations using the N-body tree code of Vine \\& Sigurdsson (1998).  Using N-body simulations we calculated, for four typical stellar masses (0.1, 0.3, 0.6 and 1.0 $M_{\\sun}$), (i) the likely safe impact parameter between stars and microhalos where no mass should be stripped, (ii) the critical impact parameter below which the microhalo will be destroyed (i.e. 95\\% of particles given escape velocity). (iii) We show for single encounters the difference between instantaneous bound fraction of microhalo mass and the fraction remaining bound after particle escape is allowed and the potential shallows. Finally (iv) we present the bound fraction of microhalo mass as a function of impact parameter for a single encounter, comparing the microhalos with radially anisotropic and randomly oriented circular initial conditions.  In \\S \\ref{sec:realsim} we present a procedure to use the Besanson model (Robin et al. 2003) of the Milky Way to simulate the encounters of microhalos with stars on typical orbits in the Milky Way in a three component (bulge, disk, halo) potential over a several Gyr orbital timescale until the microhalo is destroyed. Then in \\S \\ref{sec:res} we use this model to produce some interesting results. (i) That the survival time of microhalos and the averaged stellar density along the orbit can be related by a power law and (ii) That the survival time is keenly linked to the orbital energy and angular momentum. (iii) Particles escape via diffusion during disk crossings and (iv) the microstreams created from destroyed microhalos remain filamentary for many Gyr. We conclude and discuss our results in \\S \\ref{sec:conc}. ",
        "conclusions": "\\protect\\label{sec:conc} Microhalos of $10^{-6}M_{\\sun}$ are likely the first substructures formed in a CDM universe. Those microhalos in dense regions of galaxies would experience strong tidal stripping. Here we presented N-body simulations of microhalos in single encounters with stars as well as in a series of encounters in a realistic Milky Way.  The clear trend is that for orbits close to the solar neighbourhood the survival time falls off as a power law of the average stellar density and the only likely surviving microhalos are ones on orbits with high energy and angular momentum; also the TA microhalos have a slightly better chance of survival than the RA microhalos. \\\\ \\\\ Nevertheless, some cautionary points: \\\\ \\\\ 1. We consider only Sun-passing orbits. Those closer to the bulge are likely more severely stripped and more distant ones are less interesting for detection. \\\\ \\\\ 2. Our microhalos all have the same total mass of $10^{-6}M_{\\sun}$ and fixed initial tidal radius (0.01pc).  In realisty slightly more massive microhalos may endure stellar encounters to greater effect and exist in higher numbers today.  Also the initial tidal radius should vary with the pericenter of the microhalo. \\\\ 3. The stars in our Galaxy are not formed all at exactly 4 Gyrs ago.  The granularity and the flattening of the galactic potential presumably increased gradually over a Hubble time.   However, our method shows how the impulsive encounters with stars alongside disk shocking and tidal stripping poses a serious threat to the survival of primordial microhalos. The statistical nature of survival that we have demonstrated could be used or at least elaborated upon to predict statistically the phase space distributions of microhalos and morphology in order to make more accurate estimates of annihilation and direct detection signals."
    },
    "0608/astro-ph0608063_arXiv.txt": {
        "abstract": "{In the framework of the study of the X-ray and optical emission in supernova remnants we focus on an isolated X-ray knot in the northern rim of the Vela SNR (Vela FilD), whose X-ray emission has been studied and discussed in Paper I.} {We aim at understanding the physical origin of the X-ray and optical emission in FilD, at understanding the role of the different physical processes at work, and at obtaining a key for the interpretation of future X-ray observations of SNRs.} {To this end we have pursued an accurate ``forward'' modeling of the interaction of the Vela SNR shock with an ISM cloud. We perform hydrodynamic simulations and we directly compare the observables synthesized from the simulations with the data.} {We explore four different model setups, choosing the values of the physical parameters on the basis of our preliminary analysis of the X-ray data. We synthesize X-ray emission maps and spectra filtered through the \\emph{XMM-Newton} EPIC-MOS instrumental response. The impact of a shock front at 6 million Kelvin on an elliptical cloud 30 times denser than the ambient medium describes well the shock-cloud interaction processes in the Vela FilD region in terms of spectral properties and morphology of the X-ray and optical emission.} {The bulk of the X-ray emission in the FilD knot originates in the cloud material heated by the transmitted shock front, but significant X-ray emission is also associated to the cloud material, which evaporates, as an effect of thermal conduction, in the intercloud medium. The physical origin of the FilD optical emission is associated to thermal instabilities. In the FilD knot the X-ray emission associated to the reflected shock front is negligible.} ",
        "introduction": "\\label{Introduction}  The study of supernova remnants (SNRs) allows us to probe the physical and chemical properties of the interstellar medium (ISM). Moreover, SNRs drive the mass and energy exchange in the ISM, thus playing a fundamental role for its evolution.  The X-ray emission of SNRs is characterized by a clumpy morphology that has been interpreted as a result of the interaction between the blast wave expanding shock and inhomogeneities of the ambient medium. The impact of the shock front with ISM clouds influences the direction and the speed of the SNR expansion and determines the heating and the ionization of the clouds. Several authors (see, for example, \\citealt{dsb97}, \\citealt{bms00}, and \\citealt{pfr02}) have shown how the study of X-ray and optical observations of SNRs provides important information about the shock-cloud interaction process and about the physical and chemical conditions of the post-shock medium (temperature, chemical composition, morphology of the clouds, etc.). However, the details of the heating processes and the physical scenario which determines the observed emission are still not well understood.  In this framework, shock-cloud interaction has been extensively studied both analytically (e. g. \\citealt{mc75}, \\citealt{hb83}) and numerically (e. g. \\citealt{bw90}, \\citealt{sn92}, \\citealt{kmc94}, \\citealt{xs95}, and \\citealt{fag05}) and, recently, the combined effects of thermal conduction and radiative cooling have been investigated for the first time by \\citet{opr05}. This work has shown that, when radiative losses dominates, the shocked cloud is fragmented into cold and dense cloudlets, while thermal conduction determines the evaporation of the clouds and suppresses hydrodynamic instabilities. Despite these large theoretical efforts, however, no detailed comparison with the observations has yet been provided. In fact, although comparisons with the observations are present, for example, in  \\cite{hs84} and in \\cite{wl91}, these models are very simplified and do not include at the same time all the relevant physical effects. Recently, an observation-oriented numerical model has been presented by \\citet{pf05} to describe the shock-cloud interaction in the southwestern region of the Cygnus Loop. However, no observable quantities are produced from this model and only a simple qualitative comparison between the model density plot and the overall morphology of the H$\\alpha$ emission is presented.  \\citet{mbm05} (hereafter Paper I) analyzed an \\emph{XMM-Newton} EPIC observation of a small isolated knot (named FilD) in the northern rim of the Vela SNR. The analysis revealed a link between optical and X-ray emission and provided information about density, temperature, extension along the line of sight, and chemical composition of the FilD cloud. In particular, the spectral analysis has shown that the X-ray emission of FilD is described well by an optically-thin plasma with two thermal components at $\\sim 1\\times 10^{6}$ K and $\\sim 3\\times 10^{6}$ K, respectively. The inhomogeneity of the surface brightness in the FilD region is determined by different volume distributions along the line of sight of the two components.  Here we present a detailed hydrodynamic model of the shock-cloud interaction in the Vela FilD. Our approach consists in synthesizing observable quantities from the model (for example, maps of X-ray surface brightness and spectra), comparing them directly with the \\emph{XMM-Newton} EPIC data and with optical observations of the FilD region. We also perform on the synthesized products the same analysis we carried on for the EPIC data.  The aims of this work are: i) the proper interpretation of the X-ray observation (i. e. to understand the physical origin of the two spectral components, to ascertain the evolutionary stage of the interaction, to obtain information about the density and temperature profiles of the emitting plasma); ii) to understand the role of the different physical processes (transmitted and reflected shocks, thermal conduction, radiative cooling) in the evolution of the system; iii) to obtain a key for the interpretation of future X-ray observations of middle-aged SNRs.  The paper is organized as follows: in Sect. \\ref{FilD} we review the main results of Paper I that give us constraints for our simulations; the model equations and setups are presented in Sect. \\ref{The Model} and in Sect. \\ref{The model setup}; our results are presented in Sect. \\ref{Results} and discussed in Sect. \\ref{Discussion}. Finally, our conclusions are drawn in Sect. \\ref{Summary and conclusions}. ",
        "conclusions": "\\label{Summary and conclusions} Our analysis is focussed on the physical description of a bright isolated X-ray knot in the northern rim of the Vela SNR (Vela FilD), whose X-ray emission has been studied and discussed in Paper I. We modelled the hydrodynamics of the interaction between a SNR blast wave shock and an isolated cloud, taking into account thermal conduction and radiative cooling effects. We explored four different model setups, choosing the values of the physical parameters on the basis of our preliminary analysis of the \\emph{XMM-Newton} X-ray data. Since we synthesized the X-ray emission filtered through the \\emph{XMM-Newton} EPIC-MOS instrumental response from the model, we can compare the model results directly to the data and derive detailed and quantitative information and deep insight. Our study has shown that: \\begin{itemize} \\item The dynamics of the shock-cloud interaction and the corresponding X-ray and optical emission are significantly influenced by the pre-shock morphology of the ISM cloud. In particular, if the Vela shock front impacts on a spherical cloud, both transmitted and reflected shocks contribute to the X-ray emission, while in the case of an ellipsoidal cloud, the reflected shock contribution is negligible. \\item An ellipsoidal cloud 30 times denser than the intercloud medium, with the semimajor axis parallel to the shock velocity describes well the \\emph{XMM-Newton} observation of the Vela FilD knot in terms of X-ray emission spectral properties and morphology, and in terms of spatial relationship between optical and X-ray emission. \\item The peculiar orientation of the FilD optical filament (which is difficult to explain according to the common scenario which associates optical emission in SNRs with slow shocks travelling through dense clouds) is naturally explained by this setup as a result of thermal instabilities. \\item Our model allows us to understand the physical origin of the two thermal components we used in Paper I to describe the FilD observed spectra. The cooler component originates in the cloud material heated by the transmitted shock front, while the hotter component is the result of thermal conduction between the cloud and the hotter shocked intercloud medium. The heating of the cloud produces its evaporation and the formation of a hot diffuse halo. \\item We highlighted spectral and morphological characteristics of the emission from transmitted and reflected shocks, thus obtaining a key for the interpretation of future shock-cloud observations. \\end{itemize}  Future studies should also address other open issues, like the presence of a single optical filament observed in FilD and the effects of the pre-shock symmetry of the cloud on the morphology of the optical emission and the on the surface brightness in X-rays. Moreover, we also aim at describing the shock-cloud interaction in RegNE, which is an other bright X-ray knot, hotter than FilD, discussed in Paper I."
    },
    "0608/astro-ph0608255_arXiv.txt": {
        "abstract": "{} {We present a search for low-mass brown dwarfs in the Pleiades open cluster. The identification of Pleiades members fainter and cooler than those currently known allows us to constrain evolutionary models for L dwarfs and to extend the study of the cluster mass function to lower masses.} {We conducted a 1.8~deg$^2$ near-infrared $J$-band survey at the 3.5~m Calar Alto Telescope, with completeness $J_{cpl}\\sim$~19.0. The detected sources were correlated with those of previously available optical $I$-band images ($I_{cpl}\\sim$~22). Using a $J$~versus~$I-J$ colour--magnitude diagram, we identified 18 faint red L-type candidates, with magnitudes $17.4<J<$~19.7 and colours $I-J>$~3.2. If Pleiades members, their masses would span $\\sim$0.040--0.020~$M_{\\sun}$. We performed follow-up $HK_{\\rm s}$-band imaging to further confirm their cluster membership by photometry and proper motion.} {Out of 11 $IJ$ candidates with proper motion measurements, we find six cluster members, two non-members and three whose membership is uncertain and depends on the intrinsic velocity dispersion of Pleiades brown dwarfs. This dispersion ($>$4~mas~yr$^{-1}$) is at least four times that of cluster stars with masses $\\ga$1~$M_{\\sun}$. Five of the seven other $IJ$ candidates are discarded because their $J-K_{\\rm s}$ colours are bluer than those of confirmed members. Our least massive proper motion members are \\object{BRB~28} and 29 ($\\sim$25~$M_{Jup}$). The $J$~versus~$I-J$ sequence of the L-type candidates at $J>18$ is not as red as theoretical models predict; it rather follows the field L-dwarf sequence translated to the cluster distance. This sequence overlapping, also observed in the $J$~versus~$J-H$ and $J-K$ diagrams, suggests that Pleiades and field L dwarfs may have similar spectral energy distributions and luminosities, and thus possibly similar radii. Also, we find $\\alpha=$~$0.5\\pm0.2$ for a power-law approximation $dN/dM \\propto M^{-\\alpha}$ of the survey mass spectrum in the mass range 0.5--0.026~$M_{\\sun}$. This value is similar to that of much younger clusters, indicating no significant differential evaporation of low-mass Pleiades members relative to more massive ones.} {} ",
        "introduction": "As a result of progressive cooling, brown dwarfs are expected to evolve from late M~spectral type at very early stages (age $\\la$~10~Myr) to L-type ($T_{\\rm eff}$ in the range $\\sim$2\\,400 to $\\sim$1\\,400~K; \\citealt{dahn2002}), and at sufficiently old ages (a few Gyr), to T-type and beyond ($T_{\\rm eff}$ $\\la$~1\\,400~K). The mass-spectral type and mass-luminosity relationships at different ages remain to be derived from observations of cool dwarfs, which also permit the calibration of the theory of substellar evolution (e.g. \\citealt{chabrier2000a}). In comparison to the spectra of M dwarfs, L dwarfs are characterized in the optical by the weakening of the metal-oxide TiO and VO bands due to the condensation of the metals Ti and V in dust grains that may significantly affect the atmospheric structure and the emergent spectrum (e.g. the DUSTY models from \\citealt{chabrier2000b}). Bands of metal-hydrides CrH, FeH and CaH and bands of water become more intense especially in the near-infrared (IR), whereas the neutral alkali metals Na, K, Rb, Cs and Li become stronger in the optical \\citep{chabrier2000a,basri2000}. The emergence in the near-IR of CH$_{4}$ bands characterizes the beginning of the T dwarf sequence. Field L dwarfs have $I-J$ and $J-K_{\\rm s}$ colours redder than $\\sim$3.3 and $\\sim$1.3~mag, respectively, increasing from early to later spectral subclasses \\citep{martin1999}. For the L4--L8 spectral subclasses the near-IR $J-H$ and $J-K_{\\rm s}$ colours appear to saturate at $\\sim$1.2 and $\\sim$2.0~mag \\citep{kirkpatrick2000,burgasser2002}, respectively, whereas for T dwarfs these colours become bluer, with values decreasing to below zero \\citep{knapp2004}. From L0 to T8 spectral types the $I-J$ colour is increasing to $\\sim$5.8~mag \\citep{dahn2002}. Associated with the other colours it can give a rough indication of the spectral type.   The \\object{Pleiades} open cluster ($\\sim$120~Myr, $\\sim$130~pc) served for decades as a reference stellar laboratory where models were contrasted with observations. Since the discovery of Pleiades brown dwarfs \\citep{rebolo1995,rebolo1996}, subsequent studies have provided numerous fainter substellar candidates with spectral types down to late M (\\citealt{zapateroosorio1997a,zapateroosorio1997b,bouvier1998}, hereafter\\defcitealias{bouvier1998}{B98}B98; \\citealt{festin1998a,zapateroosorio1999Pl}), from which a fraction is confirmed by lithium detection \\citep{stauffer1998a,martin2000} or by proper motion \\citep[ hereafter M01]{moraux2001}. Very large area proper motion surveys \\citep{hambly1999,deacon2004} have obtained on the other hand a significant census of the stellar-substellar boundary population of the cluster.\\defcitealias{moraux2001}{M01} The coolest object for which a spectral type has been obtained is the L0 brown dwarf candidate \\object{Roque~25} \\citep{martin1998b}, with an estimated theoretical mass of $\\sim$0.035~M$_{\\sun}$. A dozen fainter objects than \\object{Roque~25} have been identified in deep surveys (\\citealt{festin1998a,bejar2000,dobbie2002,nagashima2003}; \\citealt[ hereafter M03] {moraux2003}; \\citealt{schwartz2005}), and some were confirmed by proper motion \\citep{bouy2006}.\\defcitealias{moraux2003}{M03}  The mass spectrum $dN/dM=f(M)$ of the Pleiades cluster, where $dN$ stands for the number of objects in the mass range $dM$, is best fitted by a lognormal function (\\citetalias{moraux2003}; \\citealt{deacon2004}), and its approximation by a power law $M^{-\\alpha}$ in the low mass range 0.6--0.03~$M_{\\sun}$ provides a spectral index $\\alpha$ $\\sim$~0.6--0.8 (\\citealt{dobbie2002}; \\citetalias{moraux2003}). Similar values are found in other much younger open clusters, as for example \\object{$\\sigma$ Orionis} ($\\sim$2--4~Myr), where $\\alpha=$~$0.8\\pm0.4$ for $0.2<$~$M(M_{\\sun})<$~$0.013$ \\citep{bejar2001}, and \\object{IC~348} ($\\sim$3~Myr), where $\\alpha=$~$0.7\\pm0.2$ for $0.5<M(M_{\\sun})<0.035$ \\citep{tej2002}. It is important to determine the behaviour of the $\\alpha$ index in the lower mass domain because it can help to discriminate among brown dwarf formation mechanisms. For instance, for the formation of gravitationally unstable cores by turbulent fragmentation \\citep{padoan2004}, the mass spectrum is predicted to be log-normal, with a peak that appears at smaller masses when the sonic Mach number and the mean density of the cores are greater.  In this paper we present a 1.8~deg$^2$ deep $J$-band survey of the Pleiades open cluster aimed to find lower-mass brown dwarfs ($<$0.035~$M_{\\sun}$), by comparison with a new analysis of $I$-band data obtained by \\citetalias{bouvier1998}. Our goal is to identify L dwarf cluster candidates and study the location of the L dwarf sequence in the colour--magnitude diagram. Preliminary results have been reported in \\citet{bihain2005}. Because the cluster has a well known age, the properties of Pleiades L-type brown dwarfs could provide useful constraints to models predicting the evolution of substellar objects. Additionally, we can estimate the cluster mass spectrum of the surveyed area at lower masses. ",
        "conclusions": "\\label{mf_section}  As a result of an $IJ$-band survey in the Pleiades, we have found 18 brown dwarf candidates, most likely of L-type. They have 17.4~$<J<$~19.7, $I-J>$~3.2 and theoretical masses $\\sim$0.040--0.020~$M_{\\sun}$. The near-IR follow-up for proper motion of 11 candidates permits us to confirm six as cluster members. Another three remain as possible members. Two are clear non-members. The near-IR photometric follow-up of the seven other candidates indicates that at least five are probable contaminants, because their $J-K_{\\rm s}$ colours are bluer than those of the cluster proper motion members. These latter determine the L brown dwarf photometric sequence of the Pleiades. The sequence in the $J$~versus~$I-J$ colour--magnitude diagram at $J>18$ is bluer than that predicted by the DUSTY atmospheres model, whereas it agrees with the predictions in the $J$~versus~$J-H$ and $J-K$ diagrams. Moreover the sequence overlaps that of the field L~dwarfs shifted to the cluster distance in the three colour--magnitude diagrams, suggesting that the Pleiades and field L~dwarfs may have comparable spectral energy distributions and luminosities, and thus possibly similar radii.  We find evidence for an intrinsic velocity dispersion of Pleiades brown dwarfs at least four times that of Pleiades stars more massive than the Sun. The estimated value $>$4~mas~yr$^{-1}$ appears consistent with the kinematical expectations for brown dwarfs in a nearly relaxed cluster.  Correcting for contaminants and comparing with the proper motion objects from \\citet{deacon2004} in the same fields, we obtain a substellar mass spectrum for the surveyed area of the cluster. For a power law $dN/dM \\propto$~$M^{-\\alpha}$ fit in the mass range 0.5--0.026~$M_{\\sun}$, we find a spectral index $\\alpha=$~0.5~$\\pm$~0.2 and in agreement with previous results. The slope is similar to that of much younger open clusters, suggesting the absence of significant differential evaporation of the low-mass brown dwarfs in the Pleiades. This supports deeper searches to detect lower mass free-floating objects in the cluster."
    },
    "0608/astro-ph0608125_arXiv.txt": {
        "abstract": "Direct-summation $N$-body algorithms compute the gravitational interaction between stars in an exact way and have a computational complexity of $\\mathcal{O}(N^2)$.  Performance can be greatly enhanced via the use of special-purpose accelerator boards like the GRAPE-6A. However the memory of the GRAPE boards is limited. Here, we present a performance analysis of direct $N$-body codes on two parallel supercomputers that incorporate special-purpose boards, allowing as many as four million particles to be integrated. Both computers employ high-speed, Infiniband interconnects to minimize communication overhead, which can otherwise become significant due to the small number of ``active'' particles at each time step.  We find that the computation time scales well with processor number; for $2\\times 10^6$ particles, efficiencies greater than 50\\% and speeds in excess of $2\\TFlops$ are reached. ",
        "introduction": "Numerical algorithms for solving the gravitational $N$-body problem \\citep{aarseth-03} have evolved along two basic lines in recent years. Direct-summation codes compute the complete set of $N^2$ interparticle forces at each time step; these codes are designed for systems in which the finite-$N$ graininess of the potential is important or in which binary- or multiple-star systems form, and until recently, were limited by their $\\mathcal{O}(N^2)$ scaling to moderate ($N\\lap 10^5$) particle numbers.  The best-known examples are the {\\tt NBODY} series of codes introduced by \\cite{aarseth-99} and the {\\tt Starlab} environment developed by McMillan, Hut, and collaborators (e.g. \\cite{STARLAB}).  A second class of $N$-body algorithms replace the direct summation of forces from distant particles by an approximation scheme.  Examples are the Barnes-Hut tree code \\citep{bh-86}, which reduces the number of force calculations by subdividing particles into an oct-tree, and fast multipole algorithms which represent the large-scale potential via a truncated basis-set expansion \\citep{vavg-77,gr-87}, or on a grid \\citep{mp-68,ee-81}.  These algorithms have a milder, $\\mathcal{O}(N\\log N)$ or even $\\mathcal{O}(N)$ scaling for the force calculations and can handle much larger particle numbers, although their accuracy can be substantially lower than that of the direct-summation codes \\citep{spurzem-99}. The efficiency of both sorts of algorithm can be considerably increased by the use of individual time steps for advancing particle positions \\citep{aarseth-03}.  The $N$-body problem is particularly challenging in the case of dense stellar systems like galactic nuclei, which may contain single or multiple, supermassive black holes in addition to stars \\citep{FF05,living}. Particle advancement must be very accurate for trajectories that pass near the black hole(s). In addition, galactic nuclei are often ``collisional'' in the sense that gravitational encounters (i.e. scattering) can redistribute energy between stars on time scales less than the age of the universe \\citep{merritt-06}. Simulating the evolution of a collisional nucleus is probably not feasible using tree or grid algorithms due to their limited accuracies. While alternative, highly-efficient algorithms based on the Fokker-Planck or fluid formalisms have been widely used to model galactic nuclei \\citep{ls-91,fb-01}, these methods can not deal with systems that are far from dynamical equilibrium or that fail to respect spatial symmetries.  A natural way to increase both the speed and particle number in an $N$-body simulation is to parallelize \\citep{dubinski-96,pc-97}. Parallelization on general-purpose supercomputers is difficult, however, because the calculation cost is often dominated by a small number of particles in a single dense region, e.g. the nucleus of a galaxy. Communication latency becomes the bottleneck: the time to communicate particle positions between processors can exceed the time spent computing the forces.  The best such schemes use systolic algorithms (in which the particles are rhythmically passed around a ring of processors) coupled with non-blocking communication between the processors to reduce the latency \\citep{makino-02,dorband-03}.  A major breakthrough in direct-summation $N$-body simulations came in the late 1990s with the development of the GRAPE series of special-purpose computers \\citep{grapebook}, which achieve spectacular speedups by implementing the entire force calculation in hardware and placing many force pipelines on a single chip.  The GRAPE-6, in its standard implementation (32 chips, 192 pipelines), can achieve sustained speeds of about 1 Tflops at a cost of just $\\sim\\$50$K. In a standard setup, the GRAPE-6 is attached to a single host workstation, in much the same way that a floating-point or graphics accelerator card is used.  Advancement of particle positions [$\\mathcal{O}(N)$] is carried out on the host computer, while interparticle forces [$\\mathcal{O}(N^2)$] are computed on the GRAPE. More recently, ``mini-GRAPEs'' (GRAPE-6A) \\citep{FMK05} have become available, which are designed to be incorporated into the nodes of a parallel computer. The mini-GRAPEs place four processor chips on a single PCI card and delivering a theoretical peak performance of $\\sim 131$ Gflops for systems of up to 128k particles, at a cost of $\\sim\\$6$K. By incorporating mini-GRAPEs into a cluster, both large ($\\gap 10^6$) particle numbers and high ($\\gap 1\\TFlops$) speeds can in principle be achieved.  In this paper, we describe the performance of direct-summation $N$-body algorithms on two computer clusters that incorporate GRAPE hardware. \\S 2 describes the hardware implementations. The parallel $N$-body code and its implementation on the GRAPEs is described in \\S 3. \\S 4 presents the results of performance tests using realistic galaxy models, and \\S 5 describes a theoretical performance model that reproduces the observed performance and which can be used to predict the performance of similar codes on different clusters. \\S 6 summarizes our results and discusses directions for future work. ",
        "conclusions": "The performance of a direct-summation, parallel $N$-body code has been evaluated on two new computer clusters incorporating GRAPE special-purpose accelerator boards. Parallelization was carried out using a hybrid method incorporating aspects of both systolic (``ring'') and broadcast (``copy'') algorithms, in order to make the most efficient use of the GRAPE pipelines. The algorithm exhibits asymptotically an $\\mathcal{O}(Np)$ scaling in communication complexity and an $\\mathcal{O}(N^2/p)$ scaling in computation time. Benchmark simulations were carried out using a set of galaxy models having realistically high degrees of central concentration. Using one million particles and 32 nodes, the clusters achieved a sustained performance of 50\\% -- 100\\% of the theoretical peak ($\\sim 4\\TFlops$). When run on general-purpose parallel computers, $N$-body codes like ours typically only achieve a few per cent of peak performance; hence, our special-purpose computer clusters are competitive with the fastest computers in the world when applied to the gravitational $N$-body problem.  We presented a simple model that predicts the performance of the $N$-body code as a function of processor number $p$, particle number $N$, and hardware constants. The model reproduces the observed performance very well, and can be used to predict the performance of the code on clusters with different hardware characteristics. The model supports our finding that the performance depends critically on having very fast and low-latency communication hardware.  It is convenient to discuss the performance of the parallel $N$-body code in terms of a working point $P$ defined by a pair of values $P=(p,N)$, where $P$ is given by the condition that the time required for communication and computation be approximately equal. When using all 32 of the processors on our clusters, this point is reached for $N\\approx 10^6$. For larger $N$, the system operates at or near optimal speedup and efficiency (see Figs. 6 and 7). At a given $p$, the linear scaling of total computation time with $N$ at low $N$ (communication-dominated) changes to the asymptotic, $N^2$ scaling (computation dominated) roughly at $P$. Increasing $N$ beyond its value at $P$ would still achieve near-optimal speedups, but there would be room to increase $p$ without sacrificing efficiency (if more nodes were available). Increasing $p$ beyond $P$ at fixed $N$ is not useful because communication overhead will start to dominate, leaving the GRAPEs idle for part of the time. The design goal of our clusters was to simulate $\\sim 10^6$ particles at high efficiencies, and we have shown by our performance tests that this goal has been achieved.  \\begin{figure}[ht] \\begin{center} \\includegraphics[width=10cm,angle=-90.]{harfst.fig15.ps} \\end{center} \\caption{This figure demonstrates what particle numbers are needed to achieve the separation of time scales discussed in the text, when simulating nuclei containing a central sink, e.g. a single or binary BH. Vertical axis is the fraction of the flux into the ``sink'' originating from orbits that are in the empty-loss-cone regime; in real galaxies, this ratio is close to unity. Curves are labelled by $r_{LC}/r_h$, where $r_{LC}$ is the linear size of the central ``sink'' and $r_h$ is the BH influence radius. For a binary SBH, $0.1\\le r_{LC}/r_h\\le 10^{-3}$. } \\label{fig:empty} \\end{figure}  How large an $N$ is required to accurately represent the evolution of a dense stellar system like a galactic nucleus? A minimum condition is that $N$ be large enough to enforce a strict separation of time scales between the period $T_{orb}$ of an orbit and the relaxation time $T_{ref}$; the latter is the time for two-body gravitational scattering to randomize velocities \\citep{spitzer-87}. In real galaxies, $T_{rel}\\gg T_{orb}$, i.e. the integrity of stellar orbits is maintained for many orbital periods. In an $N$-body simulation, one has \\begin{equation} T_{rel}\\approx {0.1 N\\over \\ln N} T_{orb} \\label{eq:tr} \\end{equation} \\citep{aarseth-03}, and the separation of time scales can be achieved with modest particle numbers, $N\\gap 10^3$. Relaxation-driven processes like core collapse can be simulated using $N$-body codes and the numerical evolution rates scaled to those in real galaxies using an equation like (\\ref{eq:tr}) \\citep{SA-96,SMK-05}.  In simulations of galactic nuclei, a much stronger separation of time scales must be achieved if the results are to be scaled to real systems, implying larger values of $N$. Galactic nuclei contain ``sinks,'' regions near the center where stars are lost or captured. For instance, the supermassive black hole (BH) at the center of the Milky Way galaxy disrupts stars that pass within a distance $r_{tidal}\\approx 10^{-6}$ pc of it. Another example is a galaxy containing a binary supermassive BH: stars that pass within a distance $\\sim a$ of either BH are ejected by the gravitational slingshot, where $10^{-3}\\ {\\rm pc} \\lap a \\lap 10^{-1}$ pc is the semi-major axis of the binary. In this case, the radius of the ``sink'' is $\\sim a \\gg r_{tidal}$. Single or binary supermassive BHs are believed to be generic components of galaxies \\citep{FF05,living}, and the structure and evolution of nuclei may be shown to depend critically on the rate at which stars are lost to the central sink \\citep{merritt-06}.  In real galaxies, relaxation times are long enough that most stars would diffuse gradually -- i.e., over many orbital periods -- onto so-called loss-cone orbits that intersect the capture sphere of radius $r_{tidal}$ or $a$. This means that the loss cone is essentially empty, since the removal time is much shorter than the diffusion time. In order to reproduce a diffusive loss-cone repopulation in an $N$-body simulation, the relaxation time must be large enough that \\begin{equation} T_{rel} > \\left({r\\over r_{LC}}\\right) T_{orb} \\gg T_{orb} \\label{eq:loss} \\end{equation} \\citep{LS-77}. Here, $r_{LC}$ is the radius of the disruption ($\\sim r_{tidal}$) or ejection ($\\sim a$) sphere at the center of the galaxy, and $r$ is the typical size of a stellar orbit. Equation~(\\ref{eq:loss}) reflects the fact that scattering into a loss-cone orbit occurs in much less than one relaxation time, hence maintaining an empty loss cone requires that relaxation times be much longer than orbital periods. This is a stricter condition than $T_{rel}\\gg T_{orb}$ and implies a larger $N$.  Fig.~\\ref{fig:empty} presents a more careful analysis along these lines. Shown there is the fraction of the flux into the central sink originating from orbits that are in the empty loss cone regime, as a function of $N$ and $r_{LC}$; the latter is expressed in units of $r_h$, the radius of gravitational influence of the central (single or binary) BH of mass $M_\\bullet$. (The influence radius is the radius containing a mass in stars equal to twice $M_\\bullet$.) In the case of a binary supermassive BH, $r_{LC}/r_h\\approx a/r_h$ is roughly $10^{-1}$ at the time of binary formation, falling to $\\sim 10^{-3}$ when the two BHs are close enough to coalesce. Fig.~\\ref{fig:empty} suggests that particle numbers accessible to a single GRAPE-6 ($\\sim 0.25$M) can only reproduce the empty loss cone regime characteristic of real galaxies during the early phases of binary evolution. The values of $N$ achievable on gravitySimulator ($\\sim 4M$) allow the evolution of a binary to be followed to separations $\\sim 1/10$ as small. These estimates are consistent with the results of recent simulations which reproduce diffusion-limited evolution of massive binaries with $N\\approx 10^6$ \\citep{MF04,BMS05}. Simulating loss cone dynamics around the much smaller tidal disruption sphere of a single BH ($r_{LC}\\approx 10^{-6}r_h$) would require considerably larger $N$, well beyond the capabilities of existing or planned computers. However the dynamics in this regime could be qualitatively reproduced by adjusting the size of the capture sphere in an $N$-body simulation such that the bulk of the stars scattered into the BH are on orbits respecting the correct ratio of $T_{rel}$ to $T_{orb}$.  The $T_{rel}\\sim N$ scaling of Eq.~(\\ref{eq:tr}) implies an effectively $\\sim N^3$ scaling for calculations that extend over one relaxation time of the system. This scaling makes it very expensive to simulate the ``collisional'' evolution of stellar systems using the full, $4\\times 10^6$ particle number allowed by the combined GRAPE memories on the RIT and ARI clusters. One way to accelerate the computations without significant loss of accuracy would be to use the Ahmad-Cohen (AC) neighbor scheme \\citep{AC73}, as implemented in codes like NBODY5 or NBODY6 \\citep{aarseth-99}. In the AC scheme, the full forces are computed only every tenth timestep or so; in the smaller intervals, the forces from nearby particles (the ``irregular'' force) are updated using a high order scheme, while those from the more distant ones (the ``regular'' force) are extrapolated using a second-degree polynomial. A parallel implementation of NBODY6, including the AC scheme, exists, but only for general-purpose parallel computers \\citep{spurzem-99,khalisi-03}; the algorithm has not yet been adapted to systems with GRAPE hardware.  Greater speedups could also be achieved by increasing the number $p$ of processors, but only if communication costs are held low. One way to do this is via a variant of the \\cite{lippert-98a,lippert-98b} hyper-systolic (HS) algorithm. In its basic form, the HS algorithm reduces the number of data transfers by storing the shifted data on each node. The number of shifts, and hence the communication time, is reduced from $\\mathcal{O}(p)$ to $\\mathcal{O}(\\sqrt{p})$ and the memory requirements are increased by a similar factor. The HS speedup is smaller when only a subset of the full $N$ particles are advanced at every time step however \\citep{dorband-03}. \\citet{makino-02} has presented a direct $N$-body summation code optimized for a quadratic layout of processor ($p$ required to be a square number), which is in fact a simplified version of the hypersystolic algorithm proposed by \\citet{lippert-98b}. In Makino's case the asymptotic scaling of communication is $\\mathcal{O}(N/\\sqrt{p})$ (for calculation cost there is no difference to our code), which would allow to use much larger numbers of processors.  However, performance tests of a HS $N$-body algorithm on actual hardware have apparently not been carried out.  Implementation of the AC and/or HS schemes should permit effective use of the RIT and ARI clusters with the full particle numbers permitted by the combined GRAPE memories. Still larger $N$ could be attained by implementing a hybrid scheme, e.g. coupling a direct-summation algorithm with a tree \\citep{MA-93} or basis-function representation \\citep{HSS-02} for the distant particles, or a hierarchical generalization of the Ahmad-Cohen neighbor scheme.  \\ack SH and DM were supported by grants AST-0420920 and AST-0437519 from the NSF, grant NNG04GJ48G from NASA, and grant HST-AR-09519.01-A from STScI. AG and SPZ acknowledge support from the Netherlands Organization for Scientific Research (NWO, \\#635.000.001), the Royal Netherlands Academy of Arts and Sciences (KNAW) and the Netherlands Research School for Astronomy (NOVA). RS acknowledges support from the Volkswagen Foundation under grant No. I80 041-043, the Ministry of Science, Education and Arts of the state of Baden-W\\\"urttemberg, Germany, and the German Science Foundation under SFB 439 at the University of Heidelberg. PB acknowledges support from the German Science Foundation under SFB 439 and INTAS grant IA-03-59-11. This research was supported in part by the National Science Foundation under Grant No. PHY99-07949. Some of the $N$-body calculations presented here were carried out at the Center for the Advancement of the Study of Cyberinfrastructure at RIT whose support is gratefully acknowledged."
    },
    "0608/astro-ph0608639_arXiv.txt": {
        "abstract": "We present a set of 11 type Ia supernova (SN~Ia) lightcurves with dense, pre-maximum sampling. These supernovae (SNe), in galaxies behind the Large Magellanic Cloud (LMC), were discovered by the SuperMACHO survey.  The SNe span a redshift range of z = 0.11 -- 0.35. Our lightcurves contain some of the earliest pre-maximum observations of SNe~Ia to date.  We also give a functional model that describes the SN~Ia lightcurve shape (in our $VR$-band).  Our function uses the ``expanding fireball'' model of Goldhaber et al. (1998) to describe the rising lightcurve immediately after explosion but constrains it to smoothly join the remainder of the lightcurve.  We fit this model to a composite observed $VR$-band lightcurve of three SNe between redshifts of 0.135 to 0.165.  These SNe have not been K-corrected or adjusted to account for reddening.  In this redshift range, the observed $VR$-band most closely matches the rest frame $V$-band.  Using the best fit to our functional description of the lightcurve, we find the time between explosion and observed $VR$-band maximum to be 17.6$\\pm$1.3({\\it stat})$\\pm$0.07({\\it sys})~ rest-frame days for a SN~Ia with a $VR$-band $\\Delta m_{-10}$ of 0.52{\\it mag}.  For the redshifts sampled, the observed $VR$-band time-of-maximum brightness should be the same as the rest-frame $V$-band maximum to within 1.1~rest-frame days. ",
        "introduction": "\\subsection{Rise-time as a tool to discriminate between SN~Ia explosion models} The realization that type Ia supernovae (SNe~Ia) can be used as standardizable candles (Phillips 1993; Riess et al. 1995, 1996; Hamuy et al. 1996, 1996b) led to an explosion in SN~Ia science.  Surveys to test the Hubble Expansion Law at larger distances found that rather than exhibiting a constant or decelerating expansion rate, the Universe has an accelerating expansion (Riess et al.  1998, Perlmutter et al. 1999).  The consensus explanation for the accelerating expansion is a negative pressure, or dark energy, permeating the Universe. Today many teams are working to use SNe~Ia as standard candles to better constrain the properties of dark energy (ESSENCE, Matheson et al. 2005; SCP, Kowalski et al. 2005; SNLS, Astier et al. 2006).  While the methods to standardize the SN~Ia luminosity vary, the interpretation of all their results rely to varying degrees on the basic assumption that SNe~Ia belong to a single-parameter family.  Methods of standardizing SN~Ia luminosity distance using the post-maximum lightcurve shape have proven successful when verified against other standard candles such as Cepheids (Suntzeff et al. 1999, Gibson et al. 2000).  These results do not necessarily indicate the existence of a single-parameter family of progenitors, only that the behavior of SNe~Ia post-maximum is similar.  Still, the most widely considered SN~Ia progenitors are carbon-oxygen (C-O) white dwarfs in binary systems.  Even accepting these systems as the progenitors, questions remain concerning the mechanism and progression of the explosion.  Many competing theories (see Hillebrandt \\& Niemeyer 2000 and references therein) predict roughly the same post-maximum behavior and vary only in the prediction of the pre-maximum, or rising, lightcurves and spectra.  Understanding the explosion mechanism may help us better understand how the population of SNe~Ia, and their progenitors, evolves over cosmological time.  Many explosion models are sensitive to progenitor element abundances which may vary depending on the environment.  Combining existing information about the differences between low- and high-z stellar populations and galaxies with a more accurate model of the SN~Ia explosion mechanism will help more tightly constrain the impact of evolution on SN~Ia lightcurve shape.  Discriminating between competing explosion models, however, requires lightcurve coverage close to the time of explosion which has been scarcely available.  The reasons for the lack of early pre-maximum lightcurve coverage are many-fold.  Some SNe~Ia searches rely on a search-and-follow method where SNe are discovered and then followed by another, larger telescope.  Discovery often occurs near maximum brightness, and dense pre-maximum temporal coverage is not available.  Other surveys, similar to SuperMACHO, revisit the same fields every few days, obtaining consistent temporal coverage over the entire lightcurve. These data sets have better pre-maximum coverage but still do not generally provide densely-sampled pre-maximum lightcurves.  In order to maximize the number of fields observed, most surveys use a long, multi-day gap between observations which is sufficient to standardize the post-maximum behavior but often misses the earliest portion of the rise.  For higher-z SNe where the multi-day gap between observations translates to a shorter gap in the SN's rest-frame, the earliest portion of the rise is often too faint to be observed.  As described more completely below, the SuperMACHO data avoid these two pitfalls. This survey provides dense coverage (every other night) and deep imaging with its custom, broadband $VR$ filter.   \\subsection{SuperMACHO and Supernova detection} The SuperMACHO project is a five-year optical survey of the Large Magellanic Cloud (LMC) aimed at detecting microlensing of LMC stars (Stubbs et al. 2002).  The goal of this survey is to determine the location of the lens population responsible for the excess microlensing rate observed toward the LMC by the MACHO project (see Alcock et al. 2000 and references therein) and, thereby, better constrain the fraction of MAssive Compact Halo Objects (MACHOs) in the Galactic halo. The survey is conducted on the CTIO Blanco 4m telescope using a custom $VR$ broadband filter.  SuperMACHO observes 68 LMC fields during dark and gray time in the months of October -- December. We completed our fifth season of observations in the second half of 2005.  We process our images with a near-real time data reduction pipeline that employs a difference-imaging technique (see Alard \\& Lupton 1998, Alcock et al. 1999, Alard 2000, and G{\\\"o}ssl 2002 ) which enables us to detect small changes in flux and to produce lightcurves uncontaminated by light from nearby, non-varying sources.  We present here a uniform set of densely sampled pre-maximum SNe~Ia lightcurves from the SuperMACHO survey.  From these we constrain the time to maximum brightness for SNe~Ia.  We present data to provide constraints on SN~Ia explosion models to aid in discriminating between competing theories.  In Section~\\ref{section:obs} we discuss our observations.  In Section~\\ref{section:data} we present our data.  In Section~\\ref{section:discussion} we use our data to place limits on the time to maximum brightness and present a functional model for the SN~Ia lightcurve shape. ",
        "conclusions": "We present $VR$-band lightcurves and optical spectra of 11 SNe~Ia behind the LMC discovered by the SuperMACHO survey\\footnote{See http://ctiokw.ctio.noao.edu/$\\sim$sm/sm/SNrise for electronic data tables}.  Our data include some of the earliest pre-maximum detections of SNe~Ia.  We provide a functional model for the observed $VR$-band lightcurve from the time of explosion to $+60$ days by fitting a composite lightcurve to three SNe in the redshift bin of z=0.135--0.165.  The data are fitted without K-corrections or reddening corrections; however, the set of SNe have been chosen to minimize the impact of these effects.  Our function uses the expanding fireball model of Goldhaber et al. (1998) to describe the lightcurve immediately following the explosion.  The best fit of our functional model to our composite, observed $VR$-band lightcurve gives a time-to-maximum of 17.6$\\pm$1.3({\\it stat})$\\pm$0.07({\\it sys})~rest-frame days for a SN~Ia with a $\\Delta m_{-10}$ of 0.52{\\it mag}.  Our simulations indicate that the $VR$-band time-of-maximum at z=0.15 should match the rest-frame $V$-band time-of-maximum to within 1.1~rest-frame days.  We present these data to be used to test competing models of the SN~Ia explosion mechanism by placing observational limits on the time to maximum and the shape of the rising lightcurve.  Analyses of our data are limited by its being in a single band.  While our broadband filter enables us to detect flux earlier, we cannot calibrate our lightcurves against the nearby sets of SNe~Ia observed in $BVRI$.\\footnote{Because the transmission curve of our $VR$ filter differs from the sum of the $V$-band and $R$-band transmissions, we cannot simply add the nearby $V$ and $R$ templates to obtain a $VR$ template.}  An ideal study should include both a broadband filter and the standard filter set."
    },
    "0608/hep-ph0608140_arXiv.txt": {
        "abstract": "Models of Universal Extra Dimensions (UED) at the TeV scale lead to the presence of Kaluza Klein (KK) excitations of the ordinary fermions and bosons of the Standard Model that may be observed at hadron and lepton colliders. A conserved discrete symmetry, KK-parity, ensures the stability of the lightest KK particle (LKP), which, if neutral, becomes a good dark matter particle. It has been recently shown that for a certain range of masses of the LKP a relic density consistent with the experimentally observed one may  be obtained. These works, however, ignore the impact of KK graviton production at early times. Whether the $G^1$ is the LKP or not, the $G^n$ tower thus produced can decay to the LKP, and depending on the reheating temperature, may lead to a modification of the relic density. In this article, we show that this effect may lead to a relevant modification of the range of KK masses consistent with the observed relic density. Additionally, if evidence for UED is observed experimentally, we find a stringent upper limit on the reheating temperature depending on the mass of the LKP observed. ",
        "introduction": "Theoretical models constructed to understand physics beyond the standard model (SM) (most notably string theory) frequently imply the existence of extra dimensions. It has been a particular challenge in phenomenology to understand how these extra dimensions would be realized and manifested in our observable $3+1$ dimensional world. The number, shape and size of these dimensions, as well as the particles allowed to propagate in them give rise to several different models, all having different phenomenological implications.  We will be considering universal extra dimensions (UED) where the SM fields propagate in all the extra dimensions. For toroidal compactification, this would imply a tower of Kaluza-Klein (KK) particles for every SM particle, each carrying KK number. Momentum conservation in the extra dimensions implies KK number conservation. However, the requirement of obtaining the proper SM chiral modes at low energies leads to constraints on the possible compactification geometries. An example for $D=5$ is the orbifold $S^1/Z_2$ where the $Z_2$ projects out half of the zero modes, leaving only SM fields. Additionally, this breaks translation invariance along the extra dimensions, so KK number is no longer conserved. A residual symmetry, KK-parity, is still present and it is sufficient to ensure the stability of the lightest KK particle (LKP). It also ensures that KK particles are always produced in pairs, allowing for good agreement between theory and experiment for small values of the compactification scale, of the order of a few hundred GeV~\\cite{Antoniadis:1990ew,Appelquist:2000nn, Appelquist:2002wb}.  The stability of the LKP allows for an interesting candidate for dark matter (DM)~\\cite{Servant:2002aq}. The LKP is expected to be weakly interacting and electrically neutral if it is to be considered a candidate for dark matter. The usual candidate is $B^1$, the KK partner of the hypercharge gauge boson. For $m_{\\rm KK}\\sim\\mathcal{O}~(1)$~TeV, $B^1$ gives excellent agreement with the observed relic density. There have been many further analyses of the relic density in UED, including a more proper treatment of coannihilation effects, and the impact of the inclusion of second KK level resonances~\\cite{Griest:1990kh, Kakizaki:2005en, Kakizaki:2005uy,Kong:2005hn,Burnell:2005hm,Kakizaki:2006dz}.  However, most of these studies ignore the gravitational sector. We shall mostly work within the context of one extra dimension. The gravitons couple extremely weakly and therefore are generally considered to be unimportant for collider studies. Cosmologically, however, they can have a significant effect~\\footnote{The radion can also be quite relevant~\\cite{Kolb:2003mm}. We shall assume that the radion is inert during inflation and that it acquires a large mass, and therefore it does not have any impact on our analysis.}, and they can also be a candidate for the LKP~\\cite{Feng:2003nr}.  We are interested in investigating the inclusion of the graviton tower in the scenario with $B^1$ as the LKP. The gravitons do not evolve thermally, and have very long lifetimes, so if they are present, we expect them to decay to the LKP sometime after Big Bang Nucleosynthesis (BBN). The preservation of the light element abundances sets a bound on the amount of energy that can be released in a decay~\\cite{Ellis:2005ii,Kawasaki:2004qu,Cyburt:2002uv}. Secondly, since the gravitons decay ultimately into the LKP, we expect the relic density to increase. Therefore, the mass of the LKP consistent with the observed relic density is lowered. Additionally, if any of the gravitons decay after matter domination, we have to consider the effects of non-thermalized photons released in the decays on the spectrum of the diffuse photon flux (DPF)~\\cite{Sreekumar:1997yg}.  This article is organized as follows. In Section 2, we review the standard calculation for relic density. Section 3 reiterates the calculation for the density of gravitons presented in Ref.~\\cite{Feng:2003nr}, emphasizing the points relevant for our analysis. In Section 4, we analyze the lifetimes for the decay of the graviton tower and study its implication on the diffuse photon spectrum. Section 5 goes over the constraints on the energy released in the decay of the KK gravitons to the LKP. In Section 6, we combine the above results with the standard relic density calculation. Our goal is to analyze the effect of the gravitons on the predicted mass of the LKP consistent with the known dark matter density. We will show that, in fact, almost any value of $m_{\\rm KK}$ lower than the one obtained in the absence of gravitons would be allowed, provided the reheating temperature, $T_R$, is large enough. For the calculation of the relic density of $B^1$, we ignore complicating factors such as coannihilation and second KK level resonance effects~\\cite{Griest:1990kh, Kakizaki:2005en, Kakizaki:2005uy,Kong:2005hn,Kakizaki:2006dz}. However, since these effects can be parameterized in the effective cross-section, we present the effect the graviton tower will have on the mass of the LKP without the gravitons, $m_{\\rm WG}$, and note that as $m_{\\rm WG}$ increases, the contribution due to the graviton tower becomes large, and should be included in any precise calculation of $m_{\\rm KK}$. We also find that all the gravitons, except for $G^1$, decay right after BBN. At these early times, the electromagnetic BBN constraints are very weak. However it has been shown that at small lifetimes, the hadronic constraints are very stringent~\\cite{Kawasaki:2004yh, Jedamzik:2004er}. Therefore, including constraints from both hadronic and electromagnetic decays, requiring consistency with BBN light element predictions, we find a stringent limit on the mass difference between $m_{G^1}$ and the LKP mass, $m_{\\rm KK}$. Since the $G^1$ is long-lived, we also derive another constraint on the mass difference such that the observed spectrum of the DPF is not destroyed. Comparing the two constraints, we find that there exists a region of parameter space where both constraints are satisfied and which gives mass differences of the same order of magnitude as that obtained by radiative loop corrections to $m_{B^1}$~\\cite{Cheng:2002iz}. Additionally, we observe that if experimental evidence for UED is found, and the relic density induced by standard interactions of the KK modes is found to be lower than the one observed experimentally, the reheating temperature may be determined by assuming that the graviton KK modes provide the contribution necessary to achieve consistency between theory and experiment. Alternatively, including the possibility of other unobserved exotic particles contributing to the relic density, an upper bound on the reheating temperature is obtained. ",
        "conclusions": "In this article, we analyzed the effects of including the KK graviton tower on the determination of the relic density associated with the lightest KK particle in the scenario of Universal Extra Dimensions. Gravitons may be copiously produced in the early universe and their subsequent decay into the LKP may lead to a modification of the observed LKP relic density. Graviton production is governed by a parameter $C \\sim {\\cal{O}}(1)$ and by the reheating temperature, which is bounded from above by the requirement of maintaining the perturbative consistency of the theory, $T_R < 40 m_{\\rm KK}$ ($T_R < 10 m_{\\rm KK}$) for $D =5$ ($D=6$). In the case that Universal Extra Dimensions are observed in laboratory experiments, we show that an upper limit on the reheating temperature can be deduced from the requirement that only the LKP contributes to the observed dark matter density.  Throughout this work we have assumed KK parity conservation, leading to the stability of the LKP particle. We found that including the graviton effectively lowers the lightest KK particle mass consistent with the observed dark matter relic density from values of about 0.92 TeV to values as low as $m_{\\rm KK}\\sim0.58$ TeV for $C=1$, $D=5$ and $T_R<40\\;m_{\\rm KK}$. Additionally, including effects which change the LKP density (coannihilation, second KK-level mode resonant contributions, etc.), we show that the graviton tower has a large impact on the predicted mass of the LKP with increasing mass. It should be stressed here that these results are independent of which KK particle is the LKP (or the NLKP).  Additionally, there are bounds on the mass difference of the LKP with the graviton KK modes induced by the requirement that the energy released in the graviton decay does not lead to a disturbance of the light element abundance or the diffuse photon flux. Under the assumption that the graviton spectrum is $m_{G^n} \\sim n/R$, we have obtained a bound on the mass difference of $G^1$ and the LKP ($B^1$) mass, that is consistent with the minimal one-loop corrections obtained in Ref.~\\cite{Cheng:2002iz} for a large range of values of $m_{\\rm KK}$. ~\\\\ ~\\\\ ~\\\\ \\large{\\textbf{Acknowledgements:}} \\normalsize We would like to thank Jonathan Feng, Csaba Balazs and Tim Tait for useful discussions and comments. Work at ANL is supported in part by the US DOE, Div.\\ of HEP, Contract W-31-109-ENG-38. This work was also supported in part by the U.S. Department of Energy through Grant No. DE-FG02-90ER40560."
    },
    "0608/astro-ph0608196_arXiv.txt": {
        "abstract": "{We present an evolutionary disc model of the edge-on galaxy NGC 5907 based on a continuous  star formation history and a continuous dynamical heating of the stellar subpopulations.} {This model explains the disparate two observational facts: 1) the exponential vertical disc structure in the optical and NIR of the non-obscured part of the stellar disc and 2) the FIR/submm luminosity enhanced by about a factor of four near the obscured mid-plane, which requires additional dust and also stellar light to heat the dust component.} {We use multi-band photometry in U, B, V, R, and I- band combined with radiative transfer through a dust component to simultaneously fit the vertical surface-brightness and colour index profiles in all bands adopting a reasonable star formation history and dynamical heating function. The vertical distribution of the stellar subpopulations are calculated self-consistently in dynamical equilibrium and the intrinsic stellar emissivity is calculated by stellar population synthesis.} {The final disc model reproduces the surface-brightness profiles in all bands with a moderately declining star formation rate and a slowly starting heating function for young stars. The total dust mass is $5.7\\times 10^7\\,\\msun $ as required from the FIR/submm measurements. Without a recent star burst we find in the midplane an excess of 5.2-, 4.0-, and 3.0-times more stellar light in the U-, B-, and V-band, respectively. The corresponding stellar mass-to-light ratios are 0.91 in V- and 1.0 in R-band. The central face-on optical depth in V-band is $\\tau_\\mathrm{V}^{f}=0.81$ and the radial scale length of the dust is 40\\% larger than that of the stellar disc.} {Evolutionary disc models are a powerful method to understand the vertical structure of edge-on galaxies. Insights into the star formation history and the dynamical evolution of stellar discs can be gained. FIR/submm observations are necessary to restrict the parameter space for the models.} ",
        "introduction": "}  \\subsection{Vertical structure of stellar discs}\\label{vertstruct}  The vertical structure of the stellar disc in late-type galaxies is mainly determined by the star formation history combined with the dynamical heating $\\sigma(t)$ of the stellar population due to gravitational perturbations. Young stellar subpopulations are confined tightly to the midplane with small velocity dispersion, whereas the older subpopulations are distributed over larger heights due to their higher velocity dispersion. This basic feature is well-known in the Galaxy (Wielen \\cite{wie77}, Freeman \\cite{fre91}, Edvardsson et al. \\cite{edv93} e.g.). Dynamical heating leads to an increasing mean age of the population with increasing height above the midplane. This results in an increasing mass-to-light ratio and vertical colour index gradients from blue to red. In principle, these colour index gradients should be directly observable in edge-on galaxies but unfortunately, strong dust extinction and reddening often dominates the vertical colour index distribution. Therefore detailed models are necessary to extract the intrinsic stellar disc properties from multi-colour photometry.  In a series of papers van der Kruit \\& Searle (\\cite{vdk81,vdk82a,vdk82b}) started to investigate the vertical and radial properties of the surface-brightness profiles of edge-on galaxies. They determined for a number of nearby galaxies (including NGC 5907) the radial scale length, the radial cutoff, and the vertical scale length. They recognised that the shape of the vertical surface-brightness profiles can be well fitted by a $sech^2$-model disregarding the central minimum in the case of strong dust lanes. In recent decades much work has been done on the phenomenological analysis of edge-on galaxies (e.g. Guthrie \\cite{gut92}, Pohlen et al. \\cite{poh02}).  In our Galaxy the stellar subpopulations of the thin disc are not homogeneously mixed. The vertical thickness of the  subpopulations increases with age (Wielen \\& Fuchs \\cite{wie88}). Wainscoat et al. (\\cite{wai89}) observed a similar vertical structure in the edge-on galaxy IC 2531. For the analysis of the intrinsic physical structure of galactic discs we started to investigate the vertical surface-brightness and colour index profiles of edge-on galaxies in order to analyse the age distribution of stellar discs (Wielen et al. \\cite{wie92}, Just et al. \\cite{jus96}). We used evolutionary stellar population models for the intrinsic light distribution and include dust extinction in the radiative transfer for the observable vertical surface-brightness profiles. We found that in edge-on galaxies the vertical colour index gradients above the innermost part, where extinction dominates, are a signature of the changing properties of the stellar population with height above the midplane. De Grijs and Peletier (\\cite{dgr00}) also found from a large sample of edge-on galaxies that the stellar population properties in the discs can be seen in the vertical colour index gradients above the extinction features.  For NGC 5907 and other nearby edge-on galaxies Xilouris et al. (\\cite{xil97,xil99}) used an extensive radiative transfer code including scattered light to model the two-dimensional light distribution of the stellar discs. Using exponential intrinsic luminosity distributions different for each band (B,V,I in the case of NGC 5907) and an exponential disc for the dust component, they determined radial and vertical scale lengths of the star light and the properties of the dust component. The resulting relative low dust mass corresponds to a gas-to-dust ratio of 810 (corrected for the Helium contribution). In recent years FIR observations offered a new, more direct way to deduce the dust properties of galaxies. Popescu et al. (\\cite{pop00}) developed a sophisticated temperature model of the dust component in edge-on galaxies in order to fit the FIR observations and applied it to NGC 891. With this method Misiriotis et  al. (\\cite{mis01}) determined for NGC 5907 a dust mass 4 times larger than Xilouris et al. (\\cite{xil99}). For the source of the FIR excess they adopted an additional young stellar component in a very thin disc, which is completely hidden by dust extinction in the optical and NIR bands. Recently Stevens et al. (\\cite{ste05}) used a two-temperature dust model to fit the IRAS and SCUBA data for a sample of spiral galaxies. For NGC 5907 they determine a dust mass only three times larger than the Xilouris et al. value.   \\subsection{Content of this paper}\\label{content}   \\begin{figure}[t] \\hbox{\\hspace{0cm} \\psfig{figure=5701fig1small.ps,width=8.5cm,clip=}} \\caption[]{Deep R image of NGC 5907, sum of 4 exposures. The field of view is $15\\times 13.25$ arcmin. N is top, E is left. (cf Sect.~\\ref{observ})} \\label{figimagr} \\end{figure}  In this paper we present an evolutionary stellar disc model for NGC 5907, which explains both observational results: 1) The exponential vertical profiles in all optical bands  are reproduced by the outer stellar disc dominated by the older thin disc population and are no longer just a set of independent fit functions in each band as in Xilouris et al. (\\cite{xil99}). 2) The large amount of dust determined from FIR/submm observations by Misiriotis et al. (\\cite{mis01}) exceeding the dust component found in Xilouris et al. by a factor of four is necessary in our model to obscure the additional light of the young stellar subpopulations with low mass-to-light ratio near the midplane of the disc. There is no additional recent star burst required to explain the additional star light needed for the dust heating.  The basis of our disc model is a vertical cut with a self-consistent detailed distribution of stellar subpopulations. The vertical distribution of the stars is determined by the dynamical equilibrium of the stellar subpopulations described by the star formation history of the disc and by the dynamical heating function (measured by the increasing velocity dispersion). The gravitational forces of the gas component and the dark matter halo are included. The intrinsic stellar light distribution is calculated from population synthesis models of simple stellar populations (SSP). The resulting vertical profile is then extended to an exponential disc up to the cutoff radius $R_\\mathrm{max}$ by scaling the surface densities with constant scale heights. The intrinsic structure is not changed as a function of radius. We compute the observable vertical surface-brightness and colour index profiles by radiative transfer calculations with an exponential dust component, where the inclination of the disc with respect to the line of sight is fitted.  The parameters of the basic vertical profile are chosen to match all constraints at some fixed radius $R_0$ of the disc, where the disc is dominating the light and the bulge can be neglected. At that radius we fix the relative contribution of the gas and dark matter component $Q_\\mathrm{g}=\\Sigma_\\mathrm{g}/\\Sigma_\\mathrm{tot}$ and $Q_\\mathrm{h}=\\Sigma_\\mathrm{h}/\\Sigma_\\mathrm{tot}$ to the total surface density $\\Sigma_\\mathrm{tot}$ up to a  maximum height $z_\\mathrm{max}$ above the midplane. We will use $R_0=10\\,\\mathrm{kpc}\\approx2\\,R_\\mathrm{s}$ with radial scale length $R_\\mathrm{s}$ of the stellar disc, which is near the maximum of $\\Sigma_\\mathrm{d}/\\Sigma_\\mathrm{h}$ in the case of an isothermal halo. The scale height of the gas component relative to the effective scale height of the stars $z_\\mathrm{g}/z_\\mathrm{0}$ and $s_\\mathrm{h}=\\sigma_\\mathrm{h}/\\sigma_\\mathrm{e}$, the ratio of the velocity dispersion of the dark matter halo and the maximum velocity dispersion of the stars are also determined at $R_0$.  In an iterative process the crucial fitting functions $SFR(t_\\mathrm{a}-t)$ and $\\sigma(t)$ are optimised together with the global parameters: inclination $i$, radial scale length $R_\\mathrm{s}$, effective scale height $z_\\mathrm{s}$, dust distribution, and total stellar and dust mass. For comparison with observations we use a series of vertical surface-brightness (and colour index) profiles parallel to the minor axes in U,B,V,R,I bands, which we obtained from corresponding deep photometric images of NGC 5907. The same type of intrinsic disc model was also successfully applied to the solar neighbourhood in order to analyse the connection between the local velocity distribution functions of the main sequence stars from the Catalogue of Nearby Stars (CNS4, Jahreiss \\& Wielen \\cite{jah97}) and the vertical density profiles of these stars (see Just \\cite{jus01,jus02,jus03}).  In Sect. \\ref{param1} we collect the basic parameters of NGC 5907, which are necessary for the basic scaling of our disc model to the observational data. In Sect. \\ref{components} we describe the building blocks of our disc model. Sect. \\ref{project} gives the iteration procedure with details on the scaling method, radiative transfer and the multi-colour analysis. In Sect. \\ref{observ} our observations and the derivation of vertical colour index and surface-brightness profiles are presented. Fig.~\\ref{figimagr} shows a deep R image (cf Section~\\ref{observ} as an example of for the high quality of data. In Sect. \\ref{result} we show the final comparison of the profiles, give the global parameters of our best model and discuss in detail the intrinsic structure of the stellar disc. A discussion of some crucial aspects follow in Sect. \\ref{discuss}. Sect. \\ref{conclusion} contains the concluding remarks. ",
        "conclusions": "}  We have observed deep photometry of NGC 5907 in the U, B, V, R, and I-band in order to construct a physical disc model based on a multi-colour analysis. We have constructed a self-consistent evolutionary disc model for NGC 5907, which fits the general characteristics of the surface brightness distribution in all bands. The stellar disc is built up by a sequence of stellar subpopulations distributed according to a smooth $SFR$ and dynamical heating function $\\sigma(t)$. The emissivity is calculated with stellar population synthesis and the appearance of the galaxy is calculated by radiative transfer including an exponential disc component of dust.  This is the first physical model of the disc of NGC 5907 that gives the correct luminosities and colours of the outer exponential parts of the disc and also the large amount of dust observed in FIR/submm. The reason for the corresponding excess in stellar light near the midplane of the disc is the concentrated luminous young subpopulation produced by the continuous $SFR$. Our model is a significant improvement of the disc model of Misiriotis et al. (\\cite{mis01}), who combined the phenomenological exponential model of Xilouris et al. (\\cite{xil99}) with an additional thin disc of uniformly mixed young stars and dust in order to produce the excess in FIR/submm luminosity.  This evolutionary disc model naturally explains three different aspects known from the literature: 1) The exponential vertical structure well above the midplane as modelled phenomenologically by Xilouris et al.(\\cite{xil99}) is the result of the older subpopulations in dynamical equilibrium in the total disc potential. 2) The large amount of dust observed in the FIR/submm (modelled by Misiriotis et al. \\cite{mis01}) is necessary to hide the bright young subpopulations confined to the midplane. The young stars are the smooth end of the $SFR$ without a recent star burst. 3) The additional light necessary to heat the dust is automatically delivered by the high mass-to-light ratio near the midplane."
    },
    "0608/astro-ph0608475_arXiv.txt": {
        "abstract": "We observed the bright unidentified TeV $\\gamma$-ray source \\hessLHM\\ with the X-ray Imaging Spectrometers onboard the Suzaku satellite.  No X-ray counterpart was found to a limiting flux of 3.1\\EHM{-13}~\\fluxHM\\ in the 2--10~keV band, which is some 60 times below the $\\gamma$-ray flux in the 1--10~TeV band. This object is bright in TeV $\\gamma$-rays but very dim in the X-ray band, and thus is one of the best examples in the Galaxy of a \"dark particle accelerator.\"  We also detected soft thermal emission with $kT \\sim$0.3--0.6~keV near the location of \\hessLHM.  This may be due to the dust grain scattering halo from the nearby bright supernova remnant RCW~103. ",
        "introduction": "Since the discovery of cosmic rays, the origin of these high energy particles has been a mystery. Recent observational studies have revealed that some supernova remnants (SNRs) have non-thermal X-ray emission (e.g. \\cite{Koyama1995,Hwang2002,Bamba2003,Vink2003,Bamba2005}), which can be interpreted as synchrotron radiation from electrons with energies approaching $E \\sim 10^{14}$~eV. TeV $\\gamma$-rays have also been detected from some of these non-thermal shell-type SNRs.  Such very high energy (VHE) gamma-rays have been explained by either (1) Inverse-Compton (IC) upscattering of cosmic microwave background (or other lower frequency) photons by the same high energy electrons giving rise to the X-ray synchrotron emission or (2) the decay of neutral pions that originate in collisions between high energy protons and dense interstellar matter (e.g. \\cite{Pannuti2003,Lazendic2004,Katagiri2005}).  Thus the combination of VHE gamma-ray and non-thermal X-ray emission is expected to be a key observational feature of high energy particle accelerators. A high sensitivity survey of the Galactic plane in VHE $\\gamma$-rays has been conducted by the High Energy Stereoscopic System (HESS) team, who discovered fourteen new VHE $\\gamma$-ray sources \\citep{Aharonian2005,Aharonian2006}. \\hessLHM\\ (hereafter \\hessHM) is one of the brightest of these new sources and, furthermore, is spatially extended with an angular diameter of $\\sim$\\timeform{16'} in the HESS data.  Three known objects, the young hard X-ray pulsar PSR~J1617$-$5055, the SNR RCW~103 (G332.4$-$0.4), and the SNR Kes~32 (G322.4$+$0.1), are in the vicinity of \\hessHM\\ ($\\sim$ \\timeform{10'}--\\timeform{15'} away from it), but none provides a convincing identification~\\citep{Aharonian2005,Aharonian2006}.  Before the HESS Galactic plane survey, a few extended $\\gamma$-ray objects with no clear counterpart in other wavebands had been found. These objects have been denoted ``dark particle accelerators''\\citep{Ubertini2005}.  They are of interest since they may be indicating the acceleration of nucleons; on the other hand, if the TeV emission comes from electrons, peculiar conditions such as extremely low magnetic fields may be necessary, due to the short radiative lifetimes of high energy electrons~\\citep{Yamazaki2006}. TeV~2032+410, discovered by the High Energy Gamma-Ray Astronomy (HEGRA) collaboration~\\citep{Aharonian2002}, was the first object of this kind.  In spite of several multiwavelength follow-up studies, no compelling counterpart has been found from radio to X-ray wavelengths~\\citep{Mukherjee2003,Butt2003,Butt2006}. HESS~J1303$-$631 is another example~\\citep{Aharonian2004}. \\citet{Mukherjee2003} observed this target with the Chandra X-ray Observatory for 5~ks, and placed an upper limit of $<$5.4\\EHM{-12}\\fluxHM\\ to its diffuse hard X-ray flux in the 2--10~keV band.  \\hessHM\\ is brighter than both TeV~2032$+$410 and HESS~J1303$-$631 in the TeV regime (see also table~\\ref{tbl:VHEsrc}) and is therefore a more suitable case for examining if dark accelerators really exist and for finding out just how ``dark'' they are.  Since many VHE $\\gamma$-ray objects are X-ray emitters, the X-ray band is the natural place to search for a counterpart.  To date there have been no X-ray observations explicitly targeting \\hessHM, so we observed \\hessHM\\ with the Suzaku satellite~\\citep{Mitsuda2006}.  The Suzaku X-ray Imaging Spectrometers (XIS)~\\citep{Koyama2006}, with their high sensitivity and stable low background~\\citep{Yamaguchi2006}, are optimal devices to search for diffuse X-ray emission. This is especially the case in the hard X-ray band ($E \\gtrsim 6$~keV), where non-thermal emission dominates over thermal and the effects of X-ray absorption are minimized. By detecting or even just setting a sensitive upper limit on the X-ray emission from \\hessHM, it should be possible to address the question of what types of particles, electrons or protons, are accelerated in the object, and to constrain the acceleration mechanism by comparing the X-ray and $\\gamma$-ray emission.  In addition to the Suzaku data we also used XMM-Newton archival data that partly covered the \\hessHM\\ region.  In this work, uncertainties are quoted at the 90\\% confidence level unless otherwise stated. ",
        "conclusions": "Neither the Suzaku nor the XMM-Newton data provide evidence for hard X-ray emission from \\hessHM; the XIS sets the most stringent upper limit of 3.1\\EHM{-13}\\fluxHM\\ in the 2--10~keV band.  The differential photon flux in the TeV $\\gamma$-ray regime can be represented by a power-law of photon index $\\Gamma=2.35$ with a total photon flux of ($43.3\\pm2.0$)\\EHM{-12}~cm$^{-2}$~s$^{-1}$ above 200~GeV~\\citep{Aharonian2006}.  This corresponds to an energy flux of 1.7\\EHM{-11}\\fluxHM\\ in the 1--10~TeV band. The ratio of the TeV $\\gamma$-ray flux to the X-ray flux ($f_{\\rm TeV}/f_{\\rm X}$) is therefore more than $\\sim$55. Table~\\ref{tbl:VHEsrc} is the list of spatially extended VHE objects with X-ray observations that we extracted from the H.E.S.S. source catalog~\\footnote{http://www.mpi-hd.mpg.de/hfm/HESS/public/HESS\\_catalog.htm}. We also added TeV~J2032+4130~\\citep{Aharonian2002} to the list.  Although recent Swift observations found possible counterparts of HESS~1614$-$518 and HESS~J1834$-$087~\\citep{Landi2006}, we do not include them in table~\\ref{tbl:VHEsrc}, since their X-ray data are not statistically enough for spectroscopy. Note that the X-ray couterparts in the list are not secure. For example, AX~J1838.0$-$0655 is the most promising candidate counterpart of HESS~J1837$-$069~\\citep{Aharonian2005,Aharonian2006}, it is slightly outside of the HESS source extension~\\citep{Landi2006}. \\hessHM\\ has the most stringent X-ray flux upper limit among the VHE objects; yet of more import is its extremely large flux ratio, $f_{\\rm TeV}/f_{\\rm X}$, compared to the others.  Evidently \\hessHM\\ is a very peculiar object.  The bright TeV $\\gamma$-ray emission strongly suggests the presence of some particle acceleration processes there; however \\hessHM\\ emits little X-ray (or lower frequency radiation) emission and thus remained undiscovered prior to the HESS survey. \\hessHM\\ is thus a strong candidate for being a ``dark particle accelerator.''   \\begin{table*} \\caption{Spatially extended VHE objects with X-ray observations.} \\label{tbl:VHEsrc} \\begin{center} \\small \\begin{tabular}{cccccccccc} \\hline \\hline Name   & Possible counterpart &Type$^a$ & $\\Gamma_{\\rm TeV}^b$ & $f_{\\rm TeV}^c$ &$N_{\\rm H}^d$ & $\\Gamma_{\\rm X}^e$  & $f_{\\rm X}^{f}$ &$f_{\\rm TeV}/f_{\\rm X}$ &Reference$^g$\\\\ \\hline HESS~J0852$-$463 &RX~J0852$-$4622 &SNR &2.1 &6.9 &4   &2.6 &$\\sim$10 &$\\sim$ 0.7 &1, 2, 3\\\\ HESS~J1303$-$631 &---           &?   &2.4 &1.0 &20  &2.0 &$<$0.64&$>1.6$ &4, 5\\\\ HESS~J1514$-$591 &PSR~B1509$-$58  &PWN &2.3 &1.6 &8.6 &2.0 &3.2 &0.5 & 6, 7\\\\ HESS~J1632$-$478 &AX~J1631.9$-$4752 &HMXB?&2.1&1.7&210&1.6 &1.7 &1.0 &8, 9\\\\ HESS~J1640$-$465 &G338.3$-$0.0    &SNR &2.4 &0.71 &96 &3.0 &0.30 &2.4 &8, 10\\\\ HESS~J1713$-$397 &RX~J1713.7$-$3946 &SNR&2.2&3.5 &8 &2.4 &54 &0.065&11, 12\\\\ HESS~J1804$-$216 &Suzaku~J1804-2142 &?   &2.7 &0.48&2&-0.3&0.025&19&8, 13\\\\ HESS~J1804$-$216 &Suzaku~J1804-2140 &?   &2.7 &0.48&110&1.7&0.043&11&8, 13\\\\ HESS~J1813$-$178 &AX~J1813$-$178    &?   &2.1 &0.89&110&1.8&0.70&1.3&8, 14\\\\ HESS~J1837$-$069 &AX~J1838.0$-$0655 &?   &2.3 &1.4 &40 &0.8&1.3&1.1&8, 15\\\\ TeV~J2032$+$4130 &---             &?   &1.9 &0.20&?  &?  &$<$0.20&$>$1.0&16\\\\ \\hline HESS~J1616$-$508 &---             &?   &2.4 &1.7&4.1&2.0 &$<$0.031&$>$55&This work\\\\ \\hline \\end{tabular} \\end{center} $a$: SNR=supernova remnant, PWN=pulsar wind nebula, HMXB=high mass X-ray binary\\\\ $b$: Photon index of TeV spectra.\\\\ $c$: Unabsorbed flux in the 1--10~TeV band (in $10^{-11}$~\\fluxHM).\\\\ $d$: Column density for X-ray spectra (in $10^{21}$~cm$^{-2}$).\\\\ $e$: Photon index of X-ray spectra.\\\\ $f$: Unabsorbed flux in the 2--10~keV band (in $10^{-11}$~\\fluxHM).\\\\ $g$: (1)~\\citet{Aharonian2005aap437}, (2)~\\citet{Slane2001}, (3) Our analysis of the ASCA archival data, (4)~\\citet{Aharonian2004}, (5)~\\citet{Mukherjee2005}, (6)~\\citet{Aharonian2005aap435}, (7)~\\citet{Delaney2006}, (8)~\\citet{Aharonian2006}, (9)~\\citet{Rodriguez2003}, (10)~\\citet{Sugizaki2001}, (11)~\\citet{Aharonian2004Natur}, (12)~\\citet{Slane1999}, (13)~\\citet{Bamba2006}, (14)~\\citet{Brogan2005}, (15)~\\citet{Bamba2003apj}, and (16)~\\citet{Aharonian2002} \\end{table*}    Assuming the origin of the TeV $\\gamma$-ray to be IC scattering of the cosmic microwave background by accelerated electrons, we can calculate the synchrotron emission from these electrons assuming magnetic field values of $B=$10, 1, and 0.1~$\\mu$Gauss (figure~\\ref{fig:sed}).  The Suzaku flux limit requires that the magnetic field in \\hessHM\\ be less than a few micro Gauss, which is the typical interstellar value and far below the estimated magnetic field values in other HESS sources such as RX~J0852$-$4622 and RX~J1713.7$-$3946. Thus other mechanisms for the TeV emission may be required. \\citet{Aharonian2006}\\ proposed an asymmetric undetected pulsar wind nebula (PWN) powered by the young pulsar PSR~J1617$-$505, which is $\\sim$\\timeform{13'} away from \\hessHM, to explain the TeV emission.  The spin-down luminosity of PSR~J1617$-$505 is 1.6\\EHM{37}\\lumiHM~\\citep{Torii1998}, and in this case the luminosity of the PWN is expected to be $\\sim$$10^{34}$~\\lumiHM~\\citep{Cheng2004} in the 2--10~keV band.  Assuming the pulsar distance to be 3.3~kpc~\\citep{Torii1998}, we obtain an unabsorbed flux of 7.6\\EHM{-12}~\\fluxHM\\ in the 2--10~keV band. To reconcile this high flux and the XIS upper limit, we would need an absorption of $N_{\\rm H}>2\\EHM{24}$~cm$^{-2}$ for a photon index of 2.  Since the column density is more than 10 times the absorption to the Galactic center~\\citet{Baganoff2003}, this seems an unlikely scenario to explain X and TeV emission.  The high $f_{\\rm TeV}/f_{\\rm X}$ ratio might be explained if \\hessHM\\ is an old SNR colliding with a giant molecular cloud, since TeV and X-ray emissions are dominated by hadronic processes and synchrotron radiation from secondary electrons, respectively~\\citep{Yamazaki2006}. However there are no radio data indicating the presence of a giant molecular cloud in the region.  In any case, it remains highly desirable to obtain detailed data of the \\hessHM\\ region in other wavelengths, especially the radio band, to detect the synchrotron emission from the accelerated electrons.   We found extended soft X-ray emission suggesting a thermal plasma with $kT\\sim0.5$~keV both in the Suzaku and XMM spectra. Both the temperature and the projection profiles of the soft band images (figures~\\ref{fig:xis_prof} and \\ref{fig:xmm_prof}) imply a close relationship between the soft emission and RCW~103.  The best-fit $N_{\\rm H}$ of 2--4\\EHM{21}~cm$^{-2}$ in table~\\ref{tbl:results} also support the connection, since these values are close to those obtained with the ASCA spectra~\\citep{Gotthelf1997}. The XMM-Newton profile (figure~\\ref{fig:xmm_prof}) suggests it extends up to \\timeform{15'} from the central point source 1E161348-5055, which is three times as large as the radius of RCW~103. Although we cannot fully reject the possibility, at least for the Suzaku data, that this soft emission can be explained by an instrumental effect, the presence of this emission in the XMM-Newton data argue strongly for an astrophysical origin.  Many bright X-ray sources in the Galactic plane have spatially extended X-ray halos, caused by the scattering of X-rays into our line of sight by interstellar dust between  us and the object~\\citep{Predehl1995}. The fact that the hard band profile of XMM-Newton does not extend beyond the rim of RCW~103 may support the dust halo hypothesis, since the scattering cross section ($\\sigma_{\\rm dust}$) depends strongly on photon energy as $\\sigma_{\\rm dust} \\propto E^{-2}$~\\citep{Predehl1995}. We will address the soft emission again in future when the calibrations have become much robust.  With careful work it may be possible to model the soft extended contamination over the region of \\hessHM\\ and further reduce the hard band X-ray flux limit.   \\begin{figure} \\begin{center} \\FigureFile(0.45\\textwidth,){figure9.eps} \\end{center} \\caption{Spectral energy distribution of \\hessHM\\ from the X-ray to TeV $\\gamma$-ray band. The synchrotron radiation from accelerated electrons, which boost the 3~K background up to the TeV energy range, is plotted toward the left for three different values of the magnetic field. } \\label{fig:sed} \\end{figure}"
    },
    "0608/astro-ph0608643_arXiv.txt": {
        "abstract": "\\indent The Shear TEsting Programme (STEP) is a collaborative project to improve the accuracy and reliability of weak lensing measurement, in preparation for the next generation of wide-field surveys. We review sixteen current and emerging shear measurement methods in a common language, and assess their performance by running them (blindly) on simulated images that contain a known shear signal. We determine the common features of algorithms that most successfully recover the input parameters. A desirable goal would be the combination of their best elements into one ultimate shear measurement method. In this analysis, we achieve previously unattained discriminatory precision via a combination of more extensive simulations and pairs of galaxy images that have been rotated with respect to each other. That removes the otherwise overwhelming noise from their intrinsic ellipticities. Finally, the robustness of our simulation approach is confirmed by testing the relative calibration of methods on real data.   Weak lensing measurement has improved since the first STEP paper. Several methods now consistently achieve better than $2\\%$ precision, and are still being developed. However, we can now distinguish all methods from {\\it perfect} performance. Our main concern continues to be the potential for a multiplicative shear calibration bias: not least because this can not be internally calibrated with real data. We determine which galaxy populations are responsible and, by adjusting the simulated observing conditions, we also investigate the effects of instrumental and atmospheric parameters. We have isolated several previously unrecognised aspects of galaxy shape measurement, in which focussed development could provide further progress towards the sub-percent level of precision desired for future surveys. These areas include the suitable treatment of image pixellisation and galaxy morphology evolution. Ignoring the former effect affects the measurement of shear in different directions, leading to an overall underestimation of shear and hence the amplitude of the matter power spectrum. Ignoring the second effect could affect the calibration of shear estimators as a function of galaxy redshift, and the evolution of the lensing signal, which will be vital to measure parameters including the dark energy equation of state. ",
        "introduction": "The observed shapes of distant galaxies become slightly distorted by the (differential) gravitational deflection of a light bundle as it passes near foreground mass structures. Such ``cosmic shear'' happens regardless of the nature and state of the foreground mass.  It is therefore a uniquely powerful probe of the cosmic mass distribution, dominated by dark matter. Observations of gravitational lensing are directly and simply linked to theories of structure formation that are otherwise ill-equipped to predict the distribution of light \\citep[for reviews, see][]{Bible,witrev,refrev}. Measurements are {\\it not} limited by astrophysical bias \\citep[\\eg][]{oferbias,meg901,hoebias,smithbias,weinbias}, which affects optical surveys, nor by unknown physics of distant supernov\\ae\\ \\citep[\\eg][]{sne_review,sne_variety,sne_environ,sne_metals}, nor by the uncertain relations between the mass of galaxy clusters and their observable X-ray luminosity or temperature \\citep[\\eg][]{hutw,pier,vlnew}. Gravitational lensing is a purely geometric effect, requiring knowledge of only deflection angles and distances. By directly observing the growth of the mass structures over cosmic time, and by investigating the large-scale geometry of the universe, it is also an effective probe of dark energy \\citep{cfhtls_deep,cfhtls_wide,jarvis06,quintessence} and can test alternative theories of gravity that move beyond general relativity \\citep{lensinggravity}.  The practical use of weak lensing in cosmology effectively began with the simultaneous detection of a coherent cosmic shear signal by four independent groups \\citep{BRE,Kaiser_orig,vW_orig,Wittman_nature}. Since then, the field of weak lensing has advanced dramatically. Large, dedicated surveys with ground- and space-based telescopes have recently measured the projected 2D power spectrum of the large-scale mass distribution and drawn competitive constraints on the matter density parameter $\\Omega_m$ and the amplitude of the matter power spectrum $\\sigma_8$ \\citep{Maoli,RRG01,vWb01,HYG02,BMRE,RRG02,Jarvis,MLB02,Hamana,Massey,RhodesSTIS,vWb04,HymzGEMS,Jarvis05, cfhtls_wide,cfhtls_deep,Hettcs,gemscs2,dahleclusters}. The results from these efforts are found to be in broad agreement and are rapidly becoming more credible, with the most recent publications presenting several different diagnostic tests to determine the levels of systematic error. Ambitious plans are being laid for dedicated telescopes both on the ground (\\eg\\  VST-KIDS, DES, VISTA darkCAM, Pan-STARRS, LSST) and in space (\\eg\\ DUNE, SNAP, JDEM). Indeed, future weak lensing surveys were recently identified as the most promising route to understanding the nature of dark energy by the joint NSF-NASA-DOE Astronomy and Astrophysics Advisory Committee (AAAC) and NSF-DOE High Energy Physics Advisory Panel (HEPAP) Dark Energy Task Force\\footnote{\\tt http://www.nsf.gov/mps/ast/detf.jsp}. The importance of weak lensing in future cosmological and astrophysical contexts seems assured.  However, the detection and measurement of weak gravitational lensing presents a technical challenge. The $\\sim1\\%$ distortion induced in the observed shapes of galaxies is an order of magnitude smaller than their typical intrinsic ellipticities, and a similar factor smaller than the spurious shape distortions created by convolution with the telescope's point spread function (PSF). Correction for these effects is crucial and complex. To test the reliability of weak lensing measurements, it has therefore been necessary since the first detections to manufacture simulated images that closely resemble real data but contain a known shear signal. \\citet{Baconsims}, \\citet{erben} and \\citet{HYGBHI} ran their shear measurement methods on such images. By comparing the input and mean measured shears, they determined the calibration error inherent to each technique, and in some cases discovered (and hence corrected) a multiplicative calibration bias. This is most important because it cannot be self-calibrated from a survey itself. Other systematics can be checked for in real data via correlation of the galaxies and the PSF, or via an $E$-$B$ decomposition \\citep{SchvWM02,CNPT02,EBring}. These early tests determined that the first successful shear measurement methods were accurate to $\\le10\\%$ of the signal.  To maximise progress in this technical field, and to foster the exchange of data and theoretical knowledge within the weak lensing community, we launched the Shear TEsting Programme (STEP). In the first STEP paper, \\citep[STEP1]{step1}, we parametrized the performance of methods in terms of their multiplicative shear calibration bias $m$, an additive residual shear offset $c$ and, in some cases, a nonlinear responsivity to shear $q$. That analysis confirmed that the main difficulty in weak lensing lies in the calibration of the shear signal, but encouragingly showed that all of the methods used on existing weak lensing surveys achieve better than $\\sim7\\%$ accuracy. Shear measurement error is therefore not currently a dominant source of error.  Unfortunately, this accuracy will not be sufficient to realise the potential of the ambitious and much larger future surveys. STEP1 found that the most accurate shear measurement methods were successfully calibrated to within a few percent, but the limited size and precision of the first STEP simulations forbade any finer analysis than this. The morphologies of galaxies in the first simulated images were also overly simplistic, in a way that did not fully test the assumptions of some shear measurement methods that galaxies lack substructure and complex shapes.  In this second STEP paper, we include complex galaxy morphologies and conduct a more precise test of current and developing shear measurement algorithms to the $\\le0.5\\%$ level. We achieve this precision through the combination of a more extensive set of simulated images and an ingenious use of galaxy pairs rotated with respect to each other \\citep{reiko}. This removes the otherwise dominant noise from galaxies' intrinsic ellipticities. The new set of simulated images has also been designed  to span a wide range of realistic observing conditions and isolate several potentially challenging aspects of shear calibration in which the accuracy of shear recovery may begin to deteriorate. The data set is sufficiently large for it to be divided into different simulated observing conditions and for independent tests to be carried out within each. We thereby test the effects of the following parameters on shear measurement precision:  \\newpage  \\begin{itemize} \\item Complex galaxy morphology \\item Galaxy size \\item Galaxy magnitude \\item Selection effects related to galaxy ellipticity \\item Direction of the shear signal relative to the pixel grid \\item PSF size \\item PSF ellipticity \\end{itemize}  Sixteen different shear measurement codes have been run on the simulated images. These can be categorised into four distinct categories. We provide a brief description of each algorithm, and outline the relative successes of each method. The STEP programme has dramatically sped the development of new shear measurement methods \\citep[\\eg][Bridle \\etal\\ in preparation]{shapelets2,bj02,shapelets3,shapeletskk,reiko}, and we particularly focus on these. However, these methods necessarily remain experimental, and development continues. The results from such methods should therefore be taken as an indication of progress rather than a judgement on their ultimate potential.  This paper is organised as follows.  In \\S\\ref{sec:data}, we describe the simulated images. In \\S\\ref{sec:methods}, we review the different shear measurement methods used by each author, translating them into a common language for ease of comparison, and categorising them into four distinct groups. In \\S\\ref{sec:results}, we compare each author's measured shear with the input signal, and split the simulations in various ways to isolate areas of potential difficulty in shear measurement. Because of the number of different methods used, this is a rather daunting process. In \\S\\ref{sec:discussion}, we provide some perspective on the results, assessing the relative performance of the different methods, and the categories of methods. In \\S\\ref{sec:conc}, we derive some general conclusions and outline suggestions for future development. ",
        "conclusions": ""
    },
    "0608/astro-ph0608369_arXiv.txt": {
        "abstract": "{% } {% We discuss the applicability and reliability of the shapelet technique for scientific image analysis. } {% We quantify the effects of non-orthogonality of sampled shapelet basis functions and misestimation of shapelet parameters. We perform the shapelet decomposition on artificial galaxy images with underlying shapelet models and galaxy images from the GOODS survey, comparing the publicly available IDL implementation with our new C++ implementation. } {% Non-orthogonality of the sampled basis functions and misestimation of the shapelet parameters can cause substantial misinterpretation of the physical properties of the decomposed objects. Additional constraints, image preprocessing and enhanced precision have to be incorporated in order to achieve reliable decomposition results. }{} ",
        "introduction": "\\label{sec:introduction} Image analysis is \\emph{the} crucial technique and the prime objective in all observational sciences. Computer-based image analysis is supposed to provide us with information that human perception is not or only hardly capable of extracting, due to either the vast amount of data or the required accuracy. The decomposition of imaged objects into an orthogonal function set and analysis of the expansion coefficients is the preferred way of conducting image analysis since sufficient computer power is available.  \\cite{shapeletsI} proposed the 'shapelets' function set, composed of a scalable version of the eigenfunctions of the harmonic oscillator in quantum mechanics. They form an orthonormal basis set of Gauss-Hermite polynomials. Because of their continuity, finite extent and their smooth peaks, they offer themselves for decomposing galaxy images or the like. In particular, they were proposed as an image processing and analysis technique for weak-lensing studies \\citep{shapeletsII,weak-lensing-LSS-FIRST,cosmic-shear-WHT, galaxy-galaxy-flexion,shear_shapelets,STEP1,STEP2}, morphological classification of galaxies \\citep{galaxy-morphology_shapelets} and Sun spots \\citep{sun-spots_shapelets}, and also in the field of medical computer tomography \\citep{skin-tomography_shapelets}.  So far, the shapelet technique has proven to reconstruct the decomposed objects from the set of expansion coefficients in a way which looks visually good; i.e. differences to the originals are marginal \\citep{shapeletsI}. \\cite{shapeletsIII} defined a goodness-of-fit measure to quantify, what a good reconstruction means: the residuals have to be at noise level. For achieving that, the decomposition has to be optimized with respect to a set of parameters. As the parameters space is highly degenerate, additional constraints are necessary to choose a particular point in this space. But it is yet unknown, if the selection of these constraints or the accuracy with which their fulfillment is verified introduces an uncertainty in or even a bias on the expansion coefficients. In fact, every study of object properties derived from shapelet coefficients might be affected by a yet unclear error contribution.  Another object of concern is the computational complexity. Since the shapelet method is known to be slow compared to other image analysis techniques, it is important to find ways to speed up the execution.  This paper is organized as follows: In sect. 2, we summarize the basic relations for the shapelet function set and describe the procedure for finding optimal decomposition parameters. In sect. 3, we discuss potential problems that can arise from the optimized decomposition procedure. In sect. 4, we show how these problems can be remedied by means of additional constraints and image preprocessing. In sect. 5, we compare the design choices, the decomposition results (of artificial and observed galaxy images), the errors made by and the computational performance of two shapelet implementations, the publicly available IDL code and our independently developed C++ code. We conclude in sect. 6. ",
        "conclusions": "The shapelet technique provides a very powerful tool to describe intrinsically smooth, compact objects. It incorporates the effects of noise such that it is able to extract all significant physical information from the object. Quantities like the flux or the ellipticity can then be computed efficiently in the much reduced shapelet space.  However, the trust in the estimators of the physical properties relies on the assumption that there is a single decomposition result. Since the shapelet decomposition depends on four external parameters (the scale size $\\beta$, the maximum decomposition order $n_{max}$ and the components of the centroid $\\mathbf{x_c}$), it is inevitable to choose appropriate parameters. Unfortunately, it is not sufficient for finding a reliable parameter set that the reconstruction has residuals at noise level, i.e. $\\chi^2\\approx 1$. In fact, increasing $n_{max}$ makes it more and more likely to find a continuous and growing range of parameters obeying $\\chi^2 \\leq 1$. This degeneracy has to be broken by ensuring that both conditions \\eqref{eq:chi2_minimum} and \\eqref{eq:chi2_converged} are fulfilled. If this is not the case, deviations of the shapelet parameters from the optimal values will introduce arbitrary errors on the shapelet coefficients and therefore on all quantities derived from them, even though $\\chi^2 \\leq 1$.  In addition, the sampled shapelet basis vectors have to remain orthonormal to guarantee the correctness of the decomposition results. The orthonormality can be violated by undersampling, boundary truncation and introduction of non-orthogonal elements. There is no way to avoid undersampling, but the use of suitably large frames and the subtraction of noise before the decomposition can remedy the other reasons for non-orthonormality.  We showed that the proposed solutions do indeed result in a faithful representation of the decomposed objects, not only at the visual level but also concerning their physical properties. By comparing the results of our C++ implementation with the publicly available IDL code, we found that our code is both more reliable and has considerable performance benefits."
    },
    "0608/astro-ph0608533_arXiv.txt": {
        "abstract": "X-rays provide one of the few bands through which we can study the epoch of reionization, when the first galaxies, black holes and stars were born. To reach the sensitivity required to image these first discrete objects in the universe needs a major advance in X-ray optics.  {\\em Generation-X} (Gen-X) is currently the only X-ray astronomy mission concept that addresses this goal. {\\em Gen-X} aims to improve substantially on the {\\em Chandra} angular resolution and to do so with substantially larger effective area. These two goals can only be met if a mirror technology can be developed that yields high angular resolution at much lower mass/unit area than the {\\em Chandra} optics, matching that of {\\em Constellation-X} ({\\em Con-X}).  We describe an approach to this goal based on {\\em active X-ray optics} that correct the mid-frequency departures from an ideal Wolter optic {\\em on-orbit}. We concentrate on the problems of sensing figure errors, calculating the corrections required, and applying those corrections. The time needed to make this in-flight calibration is reasonable. A laboratory version of these optics has already been developed by others and is successfully operating at synchrotron light sources. With only a moderate investment in these optics the goals of {\\em Gen-X} resolution can be realized. ",
        "introduction": "\\label{sect:intro}  %  The high angular resolution of the {\\em Chandra} X-ray Observatory has revolutionized X-ray astronomy, and indeed wide areas of astrophysics as a whole.  At some point though, a larger and higher angular resolution successor to {\\em Chandra} must be built (Fabbiano 1990, Elvis \\& Fabbiano 1996). Fortunately X-ray astronomy is nowhere near its physical limits: the Chandra mirror would have been diffraction limited at $\\sim$20~milli-arcsec (van Speybroeck 2000, private communication), while the thermal line widths of 1~keV temperature plasmas are of order 100~km~s$^{-1}$, so the 'thermal limit' to X-ray spectroscopy is $\\lambda/\\Delta\\lambda \\sim$6000 (Elvis 2001).  More immediately, Chandra is a small telescope: The median exposure time for {\\em Chandra} observations is a day (85~ksec), and many of the breakthrough observations needed 1~Msec (Table~1).  This should not be surprising - the large number of independent spatial and spectral bins {\\em Chandra} offers ($\\sim$10$^5$ in ACIS-I) requires large number of photons to fill them, but {\\em Chandra} has an area of only $\\sim$800~cm$^{-2}$, equivalent to a 30~cm diameter optical telescope. As a year contains only $\\sim$20~Msec of net observing time, even at the high observing efficiency of {\\em Chandra}, only a small number of these Megasecond-class programs can be carried out in a year, and only a limited number in the whole {\\em Chandra} lifetime. A reconaissance of the X-ray sky at high resolution requires an faster observation rate.  \\begin{table}[t] \\caption{Some {\\em {\\em Chandra}} Breakthroughs} \\begin{tabular}{|lccll|} \\hline {\\em {\\em Chandra}} Observations&&&Implication& reference\\\\ \\hline X-ray background resolved to 10~keV&Ms&$\\rightarrow$& accretion luminosity of the Universe&Giacconi et al. 2001\\\\ Dark Energy measured &&$\\rightarrow$& new constraints on new physics& Allen et al. 2004\\\\ No cooling flows in clusters&&$\\rightarrow$&20-year puzzle solved; AGN feedback &David et al., 2001\\\\ Cooling fronts in clusters&&$\\rightarrow$&build up of galaxies, clusters &Vikhlinin et al. 2001 \\\\ Warm-Hot Intergalactic medium&Ms&$\\rightarrow$ &'missing baryons' found& Nicastro et al. 2002 \\\\ Multi-phase Winds from AGNs &Ms&$\\rightarrow$&AGN 'feedback' to galaxy formation& Kaspi et al. 2002\\\\ Ultra-Luminous X-ray Sources&& $\\rightarrow$& intermediate mass black holes?& Fabbiano et al. 2003\\\\ Galactic center flares& &$\\rightarrow$ & Not so quiet Black Hole & Baganoff et al. 2001\\\\ Hot ISM abundances&Ms&$\\rightarrow$ &SN yields, ecology& Fabbiano et al. 2004 \\\\ Optically dull/X-ray bright galaxies&Ms&$\\rightarrow$&hidden active galaxies & Alexander et al. 2003 \\\\ 0.5$c$ wisps in Crab&&$\\rightarrow$&particle acceleration to near $c$& Hester et al. 2002 \\\\ \\hline \\end{tabular} Ms: exposure times approaching 1~Msec (2~weeks) needed to make the discovery. \\end{table}  In terms of angular resolution, {\\em Chandra} often provides just one unique example of each class of object: e.g. the Crab nebula wisps, the Antennae galaxy merger system, the 'bullet' cluster 1E~0657-06 (Markevitch et al., 2004) Only by increasing angular resolution can we expand these samples: a factor 10 improvement increases the accessible volume 100 times for Galactic objects (as they lie in a disk), and a 1000 times for extragalactic objects (and for objects within 100~pc).  At the faintest fluxes, where the first generation objects will be found, high angular resolution is required to make firm identifications with counterparts from JWST, ALMA or other next generation telescopes.  The need then is clear. To give some substance to the science goals we first describe briefly the first generation of black holes and their expected properties. Then we outline the {\\em Gen-X} Vision Mission study - a just completed 2-year NASA funded investigation. We then address the question of active X-ray optics in some detail, to show that this is a feasible prospect. ",
        "conclusions": "The future of X-ray astronomy needs a high angular resolution successor to {\\em Chandra}, with larger area. Such a mission would address a huge range of astrophysics, and in particular would open up one of the few spectral windows to the study of the very first stars, galaxies, and black holes.  Mission architecture studies show that such a mission is feasible, despite the long focal lengths (50~m - 125~m) and large diameters (6~m - 20~m). The tall pole is the development of lightweight high resolution X-ray optics.  Active X-ray optics, controlled with piezoelectric acuators, address the biggest technical challenge faced by the {\\em Generation-X} Vision Mission concept. Piezos are a good match to the thin shells required by light-weight optics and to the need not to block the narrow optical path through densely nested Wolter optics.  A 0.1~arcsec HPD image at 1~keV can be produced with 2~cm sized actuators.  Out-of-focus imaging of the converging beam will allow figure errors on the individual patches on each sector of each shell to be measured independently, and celestial sources provide a sufficient count rate to make the measurements in a reasonable time. The parallization of the problem afforded by the out-of-focus technique reduces the compute time to a modest level. Meanwhile, independent work at synchrotrons has shown that 2~cm sized (1-D) actuators can be constructed, and work well and reliably, reducing figure errors in a Kirkpatrick-Baez configuration by large factors.  The {\\em Generation-X} Vision Mission study has shown that the obstacles to creating large high resolution X-ray optics to go well beyond the capabilities of {\\em Chandra} are far smaller than might have been thought. It has not escaped our attention that once high resolution, light-weight, X-ray optics are available, then they would be the mirrors of choice for most, if not all, X-ray astronomy missions. We believe that the rapid development of active X-ray optics by means of a quite moderate sized program, is both an urgent and a reasonable short-term goal for X-ray astronomy."
    },
    "0608/astro-ph0608019_arXiv.txt": {
        "abstract": "{{\\bf ABSTRACT} The Millennium Run is the largest simulation of the formation of structure within the $\\Lambda$CDM cosmogony so far carried out. It uses $10^{10}$ particles to follow the dark matter distribution in a cubic region 500$h^{-1}$Mpc on a side, and has a spatial resolution of 5 $h^{-1}$kpc. Application of simplified modelling techniques to the stored output of this calculation allows the formation and evolution of the $\\sim 10^7$ galaxies more luminous than the Small Magellanic Cloud to be simulated for a variety of assumptions about the detailed physics involved.  As part of the activities of the German Astrophysical Virtual Observatory we have used a relational database to store the detailed assembly histories both of all the haloes and subhaloes resolved by the simulation, and of all the galaxies that form within these structures for two independent models of the galaxy formation physics. We have created web applications that allow users to query these databases remotely using the standard Structured Query Language (SQL). This allows easy access to all properties of the galaxies and halos, as well as to the spatial and temporal relations between them and their environment. Information is output in table format compatible with standard Virtual Observatory tools and protocols. With this announcement we are making these structures fully accessible to all users. Interested scientists can learn SQL, gain familiarity with the database design and test queries on a small, openly accessible version of the Millennium Run (with volume 1/512 that of the full simulation). They can then request accounts to run similar queries on the databases for the full simulations. } \\newpage ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608705_arXiv.txt": {
        "abstract": "We present colour-magnitude diagrams of open clusters, located in the range $112^\\circ < l < 252^\\circ$, manifesting stellar populations in the background of clusters. Some of the populations are found to be located beyond the Perseus arm and may be the part of Norma-Cygnus (outer) arm. The outer arm seems to be continued from $l\\sim120^\\circ$ to $l\\sim235^\\circ$. The background populations follow the downward warp of the Galactic plane around $l\\sim240^\\circ$. ",
        "introduction": "\\begin{figure*} \\includegraphics[height=6.6in,width=6.6in]{fig1.eps} \\caption{Upper and lower panels of the figure show two-colour diagrams and CMDs for the six clusters studied by us. Thin and thick lines represent ZAMS shifted to match the observed sequence of cluster and background population respectively (see  text for details). Open circles represent probable background population members.} \\end{figure*}  \\begin{figure*} \\includegraphics[height=6.6in,width=6.6in]{fig2.eps} \\caption{Same as Fig. 1 but for data taken from literature. In the case of NGC 129 the background population could not be disentangled. Upper right panel of the figure shows CMDs for the clusters NGC 1912, and NGC 2099 (see  text for details). Isochrone by Bertelli et al. (1994) for log age = 7.4 is also shown in the CMD of cluster NGC 581.} \\end{figure*}  \\begin{figure*} \\includegraphics[height=4in,width=7.0in]{fig3.eps} \\caption{Same as Fig. 1 but for nearby field regions to the clusters } \\end{figure*}  The recently discovered ring-like structure (the Monoceros/Canis Major Ring) in the anti-center direction of the Galaxy (e.g. Newberg et al. 2002, Ibata et al. 2003) could be the consequence of accretion of dwarf galaxies by the Milky way (cf. Martin et al. 2004, Bellazzini et al. 2006 and references therein). Helmi et al. (2003) suggested that the ring may be a tidal arc produced by stars stripped away from the parent satellite galaxy during a recent interaction. This would produce an asymmetric structure above and below the disc, limited in Galactic longitude and with a significant velocity gradient in the Galactic longitude direction. If the ring is symmetric without velocity gradient, it could be produced by ancient minor mergers.  From the analysis of the asymmetries in the population of Galactic M-giant stars, Martin et al. (2004) have reported detection of the Canis Major dwarf galaxy (CMa) located at 7-8 kpc from the Sun and centered at $l\\sim240^\\circ$ and $b\\sim -8^\\circ$. CMa was identified as strong elliptical-shaped over-density of M-giants by the comparison of star counts in Northern and Southern Galactic hemispheres. The $V/B-V$ colour magnitude diagram (CMD) of a field located at $\\simeq4^\\circ.2$ from the center of the CMa by Bellazzini et al. (2004) revealed a population of an intermediate/old ($age\\sim4-10$ Gyr) and moderately metal deficient stellar system located a distance of $d_\\odot=8.0\\pm1.2$ kpc. These results have also been supported by the CMD presented by Martinez-Delgado et al. (2005). A Blue Plume (BP) of possibly young stars or blue stragglers has also been detected in the CMD. Martin et al. (2005) further reported existence of CMa on the basis of distance-radial velocity gradient among CMa stars that can be explained as the effect of on-going tidal disruption of the stellar system.  \\begin{figure} \\includegraphics[height=3.5in,width=3.0in]{fig4.eps} \\caption{Upper and lower panels show TCD and CMD respectively for Do 25 cluster region. Thin and thick lines represent ZAMS shifted to match the observed sequence of foreground and cluster stars respectively. } \\end{figure}  Momany et al. (2004, 2006) suggested that the CMa over-density is an effect of the Galactic warp. The Southern stellar maximum warp occurring near $l\\sim240^\\circ$ (for $R_\\odot \\sim7$ kpc) brings down the Milky Way mid-plane by $\\sim3^\\circ$ in this direction. The regularity and consistency of the stellar, gaseous and dust warp argues strongly against a recent merger scenario for CMa. They present evidence to conclude that all observed parameters of CMa are consistent with it being a normal Milky Way outer-disk population.  Carraro et al. (2005) and Moitinho et al. (2006) have also detected BP population (age $\\le100$ Myr) in the background of several open clusters in the third Galactic quadrant. They concluded that BP population is associated to Norma-Cygnus spiral arm and the over-density towards CMa is simply a projection effect of looking along the nearby local arm.  The detection of the CMa over-density by Martin et al. (2004) has attracted the attention towards the third Galactic quadrant of the Milky Way. Young open clusters are one of the best tools to study the spiral arms and structure of the Milky Way (Becker \\& Fenkart 1970; Pandey et al. 1988, 1990). Relevant parameters like their distance, age, reddening towards the cluster can be measured more accurately than for single stars. We have embarked on a wide field CCD photometric survey around open clusters using the Kiso Schmidt telescope (Japan), since extensive studies of the coronal regions of open clusters have not been carried out so far mainly because of non-availability of photometry in a large field around open star clusters. We have undertaken a $CCD$ photometric study of a large area around open clusters using the $2K \\times 2K$ $CCD$ mounted on the 105-cm Schmidt telescope of the Kiso Observatory which gives $\\sim 50'\\times 50'$ field. Results on few open clusters have been reported by Pandey et al. (2001, 2005, 2006) and Sharma et al. (2006). We also noticed undesired background contamination in the CMDs of open clusters located in the second and third galactic quadrant, therefore in the light of above discussions it is considered worthwhile to study the background population towards the target clusters. ",
        "conclusions": ""
    },
    "0608/astro-ph0608527_arXiv.txt": {
        "abstract": "We have obtained  the  spectrum of a middle-aged PSR B0656+14  in the  4300-9000\\AA\\ range with the ESO/VLT/FORS2. Preliminary results show that at 4600-7000\\AA\\  the spectrum is almost featureless and flat with a spectral index   $\\alpha_{\\nu}$$\\simeq$-0.2  that undergoes a change  to a positive value at longer wavelengths. Combining with available multiwavelength data suggests two wide, red and blue, flux depressions whose frequency  ratio is about 2 and which could be the 1st and 2nd harmonics of  electron/positron cyclotron absorption formed at magnetic fields $\\sim$10$^8$~G  in  upper magnetosphere of the pulsar. ",
        "introduction": "\\label{intro} Multiwavelength observations of radio pulsars are an important tool for the study of not yet clearly understood radiative mechanisms and spectral evolution of rotationpowered isolated neutron stars (NSs). Optical observations are an essential part of these studies. Among a dozen  optically identified   NSs,  only seven pulsars have parallax based distances and, therefore,  minimal uncertainties ($\\le$10\\%)  in luminosities.  Their optical spectra are shown in Fig.~1. A reliable optical spectrum has been obtained  only for the young and bright Crab pulsar \\citep{Sollerman}. \\begin{figure} \\centering \\includegraphics[width=7cm,clip=]{zharikovfig1.eps} \\caption{Optical spectra of seven pulsars of different characteristic age indicated in the left-bottom corner of each panel in years. The youngest Crab pulsar is at the top and the oldest PSR B0950-08 is at the bottom. % For PSR B0656+14 dashed lines show the low energy extensions of the blackbody (BB, T=0.84MK) and power law (PL, $\\alpha =0.45$) X-ray spectral components and their sum \\citep{Koptsevich}. } \\label{fig:1}       % \\end{figure} Other pulsars are fainter and are represented mainly by broadband photometric points. Published  optical spectra  of PSR B0540-69 % \\citep{Hill, Serafim} are strongly contaminated  by a bright pulsar nebula \\citep{Serafim}, while a tentative spectrum of Geminga \\citep{Martin} is much noisier than available photometric fluxes. \\cite{Mignani2} reported on the Vela pulsar spectral observations  but these data are not published yet.  In Figure 2 we show   the evolution of luminosity L and radiation efficiency ${\\rm\\eta=L/L_{sd}}$ (${\\rm L_{sd}}$ is spindown luminosity) demonstrated by these seven pulsars % using the data of Table~1 from \\cite{Zhar2}. We note significantly non-monotonic dependencies of ${\\rm\\eta_{Opt}}$ % and ${\\rm\\eta_{X}}$  % versus pulsar age with a pronounced  minimum at the beginning of the middle-age epoch (${\\rm \\simeq10^4}$~yr) and  comparably higher  efficiencies of  younger and older pulsars.  \\begin{figure}[t] \\centering \\includegraphics[width=6.5cm,bb=0 0 560 955,clip=]{zharikovfig2.eps} \\caption{{\\sl From top to bottom:} Evolution  of the spindown, radio, optical, X-ray, and $\\gamma$-ray luminosities and respective efficiencies  % demonstrated  by the 7 optical pulsars from Fig.~1.} \\label{fig:1}       % \\end{figure} Owing to its relative brightness and proximity, the middle-aged PSR B0656+14 is one of isolated NSs most intensively studied in different wavelengths. It was discovered in radio % by \\cite{Manchester}, then identified in X-rays with Einstein, and observed in details with ROSAT,  ASCA, Chandra and XMM (for references see \\cite{Shib1, Shib2}). % The X-ray emission can be described as a combination of thermal radiation from the entire surface of a cooling NS and  from hotter % polar caps heated by relativistic particles  of magnetospheric origin. % An excess over the hot thermal component at energies $\\ge$2 keV was interpreted as nonthermal radiation from the pulsar magnetosphere \\citep{Greiv}. The pulsar has  been also marginally detected in $\\gamma$-rays  ($\\ge$50  MeV) by % \\cite{Raman}.   In the optical  PSR B0656+14 was identified by % \\cite{caraveo} with the ESO/NTT telescopes in the V band. % It was then studied in UV with the HST/FOC in the F130LP, % F430W, F342W and F195W bands \\citep{Pavlov1}  % and  with the  HST/WFPC in  the  F555W band % \\citep{Mignani1}.  % Detailed  photometric   studies in the optical-NIR were performed by  % \\cite{Kurt}, \\cite{Koptsevich}, \\cite{Komarova}, and \\cite{Shib2}. The studies % showed that the bulk of the optical radiation % is of  nonthermal origin. This was confirmed by the detection of  coherent optical pulsations with the radio pulsar period  % in the B band \\citep{Shearer}, in a wide 400-600 nm passband % \\citep{Kern} and in NUV  % \\citep{Shib1}. The  % pulse profiles are  rather sharp % with a high pulse fraction % as expected  for  nonthermal emission mechanisms. \\begin{figure} \\centering \\includegraphics[width=8cm,bb=0 0 760 403,clip=]{zharikovfig3.eps} \\caption{Relationship between the B-band optical and 2-10 keV X-ray efficiencies  % for the same 7 pulsars as in Fig.~1,2.} \\label{fig:1}       % \\end{figure} The phase integrated multiwavelength spectrum (Fig.~4)   shows  % \\citep{Koptsevich,% Shib2} that the NIR-optical-UV spectral energy distribution is,  in a first approach, compatible with the low energy extension of the sum  of the X-ray thermal blackbody (BB) spectral component from the whole NS surface  and the power law (PL) component  dominating in the high energy tail. % The  BB extension does not contribute  at longer  % wavelengths where  the optical-NIR fluxes are in a good agreement with the PL  alone (Fig.~1). This indicates  a common origin of the nonthermal optical and X-ray emission, which is strongly supported by % a good coincidence in phase and shape of % the pulse profiles in the optical and  the X-ray tail % \\citep{Shib1}. The  same origin of the nonthermal optical and X-ray photons is likely to be  a general property for other pulsars detected in both ranges, as follows from a strong correlation  between respective efficiencies (Fig.~3) found by \\cite{Zhar1,Zhar2}.    In the optical emission of PSR B0656+14 there is  an apparent, $\\simeq$(3-5)$\\sigma$, flux excess  over the PL  % ``continuum\" at Log$(\\nu)\\simeq$14.7 % (Fig.~1). This could reveal an additional, 3rd, spectral component to the  BB+PL discussed above. Here we present first results of the optical spectroscopy of the pulsar partially motivated by more detailed studies of the excess. In a broader sense, these results also allow us to consider, for the first time, optical properties of a middle-aged ordinary pulsar % at the spectroscopic level, as it has been achieved so far only for the Crab pulsar.   \\begin{figure*} \\centering \\includegraphics[width=15.cm,clip=]{zharikovfig4.eps} \\caption{{\\sl Left:} Unabsorbed multiwavelength spectrum of  PSR B0656+14 from the radio through $\\gamma$-rays. The box marks the % range zoomed in the right panel. {\\sl Right:} The spectral and photometric fluxes of PSR B0656+14 in the NIR-optical-NUV range  obtained with different telescopes and instruments, as notified in the plot. Black dashed lines show low energy extensions of the soft blackbody (BB) and  power law (PL) spectral fits and their sum (BB+PL) obtained in X-rays. The black solid lines show the same but with the PL normalization shifted up by a factor of 1.4 (PL1) to fit the upper edge of its 1$\\sigma$ error bar shown at the left side of the plot. These are possibly a better match for the optical and NUV spectra than the dashed lines. The symbol $\\oplus$ marks the Earth atmospheric absorption band  near 7600\\AA. } \\label{fig:1}       % \\end{figure*} ",
        "conclusions": "\\label{sec:3} The binned and dereddened with ${\\rm A_V}$=0.093 spectrum of the pulsar is shown in Figure~4 (red curve in the right panel). We  show also   % available multiwavelength  data (see \\cite{Shib2} for the data and  ${\\rm A_V}$ % description). The spectrum is  in a good agreement  with the % broadband VRI fluxes, while it is somewhat higher than the B band flux.  % It does not show any strong nebular emission lines that could be responsible for the apparent VR  excess,  % mentioned  above as a 3rd component, while the presence of  weak % features can not be completely ruled out. % The continuum of the dereddened spectrum in 4600-7000\\AA\\ range has a  power law shape $\\propto\\nu^{-0.2\\pm0.2}$ (green line) % in agreement  with the nonthermal nature of the bulk of the optical emission. Its slope is close to the value expected in this range from   the BB+PL   extended from X-rays. % The slope likely undergoes a change to a positive value between the R and I bands, as follows also from the photometric data.  Within uncertainties   the bluest end of the optical and  the reddest end of the NUV \\citep{Shib1} spectra are compatible with each other, suggesting a smooth connection of  both. However, the F430W photometric flux in the gap between them % drops below this connection with a significance of about 2$\\sigma$  of the flux confidence level. Unless this is a result of some unknown  systematics in calibration, it suggests a spectral dip in the pulsar emission centered near  $\\simeq$4300\\AA\\ ($\\simeq$14.83 in Log($\\nu$)).  The  optical and NUV spectra and  two NIR  photometric points, F160W and F187W, almost perfectly match the BB+PL  extension % from X-rays, if the PL normalization is taken to be a factor of 1.4 higher (solid lines) than its best X-ray fit value (dashed lines). The change is  within 1$\\sigma$ uncertainty of the  fit. Considering the solid line version % as a new optical ``continuum level\" we find an additional  and   more significant  flux depression in the red part of the  spectrum overlapping  the I  and  F110W  bands and centered near  9000-10000\\AA\\ (Log($\\nu)$$\\simeq$14.5).  Additional spectral studies are necessary to confirm the suggested ``red\" and ``blue\" features  and  to measure   their shapes and wavelengths more accurately. % Nevertheless, if they are real, an  approximate   blue-to-red frequency ratio is $\\simeq$2. This indicates that they can be  the  1st and  2nd  harmonics of an electron/positron cyclotron absorption formed in the upper magnetosphere of the pulsar at  an effective altitude % where the magnetic field B$\\simeq$10$^{8}$~G. This is $\\simeq$360 km, assuming a dipole NS field with a surface value of 4.66$\\times$10$^{12}$ G, as derived from spindown measurements. The absorbing $e^{\\pm}$  % have to be cooled enough  and provide a sufficient optical depth above a source of the nonthermal continuum. The source altitude  is, therefore,  $<$360 km and much below  the light cylinder radius of $\\simeq$18$\\times$10$^{3}$ km, likely suggesting its polar cap origin. The  features are  broad, as expected  from % the magnetospheric field inhomogeneity. Tentative ($\\leq$2$\\sigma$)  absorption features  in the NUV spectrum  at Log($\\nu)\\simeq$15 and  15.1  \\citep{Shib1} may be % the 3rd and 4th % harmonics, respectively, which are fainter as the cyclotron harmonic % intensity  decreases with its number ${n}$.  Similar, albeit less significant, spectral features  are likely seen in  the photometric and spectral data of another middle-aged pulsar Geminga (Fig.1). The absence of strong nebular % lines suggests  that the features are also of the NS magnetospheric origin. To explain the Geminga spectrum \\cite{Martin} applied  % a toy model of an ion cyclotron absorption at B$\\simeq$10$^{11}$~G  in the inner magnetosphere of the NS combined with the BB and PL  components. At this  field a low ${n}$ ion cyclotron frequencies  indeed fall in the optical range. However, it is likely that the nonthermal optical emission is generated  in the upper magnetosphere  where the magnetic field is by orders of magnitude weaker  and any  ion cyclotron absorption of the respective optical continuum is negligible. In this case the electron/positron cyclotron absorption or scattering appears to be more plausible interpretation.  The optical spectrum of the young Crab-pulsar is featureless and has a different  (positive) slope (Fig.~1). The ten times older Vela-pulsar has a flat and also featureless spectrum  (Mignani, private communication). The spectroscopy of  PSR B0656+14 demonstrates that  optical spectra of middle-aged pulsars can be  % distinct from those of  younger ones by the presence of unusual spectral features  or slope changes. To study  this  new spectral observations  of  PSR B0656+14 in the NIR and 3000-5000\\AA\\ ranges are needed. A question, whether the observed  difference in pulsar spectra is simply caused by different pulsar  geometry or by a change of  physical conditions in the emission region  with  age, demands   quantitative modeling  physical processes in pulsar magnetospheres."
    },
    "0608/astro-ph0608657_arXiv.txt": {
        "abstract": "The optical rest--frame sizes of 10 of the most massive ($\\sim$5$\\times$10$^{11}$h$_{70}$$^{-2}$M$_{\\sun}$)  galaxies found in the near--infrared MUNICS survey at 1.2$<$z$<$1.7 are analysed. Sizes were estimated both in the J and K' filters. These massive galaxies are at least a factor of 4$_{-1.0}^{+1.9}$ ($\\pm$1 $\\sigma$) smaller in the rest--frame V--band than local counterparts of the same stellar mass. Consequently, the stellar mass density of these objects is (at least) 60 times larger than massive ellipticals today. Although the stellar populations of these objects are passively fading, their structural properties are rapidly changing since that redshift. This observational fact disagrees with a scenario where the more massive and passive galaxies are fully assembled at z$\\sim$1.4 (i.e. a monolithic scenario) and  points towards a dry merger scenario as the responsible mechanism for the subsequent evolution of these galaxies. ",
        "introduction": "In the local universe, the population of galaxies with stellar masses larger than 10$^{11}$M$_{\\sun}$ is dominated by passive early--type galaxies that are a factor of $\\sim$3 more numerous than late--type galaxies above this mass threshold (Baldry et al 2004).  Their stellar populations are old, metal rich and are characterized by a short formation time-scale (e.g. Heavens et al. 2004; Thomas et al 2005; Feulner et al. 2005). In addition, these massive galaxies are large, with sizes (as parametrized by the effective radius) larger than 4 kpc (Shen et al. 2003).  Historically, two different formation scenarios have been proposed in order to explain the properties of these objects: the so--called monolithic collapse model (Eggen, Lynden--Bell \\& Sandage 1962; Larson 1975; Arimoto \\& Yoshii 1987; Bressan, Chiosi \\& Fagotto 1994) and the hierarchical merging model (White \\& Frenk 1991). In the former scenario, spheroidal galaxies formed at a very early epoch as a result of a global starburst, and then passively evolve to the present. In the merger model, spheroids are formed by violent relaxation during major merger events.  Favoring the monolithic model is the fact that the bulk of stars in massive ellipticals are old (Mannuci et al. 2001) and have high [$\\alpha$/Fe] ratios (i.e. short star formation time scales; Worthey, Faber \\& Gonzalez 1992). On the other hand, supporting a hierarchical merger scenario, current observations seem to find a decline in the number of massive galaxies at high--z. The space density of red passively evolving early--type galaxies has moderately increased since z$\\sim$1 (Daddi et al. 2000; Pozzetti et al. 2003; Bell et al. 2004; Drory et al. 2004; Faber et al. 2005). At even higher redshift, z$\\sim$1.7, their space density appears to be a factor of 2--3 smaller than that of their local counterparts (Daddi et al. 2005; Saracco et al. 2005; Drory et al. 2005). In addition, new generations of semianalytical models (e.g. de Lucia 2006) are now able to produce results in better agreement with the stellar population properties of ellipticals.  An additional test to check the validity of the above two scenarios is to explore the size evolution of the spheroids through  cosmic time. In this sense, the prediction from the monolithic model is that the stellar mass--size relation of these objects should remain unchanged after their formation, with the luminosity--size relation evolving in agreement with the fading of their stellar populations. In the hierarchical model, however, the stellar mass--size relation of these objects should change as a result of the increase in size of the remnants after each merger.  From the observational point of view,  the evolution of the luminosity--size and stellar mass--size relations of the early--type galaxies since z$\\sim$1 is consistent with the passive aging of ancient stellar populations (Trujillo \\& Aguerri 2004, McIntosh et al. 2005). Consequently, the structural properties of the more massive galaxies   seem  to not change from z$\\sim$1 to the present. At z$>$1 the situation is less clear. Large and deep near--infrared surveys are needed to collect a significative sample of massive passively evolving  galaxies at z$>$1 and to explore the sizes of the galaxies in their optical rest--frame. Using the Faint Infrared Extragalactic Survey (FIRES; Franx et al. 2000), Trujillo et al. (2004;2006) have shown that there is a hint for most massive galaxies (M$_{\\star}$$\\gtrsim$7$\\times$10$^{10}$ h$_{70}$$^{-2}$M$_\\odot$) at z$\\sim$2.5 being a factor of 2 smaller than present--day counterparts with similar masses. However, the number of galaxies in that sample is not large enough to firmly establish this result. In addition, there has been a recent claim of 4 very compact (r$_e$$\\lesssim$1 kpc in their local UV restframe) and massive (M$_{\\star}$$>$10$^{11}$ h$_{70}$$^{-2}$M$_\\odot$) passively evolving galaxies at z$\\sim$1.7 in the UDF (Daddi et al. 2005). The few number of objects together with the fact that some of these galaxies could be hosting an AGN make the above claims uncertain.  To shed some light on the above issue, we estimate the sizes (in their local optical rest--frame) of a sample of 10 very massive (10$^{11}$$<$M$_{\\star}$$<$10$^{12}$M$_{\\sun}$)  galaxies spectroscopically classified as early--type galaxies at 1.2$<$z$<$1.7, with no sign of AGN activity in their spectra (except for one object). These objects are consequently good candidates to test whether their sizes are equal or smaller than their local counterparts.  Throughout, we will assume a flat $\\Lambda$--dominated cosmology ($\\Omega_M$=0.3, $\\Omega_{\\Lambda}$=0.7 and $H_0$=70  km s$^{-1}$ Mpc$^{-1}$). All magnitudes are provided in the Vega system unless otherwise stated.  \\vspace{-0.7cm} \\section[]{Description of the data}  The near-infrared images used in this study  were taken from the Munich Near-IR Cluster Survey (MUNICS; Drory et al. 2001),  and its deeper follow-up project called MUNICS-Deep. MUNICS is a wide-field medium-deep photometric survey taken in the near-infrared and optical filters. Dedicated follow-up spectroscopy is available for a selected sub-sample (Feulner et al. 2003). The main part of the survey consists of 10 fields with a total area of $\\sim 0.3$~deg$^2$. For all these fields photometry in $K'$, $J$, $I$, $R$, $V$, and $B$ is available, with limiting magnitudes ranging from K' $\\simeq$ 19.5 and J $\\simeq$ 21 to B $\\simeq$ 24.0 mag (50\\% completeness for point sources; Snigula et al. 2002). This is sufficiently deep to detect passively evolving systems up to a redshift of z$\\lesssim$1.4, and at a luminosity of 0.5L$^{*}$. The final $K'$-selected catalog contains roughly 5000 objects and is described  in Drory et al. (2001).  The sample of massive Extremely Red Objects (ERO) studied here was selected in three survey fields (S2F1, S2F5 and S7F5). The primary selection criteria were $K'<18.5$ and $R-K'>5$, resulting in a list of 36 objects. Low-resolution near-infrared spectroscopy was carried out for  $\\sim 60$\\% of them, and ten objects are identified as early-type galaxies at  $1.2<z<1.7$, with stellar masses well exceeding $10^{11} M_\\odot$ (Saracco et al. 2003, 2005; Longhetti et al. 2005).  Seven out of the ten EROs explored here are located in the field S2F1, for which deeper near-infrared images are available from the MUNICS-Deep survey (Goranova et al., in prep.). MUNICS-Deep aims at obtaining a contiguous 1-square-degree field (overlapping the MUNICS patchs S2F1 and S2F5) in optical and near-infrared filters to a detection limit 2~mag deeper than MUNICS. To improve the size measurements of these seven EROs in the field S2F1, the deep $K'$ and $J$-band images from MUNICS-Deep were used. These were obtained with Omega2000 at the Calar Alto 3.5-m telescope at a pixel scale of 0.45 arcsec/pixel, a typical seeing of $\\lesssim$1.0 arcsec, and limiting magnitudes of K'$\\sim$21.5 and J$\\sim$23.5 mag (again 50\\% completeness limits for point sources). Basic data reduction (Goranova et al., in prep.) was performed using a modified version of the \\texttt{IRAF} external package  \\texttt{XDIMSUM}.  For the remaining three objects, the sizes were estimated on the $K'$-band and J-band images of the (shallower) MUNICS project. These images, taken with OmegaPrime at the Calar Alto 3.5-m telescope, have a pixel scale of 0.396 arcsec/pixel, a typical seeing of $\\sim$1.2 arcsec, and a limiting magnitude of K'$\\sim$ 19.5.  Stellar masses have been derived from the K'-band absolute magnitudes by means of the mass--to--light ratio M/L$_{K'}$ derived from the best fitting models (see details in Longhetti et al. 2005). The largest uncertainty in the stellar mass computation comes from the variation of M/L according to the age of the stellar population and the adopted IMF. However, it is worth noting that, given the extremely bright K'-band  magnitudes (K'$<$18.4) and the redshifts (z$>$1.2) of our galaxies, their resulting stellar masses are well in excess of 10$^{11}$ M$_{\\sun}$ leaving aside any model assumption.  In this paper we use stellar masses derived using a Kroupa IMF. However, to estimate the uncertainty in the stellar masses we used a large set of different IMFs (Longhetti et al. 2005).  \\vspace{-0.7cm} \\section[]{Size estimation}   To estimate the sizes of the galaxies we have used the GALFIT code (Peng et al. 2002). GALFIT convolves S\\'ersic (1968) r$^{1/n}$ profile galaxy models with the point-spread function (PSF) of the images and then determines the best fit by comparing the convolved models with the science data using a Levenberg--Marquardt algorithm to minimize the $\\chi^2$ of the fit.  The spatial resolution of our images does not allow to estimate accurately the shape (index $n$) of the surface brightness profiles. For that reason, and to decrease the number of free parameters in our fits, we have calculated the size of the galaxies by fixing the S\\'ersic index to n=1 (i.e. an exponential profile) and n=4 (i.e. a de Vaucouleurs profile). Both models are convolved with the image PSF. The effective radii provided by every fit (r$_{e,1}$ and r$_{e,4}$) are used to estimate a mean effective radius and indicate the range of variation of the sizes of our galaxies. The PSF that was used for every galaxy corresponds to the nearest  (bright enough but non-saturated) star to the galaxy.  Neighboring galaxies were excluded from each model fit using a mask, but in the case of closely neighboring galaxies with overlapping isophotes, the galaxies were fitted simultaneously.  \\vspace{-0.35cm} \\subsection{Testing the size estimates: simulations}  The results presented in this paper rely on our ability to measure accurate structural parameters. To gauge the accuracy of our size determination we have created 250 artificial galaxies uniformly generated at random in the following ranges: 18$\\leq$ J$\\leq$21, 0.$''$1$\\leq$r$_e$$\\leq$1.6$''$ (i.e. 0.8$\\lesssim$r$_e$$\\lesssim$13.5 h$_{70}^{-1}$ kpc in the local restframe at z$\\sim$1.4) and 0.5$\\leq$n$\\leq$8. Simulations were done in J band only, but our results can be extrapolated to the K band data because of their similar signal--to--noise quality. The mock galaxies span a large range of surface brightness shapes (i.e. they are not restricted to n=1 or n=4) to model the different galaxy profiles found in the observations (Trujillo et al. 2006). To simulate the real conditions of our observations, we add a background sky image (free of sources) taken from a piece of the MUNICS--Deep field  image in the J band. Finally, the galaxy models used (n=1 and n=4) were convolved with the observed PSF. The same procedure was used to retrieve the structural parameters both in the simulated and actual images.  Fig. \\ref{simulation} shows the comparison between the input and recovered size values in our simulations. The recovered size value, r$_{e,out}$ is evaluated as the mean value between the size recovered using n fixed to 1, and n fixed to 4: r$_{e,out}$=(r$_{e,out,n=1}$+r$_{e,out,n=4}$)/2. For that reason, the intrinsic scatter shown in Fig. \\ref{simulation} is independent of the magnitudes and mainly caused by the fact that the index $n$ of the model galaxies is fixed whereas the mock galaxies span a large range of n. According to the simulations, our sizes for the smallest objects should be considered only as an upper limit. This systematic deviation of the size of the galaxies at small radii  is probably an artefact due to the relative large size of the pixel ($\\sim$0.45\\arcsec) compared to the size of the galaxies. Simulations also provide us with a typical uncertainty in the estimation of the sizes of small galaxies of $\\sim$0.1\\arcsec. We will use this value to estimate the error bars in our measurements.  \\begin{figure} \\epsfig{file=sizesimulations.eps,width=0.5\\textwidth}  \\caption{The figure shows a comparison between the input intrinsic half--light radius (before seeing convolution) and recovered size values in our simulations for the MUNICS-Deep observations. The small points are used to indicate the position of individual mock galaxies. The open squares indicate the mean value and the bars correspond to 1$\\sigma$ dispersion. According to the simulations, the retrieved sizes for the smallest objects should be considered only as an upper limit.}  \\label{simulation} \\end{figure}  \\vspace{-0.35cm} \\subsection{Testing the size estimates: crude upper limits and the effect of different images depth}  A direct method of establishing a crude upper limit to the size of the objects is by fitting a Gaussian profile to the observed galaxies. Under the assumption that the intrinsic surface brightness profile of the galaxies are also described by a Gaussian profile, the effective radius of the galaxy is given by:  \\begin{equation} r_e=1/2\\times\\sqrt{FWHM_{obj}^2+FWHM_{PSF}^2} \\end{equation}  \\noindent In general, galaxies have a surface brightness profile which is much more concentrated than a Gaussian profile, consequently, this method  provides us with an upper (very conservative) limit of the object's size. A direct Gaussian fit to our objects provide a typical value of FWHM$_{obj}$$\\sim$1.2\\arcsec. Using the typical value of seeing that we have in our images, FWHM$_{PSF}$$\\sim$1\\arcsec, this translates into an upper limit to the sizes of the galaxies of r$_e$$\\lesssim$0.8\\arcsec (or r$_e$$\\lesssim$6.5 h$_{70}^{-1}$ kpc at z$\\sim$1.4 in the cosmology used). Local galaxies with M$_{\\star}$$\\sim$5$\\times$10$^{11}$h$_{70}^{-2}$M$_{\\sun}$ are expected to have sizes of $\\sim$10 h$_{70}^{-1}$ kpc (Shen et al. 2003). This crude upper limit estimation (according to our simulations this tecnique will produce estimates $\\sim$1.5 larger than the input values) shows that our high--z massive galaxies are more compact than their local counterparts.   The sizes of three galaxies in our sample were estimated using shallower images than  the rest of the sample. We have checked whether these shallow observations could introduce any bias in the size estimates of these three objects. For the seven galaxies where we have both shallow and deep observations we estimated the sizes in both cases, and the sizes agree within the error bars. We do not observe, in addition, any systematic difference. This result implies that using deeper observations does not unveil the contribution of light from missing wings in the surface brightness distributions. In other words, our images are catching almost all the light of these galaxies.  \\vspace{-0.9cm} ",
        "conclusions": "As stated in the Introduction, it is possible to find in previous works some examples of significantly small massive galaxies at z$\\gtrsim$1. Daddi et al. (2005) have been the first on discussing in detail the nature of these objects. In addition to the four Daddi et al. objects, there is a group of compact (r$_e$$\\lesssim$1 kpc) and  massive galaxies at z$\\sim$1 seen in Fig. 9 from McIntosh et al. 2005. There is also  evidence of compact massive galaxies in Trujillo et al. (2006) and di Serego Alighieri et al. (2005). Finally, Waddington et al. (2002) studied two z$\\sim$1.5 radio--selected early--type galaxies and found, using NICMOS imaging, sizes of $\\sim$0.3 arcsec ($\\sim$ 2.5 kpc). Therefore, the observational existence of these small galaxies appears currently to be well established.  A strong morphological K--correction has been suggested by Daddi et al. (2005) as one of the potential explananation of the compacteness of their objects observed in the UV restframe. However, our observations in the optical restframe reject this possibility. Another potential explanation suggested by Daddi et al. (2005) is the pressence of an unresolved nuclear component (i.e. an AGN). In fact, two of their 4 objects were detected in X--rays. However, the AGN hypothesis is unlikely to explain our observations. Only one of our galaxies (S2F1 443) is detected in deep XMM-Newton pointing of the S2 fields (Severgnini et al. 2005) having L$_{2-10keV}$$\\gtrsim$10$^{43}$ergs s$^{-1}$ and the spectral energy distributions of our objects lack AGN features. Consequently, if AGNs are present in the rest of our galaxies,  their luminosities should be much fainter than 10$^{43}$ ergs s$^{-1}$ or alternatively, they must be heavily obscured (contributing very weakly to the stellar continuum). The quality of our data prevents us to provide a reliable pointlike+S\\'ersic fit analysis of the surface brightness distributions (as done by Daddi et al. 2005). However, an argument  against the  biasing of our size estimates due to the AGN contribution (if at all present)  is the fact that both at J and K the sizes of our objects are very similar. The AGN should contribute much strongly in J--band (where should be the emission lines H$_{\\beta}$ and OIII at z$\\sim$ 1.4) than in K--band where there is not any significant line. The difference in sizes in J and K in our objects are in agreement (within the error bars). For that reason, we think that the AGN (if present) is not affecting (significantly) the size estimates.   The galaxies analysed in this paper are already very massive at z$\\sim$1.4 and their stellar population properties are consistent with  passive fading. However, there is no observational evidence for galaxies as massive and compact as these in the local universe. Consquently, the high--z galaxies  have to increase their sizes since that redshift. This observational fact disagrees with a scenario where the most massive and passive galaxies are fully assembled at z$\\sim$1.4 (i.e. a monolithic scenario). It is worth noting that, whatever channel is used for our galaxies to evolve in size, this process should take place quickly (i.e. $\\lesssim$ 2 Gyr), since galaxies with M$_{\\star}$$>$10$^{11}$h$_{70}^{-2}$M$_{\\sun}$  at z$\\sim$0.8  seem to be all already in place (Cimatti et al. 2006), and to have sizes very similar to their current values (McIntosh et el. 2005).  A very efficient size evolutionary mechanism (r$_e$$\\propto$M$_{\\star}^{1.3}$) is found in dissipationless mergers with radial orbits (Boylan--Kolchin et al. 2006). In this process, galaxies do not evolve  parallel to the local relation (r$_e$$\\propto$M$_{\\star}^{0.56}$). Interestingly, it turns out that these radial disipationless mergers (along the filaments) of massive galaxies  is thought to be the main channel of formation of the Brightest Cluster Galaxies (BCG). Consequently, our very massive and extremely compact galaxies at z$\\sim$1.4 are very likely candidates to evolve into current BCGs. We have explored whether the comoving number densities of the present day BCGs are in agreement with the dry merger hypothesis. The Cole et al. (2001) local stellar mass functions provides the following number densities for a Kroupa IMF: $\\sim$2.5x10$^{-5}$ Mpc$^{-3}$ for objects with M$_{\\star}$$>$3$\\times$10$^{11}$M$_{\\sun}$, and $\\sim$4x10$^{-7}$ Mpc$^{-3}$ for objects with M$_{\\star}$$\\sim$10$^{12}$M$_{\\sun}$. If we assume that the number densities of M$_{\\star}$$>$3$\\times$10$^{11}$M$_{\\sun}$ objects at z$\\sim$1.4 is a 30\\% than the present--day values, and consider that to reach the mass of a BCG we need $\\sim$4 of our galaxies, we would expect a comoving density of $\\sim$20x10$^{-7}$ Mpc$^{-3}$ for objects with M$_{\\star}$$\\sim$10$^{12}$M$_{\\sun}$ today. This is slightly higher than the value measured for Cole et al. (2001) but works reasonable well due to the large uncertainties. Consequently, we think a dry merger scenario can be considered as a reasonable mechanism for the subsequent evolution of our galaxies (Khochfart \\& Burkert 2003; Dom\\'{\\i}nguez--Tenreiro et al. 2006). Alternative mechanisms of galaxy evolution, like dissipative merging, will increase the mass of the galaxies very effectively but will basically maintain unchanged the sizes (Dekel \\& Cox 2006). This will make the discrepancy in sizes between the high-z and the local galaxies even larger. So, we think wet merging is disfavoured as an evolutionary path for our objects.   An interesting open question is understanding how galaxies as massive as those we are dealing with can be so compact in the past. Recently Khochfar \\& Silk (2005) have investigated the effect of  dissipation in major mergers within the CDM paradigm. They find that  early-type galaxies at high redshifts merge from progenitors that have  more cold gas available than their counterparts at lower redshifts. As a  consequence, they claim that the remnant should be smaller in size at high  redshift. Khochfar \\& Silk (2006) have predicted that the size of objects at z$\\sim$1.5 with M$_{\\star}$$\\gtrsim$5$\\times$10$^{11}$h$_{70}$$^{-2}$M$_{\\sun}$ is a factor of $\\sim$3 times smaller than local counterparts. These estimates agree very well with our observations. If this scenario is correct, the progenitor galaxies that merge to form massive spheroids galaxies are progressively  less and less devoided of gas at lower redshift.   \\vspace{-0.9cm}"
    },
    "0608/astro-ph0608182_arXiv.txt": {
        "abstract": "Particle acceleration in relativistic shocks is studied analytically in the test-particle, small-angle scattering limit, for an arbitrary velocity-angle diffusion function $D$. Accurate analytic expressions for the spectral index $s$ are derived using few ($2-6$) low-order moments of the shock-frame angular distribution. For isotropic diffusion, previous results are reproduced and justified. For anisotropic diffusion, $s$ is shown to be sensitive to $D$, particularly downstream and at certain angles, and a wide range of $s$ values is attainable. The analysis, confirmed numerically, can be used to test collisionless shock models and to observationally constrain $D$. For example, strongly forward- or backward-enhanced diffusion downstream is ruled out by GRB afterglow observations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608307_arXiv.txt": {
        "abstract": "We study the cascading of fast MHD modes in magnetically dominated plasma by performing one-dimensional (1D) dynamical simulations. We find that the cascading becomes more efficient as an angle between wave vector and underlying magnetic field increases and fast mode becomes more compressive. We also study imbalanced turbulence, in which wave packets propagating in one-direction have more power than those in the opposite direction.  Unlike imbalanced \\Alfvenic turbulence, the imbalanced fast mode turbulence shows faster cascading as the degree of imbalance increases. We find that the spectral index of the velocity and magnetic field, which are carried by the fast-mode turbulence, quickly reaches stationary value of $-2$. Thus we conclude that the dissipation of fast mode, at least in 1D case, happens not due to weak or strong turbulent cascading, but mostly due to nonlinear steepening. The density fluctuations, which are carried by slow-mode perturbation in the larger scale and entropy-mode perturbation in the smaller scale, depend on initial driving spectrum and a ratio of specific heat. ",
        "introduction": "Magnetohydrodynamical (MHD) turbulence is ubiquitous in astrophysics. All interstellar medium (ISM) phases are well magnetized with Larmor radius of thermal proton much smaller than outer scale. Molecular clouds are no exception, with compressible, often high-Mach number, magnetic turbulence determining most of their properties (see Elmegreen \\& Falgarone 1996, Stutzki 2001). Star formation (see McKee \\& Tan 2002; Elmegreen 2002; Pudritz 2002; Ballesteros-Paredes et al.2006), cloud chemistry (see Falgarone 1999 and references therein), shattering and coagulation of dust (see Lazarian \\& Yan 2002 and references therein) are examples of processes for which knowledge of turbulence is absolutely essential. It is also believed that MHD waves and turbulence are important for the acceleration of the solar wind (e.g. Tu \\& Marsch 1995) as well as winds from stars which possess surface convective layers (e.g. Tsuji 1988).   As MHD turbulence is an extremely complex process, all theoretical constructions should be tested thoroughly. Until very recently matching of observations with theoretical constructions used to be the only way of testing. Indeed, it is very dangerous to do MHD turbulence testing without fully 3D MHD simulations with a distinct inertial range. Theoretical advances related to the anisotropy of MHD turbulence and its scalings (see Shebalin et al. 1983, Higdon 1984, Montgomery, Brawn \\& Matthaeus 1987, Shebalin \\& Montgomery 1988, Zank \\& Matthaeus 1992) were mostly done in relation with the observations of fluctuations in solar wind. Computers allowed an important alternative way of dealing with the problem. While still presenting a limited inertial range, they allow to control the input parameters making it easier to test theoretical ideas.  It is well known that linear MHD perturbations can be decomposed into Alfv\\'{e}nic, slow, fast, and entropy modes with well-defined dispersion relations (see Alfv\\'{e}n \\& F\\\"{a}lthammar 1963). The separation into Alfv\\'en and pseudo-Alfv\\'en modes, which are the incompressible limit of Alfv\\'en and slow MHD modes, is an essential element of the Goldreich-Sridhar (1995, henceforth GS95) model of turbulence. There the arguments were provided in favor of Alfv\\'en modes developing a cascade on their own, while the pseudo-Alfv\\'en modes being passively advected by the cascade. The drain of Alfv\\'enic mode energy to pseudo-Alfv\\'en modes was shown to be marginal.  The separation of MHD perturbations into modes in compressible media was discussed further in Lithwick \\& Goldreich (2001) and Cho \\& Lazarian (2002, 2003 henceforth CL02, CL03, respectively). Even though MHD turbulence is a highly non-linear phenomenon, the modes does not constitute an entangled inseparable mess.  The actual decomposition of MHD turbulence into % each mode was a challenge that was addressed in CL02, CL03. They studied a particular realization of turbulence with mean magnetic field comparable to the fluctuating magnetic field. This setting is, on one hand, rather typical of the most MHD flows in Galaxy and, on the other hand, allows them to use a statistical procedure of decomposition in the Fourier space, where the bases of the Alfv\\'en, slow, fast, and entropy perturbations were defined. The entirely different way to decompose modes in high or low-$\\beta$ cases in real space was shown in CL03 to correspond well to this procedure\\footnote{ In CL03, an isothermal equation of state is used so that entropy-mode fluctuation is prohibited.}.  In particular, calculations in CL03 demonstrated that fast modes are marginally affected by Alfv\\'{e}n modes and follow acoustic cascade in both high and low $\\beta$ medium.  In Yan \\& Lazarian (2002; 2004; henceforth YL04) the fast modes were identified as the major agent in scattering of cosmic rays. Similar results in relation to stochastic acceleration of cosmic rays were reached in Cho \\& Lazarian (2006). Interestingly enough, the dominance of fast modes for both acceleration and scattering is present in spite of the dissipative character of fast modes. The reason for that is isotropy\\footnote{ Recent calculations of fast modes interacting with Alfv\\'{e}n modes by Chandran (2005) indicate some anisotropy of fast modes. In particular, he shows that contours of isocorrelation for fast modes get oblate with shorter direction along the magnetic field. However, the calculations are performed assuming that Alfv\\'{e}n modes are in the weak turbulence regime.  This regime has a limited inertial range. Moreover, anisotropic damping discussed in YL04 is probably a more strong source of anisotropy. Anyhow, the above conclusions about the dominance of fast modes are not altered by the anisotropies that corresponds to squashing contours of isocorrelations in the direction of magnetic field.} of this mode reported in CL03. The Alfv\\'{e}n modes, that are still the default for scattering and acceleration for many researchers, are inefficient due to the elongation of eddies along the magnetic field (Chandran 2000, Yan \\& Lazarian 2002).  Apart from cosmic rays, acceleration of charged interstellar dust is also dominated by fast modes (Lazarian \\& Yan 2002, Yan \\& Lazarian 2003). In stellar magnetospheres with spiral magnetic fields, fast-modes possibly play a role in energetics and dynamics of winds from usual stars \\citep{suz06} and pulsars \\citep{lyu03}. These issues provide an additional motivation to studies of properties of fast mode cascade.    The most efficient interaction of fast modes is expected to happen for waves moving in the same direction (see CL02).  This indicates that, unlike studies of Alfv\\'{e}n modes for which the 3D character of interactions is essential, 1D simulations could still give, at least in some respects, correct physical properties of the fast-mode turbulence. 1D simulations have a big advantage that we can use a number of grid points for a broader inertial range than 2D or 3D simulations. Based on these issues, we carry out 1D decay simulations of fast-mode turbulences. We should cautiously analyze simulation data, because the 1D geometry systematically enhances shocks without dilution to the tangential directions.  In what follows we discuss our code in \\S 2, calculate power spectra in \\S 3, and study cascading of fast modes in \\S 4. We discuss our findings in \\S 5. ",
        "conclusions": "We have shown that the imbalanced fast modes dissipate in a different way than \\Alfven ones. For the \\Alfven turbulence cascading is due to wave packets moving in the opposite direction, therefore, the imbalance makes the cascading less efficient. For fast modes waves moving in the same direction interact more efficiently. We confirmed that the oppositely moving fast modes are not essential for the cascading as discussed earlier (see CL02).  However, the obtained power spectrum ($k^{-2}$) is different from the weak acoustic turbulence ($k^{-3/2}$) suggested in CL02, which may be owing to the different geometries (1D vs. 3D) as discussed in \\S \\ref{sec:teps}. We should note that this problem of the scaling of acoustic turbulence is not settled down because the approximation of weak cascade is controversial \\citep{fm96,zlf92}. It is clear that further research by 3D simulations with a sufficient inertial range is necessary.  We have also found that as the amplitude of compressions increase, the non-linear damping of turbulence speeds up. In our adiabatic calculations the restoring force was increasing due to medium heating and as the result the non-linear steepening and cascading slowed down. Similarly, for the modes that compress both magnetic field and gas, the rate of dissipation was less compared to the modes that mostly compress the gas.    The density perturbation is mainly carried by slow mode perturbation in the large scale and entropy mode in the small scale perturbation in non-isothermal gas. The slow-mode density spectrum is dominated by shocks, $\\propto k^{-2}$, while the entropy-mode, which is a consequence of the dissipation of the initial fast-mode turbulence, preserves the initial turbulent spectrum. As a result, the spectral index of the density perturbation is subject to the initial spectrum of given fast-mode turbulence. In isothermal gas, density spectrum is different in a small scale because entropy-mode fluctuation does not exist. This shows that to study density turbulence one has to carefully treat driving properties of turbulence and to take into account appropriate thermal processes (e.g. Koyama \\& Iutsuka 2002; Larson 2005) rather than assume a polytropic equation of state. While the density turbulence is an interesting issue, obviously 3D treatment is required since the slow- and entropy-mode perturbation is different in 1D and 3D circumstances; we will study this topic in a future paper.  In this paper we have focused on the fast-mode in low-$\\beta$ medium. Various interstellar/intergalactic medium is dominated by magnetic pressure rather than gas pressure, i.e., low-$\\beta$ condition. Referring to the tabulation in YL04 (Table 1 in their paper), galactic halo, and warm and cold ISM are mildly low-$\\beta$ ($\\beta=0.1-1$), and dark clouds are more magnetically dominated, $\\beta=0.01-0.1$. Our simulations are directly applicable to these media.  Magnetospheres of general stars with surface convection (like the sun) and compact objects (white dwarfs and neutron stars) are also in low-$\\beta$ condition, because the density rapidly decreases upwards. Moreover, the turbulence tends to be imbalanced there because it is mainly generated from central stars and dominated by the outward component. Our results show that the fast mode is likely to be the culprit in the heating and acceleration of winds because it dissipates effectively even without the inward component. The advantage of the fast-mode versus Alfv\\'{e}nic turbulence in the stellar wind acceleration is that the latter requires production of counter-waves via reflection or by some other mechanism.  Cascading of imbalanced fast modes is important for the problem of streaming instability evolution, which is the part of the picture of both galactic leaky box model and the acceleration of particles in shocks. This instability arises from the flow of particles along the magnetic field direction and reflects the particles back (Kulsrud \\& Pearce 1968). Our results indicate that the fast mode cascading limits the instability even in the absence of both damping and the ambient turbulence (YL02, Farmer \\& Goldreich 2004).  In our simulation the dissipation is mainly due to MHD shocks besides sub-grid scale damping. However, if we discuss the dissipation in more detail, we need to take into account the microphysics. Galactic halo and stellar winds generally consist of collisionless plasma. In such conditions, the fast-mode as well as slow-mode fluctuations possibly dissipate by transit-time damping (e.g. Barnes 1966; Suzuki et al.2006). In warm ionized ISM, the dissipation is mainly due to Coulomb collisions. If there is a considerable fraction of neutral atoms, e.g. in cold ISM and dark clouds, collisions between neutrals and ions dominate (see Table 1 in YL04).   Turbulence may be very different in 1D and 3D. We find, however, that our attempts to get insight into the dynamics of fast modes using 1D model is meaningful, as the possible transversal deviation $\\delta k$ is limited. Assuming that the fast modes follow the acoustic cascade with the cascading time \\begin{equation} \\tau_{k}=\\omega/k^{2}v_{k}^{2}=(k/L)^{-1/2}\\times V_{ph}/V^{2}, \\label{fdecay} \\end{equation} YL04 used the uncertainty condition $\\delta \\omega t_{cas}\\sim 1$, where $\\delta \\omega\\sim V_{ph} \\delta k (\\delta k/k)$ to estimate $\\delta k$ \\begin{equation} \\delta k/k\\simeq \\frac{1}{(kL)^{1/4}}\\left(\\frac{V}{V_{ph}}\\right)^{1/2}. \\label{dk} \\end{equation} Eq. (\\ref{dk}) provides a very rough estimate of how much the randomization of vectors is expected due to the fast mode cascading. This estimates suggest that deep down the cascade where $kL\\gg 1$ the approximation of waves moving in the same direction could be accurate enough.  In this paper we dealt with fast modes irrespectively of the evolution of the other modes. This possibility corresponds to the earlier theoretical and numerical studies (GS95, Lithwick \\& Goldreich 2001, CL02, CL03). Recently, however, Chandran (2005) discussed the interaction between fast and Alfv\\'{e}n modes under the approximation that the latter constitute {\\it weak} turbulence, i.e. turbulence that cascades slowly due to weak non-linearities. We expect appreciable mitigation of the Alfv\\'en-fast mode interaction compared to those in Chandran (2005), when Alfv\\'{e}nic turbulence is {\\it strong}, i.e. when the turbulence evolves fast, over just one wave period. We feel that the resulting large disparity in the rate of evolution of the Alfv\\'enic and fast modes should result in much less cascading of fast modes by the Alfv\\'enic modes compared to the case of the weak Alfv\\'enic turbulence. Note, that strong Alfv\\'enic turbulence is default for most astrophysical situations, while the inertial range of the weak Alfv\\'enic turbulence is limited\\footnote{It is easy to see that if the injection happens at scale $L$ with Alfv\\'en Mach number $M_A=V/V_A$, where $V_A$ is the Alfv\\'en velocity that arises from the total magnetic field in the fluid, the range for weak turbulence is limited by $[L, LM_A^2]$. While the regime of weak Alfv\\'enic turbulence can still be important (see Lazarian 2006), in most astrophysical situations that regime of strong  Alfv\\'enic turbulence covers much longer inertial range.} (Goldreich  \\& Sridhar 1997, Lazarian \\& Vishniac 1999). In our next paper we shall provide calculations of the fast modes interaction with the Alfv\\'enic modes.  {\\bf Acknowledgments}  T.K.S. is supported by JSPS Research Fellow for Young Scientist, 4607, and a Grant-in-Aid for Scientific Research, 18840009, from MEXT of Japan. AL acknowledges the support of the NSF Center for Magnetic Self-Organization in Laboratory and Astrophysical Plasmas and the NSF grant AST-0307869 and NSF-ATM -0312282. AB acknowledges the support of the Ice Cube project."
    },
    "0608/astro-ph0608131_arXiv.txt": {
        "abstract": "We apply the stellar population synthesis code by Cid Fernandes et al. to model the stellar contribution for a sample of 209 type II AGNs at $0.3<z<0.83$ from the Sloan Digital Sky Survey. The reliable stellar velocity dispersions ($\\sigma_{*}$) are obtained for 33 type II AGNs with significant stellar absorption features. According to the $L_{\\rm [O III]}$ criterion of $3\\times 10^{8} \\Lsolar$, 20 of which can be classified as type II quasars. We use the formula of Greene \\& Ho to obtain the corrected stellar velocity dispersions ($\\sigma_{*}^c$). We also calculate the supermassive black holes masses from $\\sigma_{*}^c$ in these high-redshift type II AGNs. The [O III] luminosity is correlated with the black hole mass (although no correlation between the extinction-corrected [O III] luminosity and the black hole mass), and no correlation is found between the Eddington ratio and the [O III] luminosity or the corrected [O III] luminosity. Three sets of two-component profile are used to fit multiple emission transitions ([O III]$\\lambda\\lambda$4959, 5007 and [O II]$\\lambda\\lambda$3727, 3729) in these 33 stellar-light subtracted spectra. We also measure the gas velocity dispersion ($\\sigma_{g}$) from these multiple transitions, and find that the relation between $\\sigma_{g}$ and $\\sigma_{*}^c$ becomes much weaker at higher redshifts than in smaller redshifts. The distribution of $<\\sigma_{g}/\\sigma_{*}^c>$ is $1.24\\pm0.76$ for the core [O III] line and $1.20\\pm0.96$ for the [O II] line, which suggests that $\\sigma_{g}$ of the core [O III] and [O II] lines can trace $\\sigma_{*}^c$ within about 0.1 dex in the logarithm of $\\sigma_{*}^c$. For the secondary driver, we find that the deviation of $\\sigma_{g}$ from $\\sigma_{*}^c$ is correlated with the Eddington ratio. ",
        "introduction": "Recent advances in the study of normal galaxies and active galactic nuclei (AGNs) are ample observational evidence for the existence of central supermassive black holes (SMBHs) and the relationship between SMBHs and bulge properties of host galaxies (Gebhardt et al. 2000; Ferrarese \\& Merritt 2000; Tremaine et al. 2002; Begelman 2003; Shen et al. 2005). We can use the stellar and/or gaseous dynamics to derive the SMBHs masses in nearby inactive galaxies. However, it is  much more difficult  for the case of AGNs. With the broad emission lines from broad-line regions (BLRs)(e.g. H$\\beta$, Mg II, CIV; H$\\alpha$), the reverberation mapping method and the empirical size-luminosity relation can be used to derive the virial SMBHs masses in AGNs (Kaspi et al. 2000; Vestergaard 2002; McLure \\& Jarvis 2002; Wu et al. 2004; Greene \\& Ho 2006a).  It has been found that nearby galaxies and AGNs follow the same tight correlation between the central SMBHs masses ($M_{\\rm BH}$) and stellar bulge velocity dispersion ($\\sigma_{*}$) (the $M_{\\rm BH}-\\sigma_{*}$ relation) (Nelson et al. 2001; Tremaine et al. 2002; Greene \\& Ho 2006a, 2006b), which also implied that the mass from reverberation mapping method is reliable.  According to AGNs unification model (e.g. Antonucci 1993; Urry \\& Padovani 1995), AGNs can be classified into two classes depending on whether the central engine and BLRs are viewed directly (type I AGNs) or are obscured by circumnuclear medium (type II AGNs). In type I AGNs, by using the broad emission lines from BLRs  (the reverberation mapping method or the empirical size-luminosity relation), we can derive virial SMBHs masses. It is not easy to study their host galaxies because their optical spectra are dominated by the non-stellar emission  from the central AGNs activity. This is especially true for luminous AGNs, where the continuum radiation from central source outshines the stellar light from the host galaxy.  In type II AGNs, the obscuration of BLRs makes both the reverberation mapping method and the empirical size-luminosity relation infeasible to derive SMBHs masses. However, we can use the well-known $M_{\\rm BH}-\\sigma_{*}$ relation to derive SMBHs masses if we can accurately measure the stellar bulge velocity dispersion ($\\sigma_{*}$). There are mainly two different techniques to measure $\\sigma_{*}$, one is the \"Fourier-fitting\" method (Sargent et al. 1977; Tonry \\& Davis 1979), the other is the \"Direct-fitting\" method (Rix \\& White 1992; Greene \\& Ho 2006b and reference therein). These years it has been successful to derive $\\sigma_{*}$ through fitting the observed stellar absorption features, such as Ca H+K $\\lambda \\lambda$ 3969, 3934, Mg Ib$\\lambda \\lambda$ 5167, 5173, 5184 triplet, and Ca II$\\lambda \\lambda$ 8498, 8542, 8662 triplet, etc,  with the combination of different stellar template spectra broadened by a Gaussian kernel (e.g. Kauffmann et al. 2003; Cid Fernandes et al. 2004a; Greene \\& Ho 2006b).  On the other hand, Nelson \\& Whittle (1996) find that the gaseous velocity dispersion ($\\sigma_{g}$) of [O III]$\\lambda$5007 from the narrow-line regions (NLRs) is nearly the same as $\\sigma_{*}$ for a sample of 66 Seyfert galaxies, and suggest that the gaseous kinematics of NLRs be primarily governed by the bulge gravitational potential. Nelson (2001) find a relation between $M_{\\rm BH}$ and $\\sigma_{\\rm [O III]}$ (the [O III]$\\lambda$5007 velocity dispersion) for AGNs, very similar to the relation of $M_{\\rm BH}-\\sigma_{*}$, although with more scatter, which strongly suggests that $\\sigma_{g}$ can be used as a proxy for $\\sigma_{*}$. For lower-redshift type II AGNs with $0.02 < z < 0.3$, Kauffmann et al. (2003) have investigated the properties of their hosts from the Sloan Digital Sky Survey (SDSS) Data Release One (DR1), measured $\\sigma_{*}$ and estimated the SMBHs masses from $\\sigma_{*}$ (Brinchmann et al. 2004). By using this sample, Greene \\& Ho (2005) have measured the gaseous velocity dispersion ($\\sigma_{\\rm g}$) from multiple transitions ([O II] $\\lambda$3727, [O III] $\\lambda$5007, and [S II] $\\lambda\\lambda$6716, 6731) and compared $\\sigma_{*}$ and $\\sigma_{g}$. They find that $\\sigma_{g}$ from these multiple transitions trace $\\sigma_{*}$ very well, although some emission features they show considerable scatters.  Type II quasars are the luminous analogs of low-luminosity type II AGNs (such as Seyfert 2 galaxies). The obscuration of BLRs makes quasars appear to be type II quasars (obscured quasars). Some methods have been used to discover type II quasars, but only a handful have been found.  Recently, Zakamsa et al. (2003) present a sample of 291 type II AGNs at redshifts $0.3 < Z < 0.83$ from the SDSS spectroscopic data. About half are type II quasars if we use the [O III] $\\lambda$5007 line luminosity to  represent the strength of the nuclear activity. What is the $\\sigma_{*} - \\sigma_{g}$ relation for type II quasars? And what are their SMBHs masses and the Eddington ratios $L_{\\rm bol}/L_{\\rm Edd}$ (i.e. the bolometric luminosity as a fraction of the Eddington luminosity)?  Here we use the sample of Zakamsa et al. (2003) to study these questions for type II quasars. In section 2, we introduce the data and the analysis. Our results and discussion are given in Sec. 3. All of the cosmological calculations in this paper assume $H_{0}=70 \\rm {~km ~s^ {-1}~Mpc^{-1}}$, $\\Omega_{M}=0.3$, $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "The stellar population synthesis code is used to model the stellar contribution for a sample of 209 type II AGNs at redshifts $0.3<z<0.83$ from SDSS. According to the $L_{\\rm [O III]}$ criterion of $3\\times 10^{8} \\Lsolar$, 20  can be classified as type II quasars. The main conclusions can be summarized as follows.   \\begin{itemize}  \\item{The reliable $\\sigma_{*}$ are measured for 33 type II AGNs with significant stellar absorption features. We use the formula of Greene \\& Ho to obtain the corrected stellar velocity dispersions ($\\sigma_{*}^c$ and SMBHs masses are calculated from the $M_{\\rm BH}-\\sigma_{*}^{c}$ relation. A median strong relation between the [O III] luminosity and the SMBH mass is found (although no correlation between the extinction-corrected [O III] luminosity and the SMBH mass); no correlation is found between the Eddington ratio and the [O III] luminosity or the extinction-corrected [O III] luminosity.}  \\item{The gas velocity dispersion ($\\sigma_{g}$) in NLRs is measured using three sets of two-gaussian profiles to fit [O III]$\\lambda\\lambda$4959, 5007 and [O II]$\\lambda\\lambda$3727, 3729) in these 33 stellar-light subtracted spectra. We find that the relation between $\\sigma_{g}$ and $\\sigma_{*}^c$ becomes much weaker at higher redshifts.}  \\item{The distribution of $<\\sigma_{g}/\\sigma_{*}^c>$ is $1.24\\pm0.76$ for the core [O III] line and $1.20\\pm0.96$ for the [O II] line, which suggests that $\\sigma_{g}$ can trace $\\sigma_{*}^c$ within about 0.1 dex in the logarithm of $\\sigma_{*}^c$. The deviation of $\\sigma_{g}$ from $\\sigma_{*}^c$ is correlated with the Eddington ratio.}   \\end{itemize}"
    },
    "0608/astro-ph0608292.txt": {
        "abstract": "%Context {A recent comparison between deep VLA H{\\sc i} observations and dynamical models of the Virgo cluster spiral galaxy NGC~4654 has shown that only a model involving a combination of a tidal interaction and ram pressure can reproduce the data.} %aims {Deep radio polarization studies, together with detailed MHD modeling, can independently verify those conclusions, that are based on H{\\sc i} observations and dynamical models.} %methods {We performed deep polarized radio-continuum observations of the Virgo cluster spiral galaxy NGC~4654 with the Effelsberg 100m telescope at 8.35~GHz and the VLA at 4.85~GHz. Detailed 3D MHD simulations were made to determine the large-scale magnetic field and the emission distribution of the polarized radio continuum in the model, during the galaxy evolution within the cluster environment.} %results {This direct comparison between the observed and simulated polarized radio continuum emission corroborates the earlier results, that the galaxy had a recent rapid close encounter with NGC~4639 and is undergoing weak ram pressure by the intracluster medium.} %conclusions {This combination of deep radio polarization studies and detailed MHD modeling thus gives us unique insight into the interactions of a galaxy with its cluster environment. It represents a diagnostic tool that is complementary to deep H{\\sc i} observations.} ",
        "introduction": "}  The bright Virgo cluster spiral galaxy NGC~4654 is located at the periphery of the cluster at a projected distance of $3.4^{\\circ}=1$~Mpc\\footnote{We use a distance of 17~Mpc to the Virgo cluster} from the cluster center (M87). Its two-armed spiral structure is asymmetric with a shorter spiral arm to the northwest and a larger wider spiral arm to the southeast. Its H{\\sc i} distribution shows a sharp edge to the northwest and a large, very low column density tail to the southeast (Phookun \\& Mundy \\cite{pm95}). While the northwestern part of the H{\\sc i} velocity field corresponds to a rising rotation curve, rotation is almost absent in the very low column-density tail. The H$\\alpha$ image (Koopmann et al. \\cite{koo01}) shows enhanced star formation at the northwestern edge of the galactic disk.  Another spiral galaxy, NGC~4639, is located only $18'=90$~kpc in projection from NGC~4654 and might be considered to be a companion galaxy of NGC~4654. NGC~4639 has a ring-like H{\\sc i} distribution. Since it is located at the edge of the field observed by Phookun \\& Mundy (\\cite{pm95}) and Cayatte et al.~(\\cite{cay94}), it is not clear if its gas distribution and kinematics are perturbed. While Warmels (\\cite{war88}) observed extended gas in the east of NGC~4639, Cayatte et al. (\\cite{cay94}) found atomic gas to the southwest.  Two kinds of interactions can in principle be responsible for the observed asymmetries and perturbations of NGC~4654: (i) a gravitational interaction with a massive galaxy or (ii) ram pressure due to the galaxy's rapid motion within the hot and tenuous intracluster medium. In order to investigate the recent evolutionary history of NGC~4654, Vollmer (\\cite{vol03}) made numerical simulations including a gravitational interaction, ram pressure, and a combination of both, concluding that \\begin{itemize} \\item a past ram-pressure stripping event where the galaxy has already passed the cluster center 800~Myr ago and is now emerging from the cluster core gives rise to a low surface-density tail in the southeast of the galaxy center. This tail shows clear signs of rotation; \\item a gravitational interaction between NGC~4654 and NGC~4639 cannot produce a low surface density tail; \\item only a gravitational together with a relatively small amount of ram pressure leads to the observed low surface-density tail where rotation is almost absent. \\end{itemize}  The results of Vollmer (\\cite{vol03}) suggest that NGC~4654 had a tidal interaction with another massive galaxy in the recent past ($\\sim$500~Myr ago). A model where the position and radial velocity of the perturbing massive galaxy are reasonably close to those of NGC~4639 can explain the perturbed stellar distribution and the gas distribution of the inner disk. Therefore, Vollmer (\\cite{vol03}) identified NGC~4639 as the perturbing galaxy.  A caveat of the gravitational interaction scenario is that NIR Tully-Fisher measurements yield a line-of-sight distance between NGC~4654 and NGC~4639 of more than $5$~Mpc (Gavazzi et al. \\cite{gav99}). While NGC~4654 is located $1-2$~Mpc in front of the cluster core, NGC~4639 is located far behind ($\\sim 5\\pm1$~Mpc). Observations of Cepheids give independent distance estimates for a number of Virgo cluster galaxies (Gibson et al. \\cite{gib00}, Freedman et al. \\cite{fre01}). Five galaxies are in both, the Gavazzi et al.~(\\cite{gav99}) and the Freedman et al.~(\\cite{fre01}) samples. For 4 galaxies (NGC~4321, NGC~4535, NGC~4536, and NGC~4536), the difference between the distance moduli obtained from the two methods are small ($|\\Delta \\mu_0| < 0.2$). However, this difference is much larger for NGC~4639 ($|\\Delta \\mu_0| =0.4$). The Cepheid distance estimates place NGC~4639 closer, but still a few Mpc behind the center of the Virgo cluster (M87). Thus these measurements corroborate the large NIR Tully Fisher distance of NGC~4639.  Since the preferred scenario of Vollmer et al. (\\cite{vol03}) includes ram pressure of $p_{\\rm ram}=200$~cm$^{-3}$(km\\,s$^{-1}$)$^{-1}$, a galaxy velocity of $\\sim 1000$~km\\,s$^{-1}$ implies an intracluster medium (ICM) density of $n_{\\rm ICM} \\sim 10^{-4}$~cm$^{-3}$. At the distance of NGC~4639, the ICM density according to Schindler et al. (\\cite{sch99}) is lower than $10^{-5}$~cm$^{-3}$. This suggests that either (i) NGC~4654 had a gravitational interaction with another massive object, (ii) the Tully Fisher method underestimates the line-of-sight position of NGC~4654, (i.e. NGC~4654 is too bright with respect to its rotation velocity), or (iii) the ICM distribution is elongated along the line-of-sight, similar to the galaxy distribution (Gavazzi et al. \\cite{gav99}). In the case of scenario (i) the perturber has to be a massive dark galaxy (see, e.g., Minchin et al. \\cite{min05}) since there is no other visible companion than NGC~4639.  In this article we test the viability of the preferred scenario of Vollmer (\\cite{vol03}). We used polarized radio-continuum observations of NGC~4654, together with an MHD model, to put further constraints on the models. The radio continuum emission is due to relativistic electrons that gyrate around the galactic magnetic field (synchrotron emission). If the polarized radio-continuum emission (called PRCE hereafter) is free of Faraday rotation (typically at frequencies above $\\sim 4$~GHz), it traces the ordered large-scale magnetic field directly. This magnetic field is extremely sensitive to shear and compression motions, which are not easily detected in 2D velocity fields. Using polarized radio-continuum emission at 6~cm, Vollmer et al. (\\cite{vol04}) could unambiguously identify the ram-pressure compressed region in the Virgo cluster spiral galaxy NGC~4522. We were able to reproduce the H{\\sc i} distribution, velocity field, and the distribution of polarized radio-continuum emission with our model simulations (Vollmer et al. 2006).  In order to understand the behavior of the large-scale magnetic field and thus the polarized radio-continuum emission during a ram-pressure stripping event, Otmianowska-Mazur \\& Vollmer (\\cite{kom03}) solved the induction equation on the velocity fields provided by a dynamical model. They showed that a prominent maximum of polarized radio emission is observed during compression and in a state where the galaxy has passed the cluster center, the ram-pressure decreases, and the gas pushed by ram pressure falls back onto the galactic disk. We applied the method described in Otmianowska-Mazur \\& Vollmer (\\cite{kom03}) to NGC~4654; i.e., we solve the induction equation on the velocity fields of the simulations of Vollmer (\\cite{vol03}) and compare the results to our polarized radio-continuum observations.  The Effelsberg and VLA radio continuum observations are presented in Sect.~\\ref{sec:observations}. In Sect.~\\ref{sec:model} we explain the 3D MHD model and how we calculated the model of polarized radio-continuum emission maps. The results of our model calculations are shown in Sect.~\\ref{sec:results} followed by direct comparison between our observations and simulations (Sect.~\\ref{sec:comparison}). These results are discussed in Sect.~\\ref{sec:discussion} and we give our conclusions in Sect.~\\ref{sec:conclusions}. ",
        "conclusions": "}   The Virgo spiral galaxy NGC~4654 was observed with the Effelsberg 100m telescope at 8.35~GHz and the VLA at 4.85~GHz in polarization. We find an asymmetric distribution of the polarized intensity with a strong peak in the south of the galaxy. This shows that asymmetric polarized radio continuum distributions are typical of cluster spiral galaxies (Vollmer et al. \\cite{vol04}, Urbanik \\cite{urb05}).  The polarized radio continuum emission was used as an additional diagnostic tool for the interaction between NGC~4654 and its environment. To do so, we solved the induction equation for the dynamical model evolution of the interstellar medium of a cluster spiral galaxy and calculated the large-scale magnetic field in this way. With an assumed distribution of relativistic electrons, the distribution of polarized radio-continuum emission can be calculated and directly compared to our observations. The input velocity fields are those from the simulations of Vollmer (\\cite{vol03}), who simulated three different interaction scenarios: (i) a gravitational interaction between NGC~4654 and NGC~4639, (ii) ram pressure stripping by the intracluster medium, and (iii) a gravitation interaction, along with a small constant ram pressure.  The direct comparison between our MHD simulations and the polarized radio-continuum observations shows that only the model with gravitation interaction and low, constant ram pressure can reproduce observations in a satisfactory way. We found that the low-resolution Effelsberg data are sufficient for distinguishing between the models. Thus, our study corroborates the results of Vollmer et al. (\\cite{vol04}). This mixed interaction succeeds in reproducing the H{\\sc i} gas distribution and velocity field (Vollmer \\cite{vol03}) and the polarized radio-continuum distribution.  Since Cepheid measurements place NGC~4639 $\\sim 5$~Mpc behind the center of the Virgo cluster, we suggest that the distribution of the Virgo intracluster medium is highly elongated thereby providing the necessary ram pressure. Thus, polarized radio-continuum emission, in combination with detailed MHD modeling, provides an important additional diagnostic tool for studying interactions between a galaxy and its cluster environment. This tool is complementary to deep H{\\sc i} observations."
    },
    "0608/astro-ph0608077_arXiv.txt": {
        "abstract": "{} {We are studying the extremely young cluster of M17 to investigate the birth of high-mass stars and the initial mass function.} {Deep $JHKL$ imaging and $K$-band spectroscopy from the VLT of 201 stars toward the cluster is presented.} {The majority of 104 stars show the CO band-head in absorption. Half of them emit X-rays and/or have infrared excess, indicative of very young objects. Their intrinsic IR luminosity is compatible with intermediate and high-mass pre-main sequence stars. Nine additional stars have the CO feature in emission, while sixty sources are lacking any stellar spectral feature due to veiling by circumstellar dust.} {We suggest that CO absorption is -- as in the case of low-mass stars -- also a common feature during the early evolution of stars with higher masses. According to model calculations the observed CO absorption is most likely a sign of heavily accreting protostars with mass accretion rates above $10^{-5}\\,$M$_\\odot$\\,yr$^{-1}$.} ",
        "introduction": "Numerous young stellar objects (YSOs) display the CO band-heads at $2.3$ $-$ $2.4\\,\\mu$m in emission. There is consensus that this CO emission originates in extremely dense ($n_{\\rm H} \\ge 10^{10}$ cm$^{-3}$) and warm (1500\\,K $<$ $T$ $<$ 4500\\,K) regions associated with YSOs and is likely a result of the disk accretion process during early stellar evolution. A variety of mechanisms and models have been proposed to explain their precise origin. These include circumstellar disks, stellar winds, magnetic accretion mechanisms, and inner disk instabilities (e.g., Carr 1989; Carr et al. 1993; Chandler et al. 1993; Biscaya et al. 1997). Recently, the CO emission has been used to infer properties of the associated circumstellar disks (Bik \\& Thi 2004; Blum et al. 2004).  CO band-heads in absorption, in contrast, are typical of the photospheres of cool, mostly evolved stars. Nevertheless, initial surveys of star--forming regions (e.g., Casali \\& Eiroa 1996) have demonstrated that low-luminosity Class\\,II sources, i.e., stars with flat or decreasing spectral energy distributions (SEDs), such as T\\,Tauri stars, may also display strong CO absorption features. Sources with steeply rising SEDs (Class\\,I) showed much weaker, or even undetectable, CO absorption. \\cite{greene96} also found that the CO bands are much weaker in Class\\,I sources. On the other hand, CO absorption arising from an expanding shell or from an accretion disk has been discussed in the context of FU Orionis objects (Hartmann et al. 2004).  Models have shown that variations in the strength and the appearance of the CO feature -- either in emission or in absorption -- depend on parameters like the stellar temperature, the geometry of the circumstellar material, and the mass accretion rate onto the protostellar object (Carr 1989; Calvet et al. 1991). The weakness or even the absence of any CO feature may therefore be a complex combination of various competing mechanisms.  We investigate the stellar content of the extremely young cluster NGC 6618 that is exciting the luminous \\ion{H}{ii} region of M17 at a distance of 1.9\\,kpc (Schmidt et al., in prep.). As we will show in the following, we have found a considerable population of CO sources in this cluster. Only a few of them show the CO band-head in emission (COES); the majority of objects display CO in absorption (COAS). Associated X-ray emission and IR excess suggest that the youngest generation of intermediate and high-mass stars is in the process of heavy accretion. ",
        "conclusions": "\\label{discussion}  Compared to previous CO band-head studies in star--forming regions, the present survey is the most extensive one; in addition, it is the only one that is fairly unbiased. \\cite{casali96} suggested from their sample of 44 YSOs in six regions that CO absorption is very common in low-mass YSOs; from their Fig.\\,2 we estimate the fraction of COAS to be about 50\\%. Likewise, \\cite{casali96} argue that Class\\,II sources tend to show CO absorption, while Class\\,I sources are featureless. In the following, we will explore the nature of the M17 sources in more detail.    \\subsection{Cluster members or field stars?}  In accordance with previous studies we interpret the 9 COES as YSOs within M17. Models for the thermal continuum emission from dusty infalling envelopes around protostars indicate that the envelope emission can exceed the stellar plus disk photospheric emission by almost an order of magnitude (Calvet et al. 1997), thus producing featureless spectra. The veiling of an envelope weakens the CO absorption lines while a disk will only amplify the CO feature. Therefore, the 60 featureless objects are very likely Class\\,I sources within M17. Those 54 COAS with X-ray emission and/or IRE are also young cluster members, as demonstrated above for the two objects CEN\\,30 and 57. Only the nature of the 50 remaining COAS is a priori less clear because there is no direct ``youth indicator'' for classification purposes. However, there are several issues indicating that there might be further YSOs in this group.  The first argument comes from the spatial distribution: most of the COAS are located in the immediate vicinity of other young cluster members and in regions with pronounced nebular emission. This makes it likely that a considerable fraction of them is related to the M17 cluster. Secondly, the infrared colours of COAS without X-rays are on average redder than for sources with X-ray emission (see Fig.\\,\\ref{f2}); likewise, their apparent IR-brightness is fainter by one or two magnitudes compared to the other stars in the sample. Given that the scatter of intrinsic $HKL$ colours for all luminosity classes is fairly small with respect to the observed colours in Fig.\\,\\ref{f2}, one must assume that this large reddening is primarily due to interstellar dust. Thus, all of them have 20 $<$ $A_V$\\,$<$ 45\\,mag, a fact that might explain the non-detections by Chandra and that argues in favour of deeply embedded cluster members, at least for some of them. Finally we want to note that the groups of COES and featureless stars also contain 22\\% and 32\\% of objects that neither show X-rays nor IRE; nevertheless, these stars are most likely YSOs. Obviously, the absence of X-ray emission or IRE is no conclusive argument against youth.   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics*[width=\\textwidth]{5946fig3.eps}} \\caption{$H$-band luminosity function of probable cluster members in M17. The colour code is identical to that of Fig.\\,\\ref{f1}. The arrows denote the range where stars with X-ray emission or IRE occur. \\label{f3}} \\end{figure}    \\subsection{High or low-mass stars?}  The question concerning the mass of these YSOs is very important. To estimate their luminosity we use the $J-H$ colours, where contamination by dust emission is likely to be small. Dereddening is achieved by adopting $(J-H)_0 \\sim 0.16\\,$mag, the mean intrinsic colour of main sequence stars. As the range of intrinsic colours for dwarf and giant stars is sufficiently small ($-0.1$ $<$ $(J-H)_0$ $<$ $0.3\\,$mag for types O to K), the dereddening will not be affected much by the stellar temperature. In this way, we obtain absolute magnitudes of $-7.7 < M_H < 2.6\\,$mag for the stars at the distance of M17. This transforms  into intermediate to high-mass stars, mostly earlier than A0, adopting main sequence luminosities. Some COAS that can be independently dereddened on the basis of their X-ray spectra corroborate this result: most of them appear to be early B-types; this holds also for the COES.  Figure~\\ref{f3} summarises the major statistical findings in the form of an $H$-band luminosity function, adopting main sequence luminosities. Those 50 COAS that do not show unique youth indicators have been omitted from Fig.~\\ref{f3}, although there might be further YSOs among them as discussed above. The distribution attains its maximum around B3; there are only 12\\% A- and 2\\% F-type stars. The remaining stars have absolute $H$-magnitudes compatible with intermediate and high mass stars, i.e., earlier than A0. Of course, at the faint end this distribution is influenced by the limiting magnitude of our spectroscopy. The faintest source for which a spectrum could be obtained has an apparent $H$-band brightness of 17.6\\,mag. Thus, it follows that we can see stars of, e.g., type A3 only if their visual extiction is below $A_V \\sim 25$\\,mag; types K6 can only be detected with $A_V \\le 5\\,$mag.  Concerning the individual groups in Fig.~\\ref{f3}, the featureless objects are basically later than B3, while the COES are earlier than B5. The COAS are distributed across the entire luminosity range. The range $-5 < M_H < -2\\,$mag contains 8 COAS that have X-ray emission and/or IRE. This suggests the presence of YSOs with extremely high luminosity.  The above results have assumed the stars to be on the main sequence, although both CO band-heads and IRE indicate pre-main sequence objects. While high-mass stars evolve mainly at constant luminosity and therefore will not alter the overall distribution of spectral types (or masses) in Fig.~\\ref{f3}, low-mass pre-main sequence objects change their luminosity by more than three orders of magnitude and thus will simulate earlier types and higher masses. If it is true that the $H$-band luminosity function is contaminated by objects of lower masses, these objects must be in an extremely early, i.e., protostellar evolutionary stage (due to their high luminosity).  According to the model by \\cite{calvet91} the \"young\" COAS, i.e., those with X-ray emission or IRE, differ from the COES in the sense that either the underlying YSO is cooler or the mass accretion rate is higher. Corroborated by our early-type spectra for CEN\\,30 and 57, we favour the latter interpretation and identify those COAS as intermediate to high-mass protostars with mass accretion rates above $10^{-5}\\,$M$_\\odot$\\,yr$^{-1}$. Sources with featureless spectra may be similar to other hot YSOs like the Herbig Ae/Be star AB Aurigae (Hartmann et al. 1989) and may form a population between the CO emission and absorption sources."
    },
    "0608/astro-ph0608506_arXiv.txt": {
        "abstract": "Many high contrast coronagraph designs have recently been proposed. In this paper, their suitability for direct imaging of extrasolar terrestrial planets is reviewed. We also develop a linear-algebra based model of coronagraphy that can both explain the behavior of existing coronagraphs and quantify the coronagraphic performance limit imposed by fundamental physics. We find that the maximum theoretical throughput of a coronagraph is equal to one minus the non-aberrated non-coronagraphic PSF of the telescope. We describe how a coronagraph reaching this fundamental limit may be designed, and how much improvement over the best existing coronagraph design is still possible. Both the analytical model and numerical simulations of existing designs also show that this theoretical limit rapidly degrades as the source size is increased: the ``highest performance'' coronagraphs, those with the highest throughput and smallest Inner Working Angle (IWA), are the most sensitive to stellar angular diameter. This unfortunately rules out the possibility of using a small IWA ($<\\lambda/d$) coronagraph for a terrestrial planet imaging mission.  Finally, a detailed numerical simulation which accurately accounts for stellar angular size, zodiacal and exozodiacal light is used to quantify the efficiency of coronagraph designs for direct imaging of extrasolar terrestrial planets in a possible real observing program. We find that in the photon noise limited regime, a 4m telescope with a theoretically optimal coronagraph is able to detect Earth-like planets around 50 stars with 1hr exposure time per target (assuming 25\\% throughput and exozodi levels similar to our solar system). We also show that at least 2 existing coronagraph design can approach this level of performance in the ideal monochromatic case considered in this study. ",
        "introduction": "Direct imaging and characterization (through low resolution spectroscopy) of extrasolar terrestrial planets (ETPs) is an exciting but challenging scientific application of coronagraphy. While the angular separation is well within the reach of a mid-sized telescope in the visible, the high star-planet contrast (about $10^{10}$ in the visible) requires a wavefront stability that can only be obtained with a space telescope. The low apparent luminosity of an Earth-like planet ($m_V \\approx 29$ for an Earth at 10pc) requires the precious planetary photons to be well isolated from the stellar light; otherwise, photon noise from the stellar light drives the exposure time unreasonably high and speckles from the stellar image, if not calibrated, prevent detection. A system capable of delivering $10^{10}$ contrast at less than 0.1'' is therefore required for efficient detection and characterization of ETPs around a reasonably large ($\\approx$30) sample of nearby stars.   The purpose of this work is to identify and quantify how fundamental physics imposes hard limits on what coronagraphs can achieve for the detection of terrestrial planets, and to compare these limits with what existing coronagraph designs can achieve.   Coronagraphs designs simulated in this work are introduced and \\S\\ref{sec:existingC}, where they are tentatively grouped in a four families. The configuration adopted for each design in this study is also described. In \\S\\ref{sec:idealcase}, a metric to quantify coronagraph performance is proposed, allowing existing coronagraph designs to be evaluated in the ``ideal'' case (monochromatic light, unresolved central star). A theoretical limit of coronagraph performance in this ideal case is also derived. In the next 2 sections, we closely look at the most fundamental (unavoidable) deviations from this ideal case, and their effects on coronagraphic performance: stellar angular size (\\S\\ref{sec:stellarsize}) and presence of a zodiacal/exozodiacal background (\\S\\ref{sec:zodi}). Finally, a detailed simulation of coronagraphic observations of a sample of nearby stars is performed in \\S\\ref{sec:montecarlo} to derive the efficiency of coronagraphs for detection of ETPs. ",
        "conclusions": "We have quantified the performance limit imposed by fundamental physics on the direct imaging of ETPs. A solution to build an ``optimal'' coronagraph reaching this limit was proposed, but may be technologically difficult to implement. Fortunately, we have also demonstrated that two existing coronagraph designs (PIAAC and OVC6) yield performances reasonably close (within a factor 2 in number of accessible targets) to this fundamental limit. These recently proposed designs could image ETPs around up to about 100 stars with a 4m telescope, or several hundred stars with an 8m telescope.  For the coronagraphs to perform as detailed in this study, two serious technological challenges must however be solved: \\begin{itemize} \\item{{\\bf Wavefront control at the sub-\\AA$\\:$level}. Thanks to the large number of stellar photons available, recent detector developments, and the expectation that a space telescope can be designed to offer a relatively stable wavefront (with little rapidly varying aberrations, as opposed to atmospheric turbulence for ground-based telescopes), it may be reasonable to hope for this level of correction. Several authors have recently explored promising wavefront control schemes free of aliasing and non-common path aberrations \\citep{codo04,labe04,guyo05b,bord06}, and contrast levels obtained in laboratory \\citep{trau03} are approaching the desired performance.} \\item{{\\bf Chromatic effects} can seriously limit the coronagraph's usable spectral band. A large number of solutions have been proposed to ``achromatize'' coronagraphs, many of them specific to a single coronagraph design \\citep{abe01,soum03a,aime05a,mawe05,pluz06,swar06}. In addition to coronagraph-induced chromatic effects, the beam delivered by the telescope is expected to be chromatic in amplitude and optical pathlength differences (non-ideal coatings, chromatic diffraction propagation between optics). In general, chromatic effects can be mitigated by reducing chromaticity in individual components (beam splitters, phase-shifters, apodization masks) or/and by splitting the spectral band into several narrower bands with dichroics.} \\end{itemize}  More fundamentally, the assumption that detection is photon noise limited seems highly optimistic: ETPs are embedded in an exozodiacal light unlikely to be perfectly smooth and symmetrical. Since good angular resolution may help to distinguish between a planet and exozodiacal light features (such as arcs or rings), it is important to avoid using coronagraphs that clip or strongly attenuate the edge of the telescope pupil if at all possible.  The results obtained in this work should therefore be cautiously considered as upper limits on what may be achieved with a particular coronagraph/space telescope design for the specific goal of imaging ETPs around nearby stars. We also note that several low throughput/large IWA coronagraphs studied in this paper can be more favorably designed for ground-based applications, where contrast requirements are relaxed by about 4 orders of magnitude.  The linear algebra approach to coronagraphy presented in this work successfully explains key characteristics and fundamental limits of coronagraphy. While we have focused our study on the fundamental limitations of coronagraphy (due to stellar angular size and background contamination), this model can be extended to evaluate the sensitivity of coronagraphs to wavefront aberrations other than tip-tilt. The same model can also be used to derive fundamental limits and guide the design of coronagraphs for more ``exotic'' pupil shapes, from segmented pupils telescopes to short or long baseline interferometric arrays."
    },
    "0608/astro-ph0608489_arXiv.txt": {
        "abstract": "We report the results of a globally coordinated photometric campaign to search for transits by the $P\\sim30 \\, {\\rm d}$ and $P\\sim60 \\, {\\rm d}$ outer planets of the 3-planet system orbiting the nearby M-dwarf Gl 876. These two planets experience strong mutual perturbations, which necessitate use of a dynamical (four-body) model to compute transit ephemerides for the system. Our photometric data have been collected from published archival sources, as well as from our photometric campaigns that were targeted to specific transit predictions. Our analysis indicates that transits by planet ``c'' ($P\\sim30 \\, {\\rm d}$) do not currently occur, in concordance with the best-fit $i=50^{\\circ}$ co-planar configuration obtained by dynamical fits to the most recent radial velocity data for the system. Transits by planet ``b'' ($P\\sim60 \\, {\\rm d}$) are not entirely ruled out by our observations, but our data indicate that it is very unlikely that they occur. Our experience with the Gl 876 system suggests that a distributed ground-based network of small telescopes can be used to search for transits of very low mass M-stars by terrestrial-sized planets. ",
        "introduction": "The Gl 876 planetary system ranks with the most remarkable dicoveries that have emerged from the first decade of extrasolar planet detections. As described by Rivera et al. (2005), Laughlin et al. (2005), Butler et al. (2001) Marcy et al. (2001), Delfosse et al. (1998) and Marcy et al. (1998), there are a number of reasons why the Gl 876 planets are extraordinary. Planet ``d\", with its 1.974 day orbit and $M \\sin(i)=5.9 M_{\\oplus}$ is the lowest-mass exoplanet known to orbit a nearby main-sequence star (Rivera et al. 2005). Planets ``c\" and ``b\" are the only Jovian-mass companions known to orbit an M-dwarf, and their clear participation in 2:1 mean-motion resonance gives important dynamical clues to the formation and evolution of the system (e.g. Lee and Peale 2002). Furthermore, the Gl 876 planets are the closest exoplanets that have been reliably characterized. Their radial velocity amplitudes induce an overall signal-to-noise of nearly 100, and they have have been observed for well over a decade.  The detection of Gl 876 d by Rivera et al. (2005) underscored the rapid development of the Doppler technique. The detection of planets having only a few Earth masses is the culmination of advances described by Butler et al. (1996), Marcy \\& Butler (1998), Marcy, Cochran, \\& Mayor (2000), Marcy et al. (2005), and Lovis et al. (2006). Nevertheless, while radial velocity precision has improved to better than 1 m/s, Keplerian orbital fits to Doppler velocities give estimates of $M_{\\rm pl}\\sin(i)$ rather than $M_{\\rm pl}$. In order to ascertain the true mass of a planet, one must obtain an independent measure of the orbital inclination, $i$, which is most easily measured when the planet is observed to transit across the face of the parent star.  The now-celebrated companion to HD 209458 provided the first example of a transiting extrasolar planet, and this body has generated an incredible variety of observational and theoretical investigations. HD 209458b was initially detected with the radial velocity method, and was then discovered to transit on the basis of follow-up photometric observations at the predicted transit times (Henry et al. 1999, 2000; Charbonneau et al. 2000). HD 209458b's well-sampled light curve, when combined with high-precision radial velocity data, has permitted precise measurements of the planet's basic attributes (e.g. Mazeh et al. 2000, Brown et al. 2001, Laughlin et al. 2005, Winn \\& Holman  2005).  To date, nine additional planets have been found to transit their parent stars, and the varied uses to which these transits have been put are described in the review article by Charbonneau et al. (2006). Our goal in this paper is to ascertain whether the two outer planets orbiting Gl 876 (IL Aqr, GJ 876, HIP 113020) can similarly be observed in transit. A positive detection would give a further major improvement in our understanding both of this specific system and of planetary properties in general.  Rivera et al. (2005) carried out a photometric search for transits by planet d ($P\\sim1.974 {\\rm d}$), and showed that such transits do not occur. To date, photometric searches for transits by planet c ($P\\sim30 {\\rm d}$) and b ($P\\sim60 {\\rm d}$) have not been reported. Although the \\emph{a priori} geometric transit probabilities for Gl 876 c and b are only $P_{\\rm tr}=1.4\\% $ and $P_{\\rm tr}=0.9\\% $, respectively, the expected photometric transit depths are greater than 10\\%. Librational motions arising from the resonance between the two outer planets lead to predicted transit epochs for the outer planets that are not spaced evenly in time. (The resonant structure of the system is discussed further in Laughlin \\& Chambers 2001; Rivera \\& Lissauer 2001; Kley 2003; and Beaug\\'{e}, Michtchenko, \\& Ferraz-Mello 2005).  Benedict et al. (2002) used HST astrometry to derive an orbital inclination for Gl 876b of $i=84\\pm 6^\\circ $. Their result thus suggests that planet b has a mass equal or close to the RV-determined value of $M\\sin (i)=1.89M_{\\mbox{JUP}}$, and that the system is within 1-sigma of being viewed edge-on. The implication of the Benedict et al. (2002) measurement is that occurence of transits is much more likely than the a-priori geometric probability would suggest. The Benedict et al. (2002) result is in conflict, however, with the dynamical models of the system presented by Rivera et al. (2005), who find a best fit to the radial velocity data when the system (assumed co-planar) is inclined by $i\\sim50\\pm3^{\\circ}$ with respect to the plane of the sky. Therefore, if transits are detected, then the Benedict et al. (2002) results will be confirmed at the expense of the conclusions of Rivera et al. (2005).  Gl 876b and Gl 876c are members of an as-yet unobserved class of giant planets with intermediate surface temperatures. The expected effective temperatures of ${T_{\\rm eff}}_b =160$k and ${T_{\\rm eff}}_c =200$K of the planets are much lower than the value $T_{\\rm eff} \\sim 1350$K measured with Spitzer for HD 209458b (Charbonneau et al. 2005), but still considerably higher than the $T_{\\rm eff} \\sim 135$K measured for Jupiter. Determination of their physical properties would thus provide a useful link between Jupiter and Saturn on the one hand, and objects such as HD 209458b on the other. Accurate measurement of their radii would also give clues to their interior structure (see Charbonneau et al. 2006).  Finally, the Gl 876 system with its strongly gravitationally interacting planets and large planet-to-star radius ratios, has the potential to display an extraordinarily informative set of transit light curves. Our model light curves herein show what one can expect when interacting transiting systems are discovered, for example, with ongoing radial velocity surveys or with space-based missions such as Kepler (Borucki et al. 2003) or COROT (Baglin 2003), and allow for a more complete understanding of the multiplanet dynamics.  The plan of this paper is as follows. In \\S 2, we describe our dynamical model of the system, and illustrate some of the photometric properties that the in-transit light-curves would be expected to display. In \\S 3, we describe an evaluation of the Doppler radial velocity measurements that have been taken during predicted transit intervals. In \\S 4, we evaluate a range of archival and newly obtained photometric data sets for evidence of transits. In \\S 5, we discuss our results, as well as the current status and outlook for the ongoing {\\it TransitSearch} collaboration, and in \\S 6, conclude. ",
        "conclusions": "Our study of Gl 876 has produced several interesting results. We have demonstrated that transits of planets b and c are not currently occurring, and we have thus modestly constrained the parameter space available to observationally allowed orbital configurations of the Gl 876 system. The absence of transits is consistent with the dynamical study of Rivera et al. (2005) who find that a co-planar system inclination $i\\sim50^{\\circ}$ yields the best fit to the radial velocity data. Our dynamical analysis indicates that when planetary transits in a strongly interacting system similar to Gl 876 are eventually observed, then a tremendous amount of information will be obtained. We look forward to the discovery of the first such system.  We have also shown that a globally distributed network of small-telescope observers can effectively provide definitive photometric follow-up of known planet-bearing low-mass stars. Our experience with Gl 876 indicates that intensive photometric monitoring campaigns of individual, nearby M dwarfs constitute a tractable strategy for planet detection. The Sun's nearest low-mass stellar neighbors should therefore be scrutinized very carefully for transits by short-period terrestrial planets. A low-budget survey of this type provides a small, but nevertheless non-negligible chance to discover a truly habitable world from the ground, and within a year or two. Such a detection would be a remarkable and exciting discovery.  \\subsection{Acknowledgements} We thank K. Johnston, R. Gaume, K. Malatesta, J. Lazio, D. Boboltz, C. Shaw, B. Casey, E. Albert, A. Henden, J. Larsen, D. Katz, C. Phillips, W. Bollwerk, Z. Dugan, J. Jaglowicz, and T. Castellano for useful discussions or other support.  We are grateful to the AAVSO and its many members who provided campaign support in addition to those authors listed above. We also gratefully acknowledge the insightful suggestions and constructive feedback offered by our referee.  This material is based upon work supported by the National Aeronautics and Space Administration under Grant NNG04GN30G issued through the Origins of Solar Systems Program to GL. GWH acknowledges support from NASA grant NCC5-511 and NSF grant HRD-9550561. This research has made use of the Simbad, USNO and {\\it TransitSearch} databases, operated at CDS, Strasbourg, France, USNO, Washington, D.C., USA, and UCSC, Santa Cruz, California, USA, respectively.  \\clearpage"
    },
    "0608/astro-ph0608440_arXiv.txt": {
        "abstract": "{Based on  \\ion{O}{2} recombination lines we present a new calibration (called \\ion{O}{2}$_{RL}$) of Pagel's $O_{23}$ indicator to determine the O/H abundance ratio in extragalactic \\ion{H}{2} regions and emission line galaxies. The \\ion{O}{2}$_{RL}$ calibration produces O/H abundances about a factor of two higher than those derived from the $T$(4363) method with $t^2$ = 0.00. The \\ion{O}{2}$_{RL}$ calibration has implications for the study of different properties of emission line galaxies such as their metallicity, star formation rate, and initial mass function. The \\ion{O}{2}$_{RL}$ calibration also affects the abundance determinations based on other O/H indicators, that include collisionally excited lines, like those known as $O_3N_2$, $N_2$, $S_{23}$, $Ar_3O_3$, and $S_3O_3$. We argue that the controversy between the $T$(4363) method and the photoionization models method to derive O/H values is mainly due to temperature variations inside the observed \\ion{H}{2} regions.}   \\resumen{Presentamos una nueva calibraci\\'on del indicador $O_{23}$ de Pagel para determinar los cocientes de O/H en regiones \\ion{H}{2} extragal\\'acticas y galaxias con l\\'i{}neas de emisi\\'on. Esta calibraci\\'on la llamamos \\ion{O}{2}$_{RL}$ y esta basada en l\\'i{}neas de recombinaci\\'on de \\ion{O}{2}. Nuestra calibraci\\'on produce abundancias de O/H alrededor de un factor de dos mayores que las obtenidas a partir del m\\'etodo $T_e$(4363) con $t^2$ = 0.00. La calibraci\\'on \\ion{O}{2}$_{RL}$ tiene implicaciones para el estudio de diferentes propiedades de las galaxias con l\\'i{}neas de emisi\\'on tales como la metalicidad, la tasa de formaci\\'on estelar, y la funci\\'on inicial de masas. La calibraci\\'on \\ion{O}{2}$_{RL}$ tambi\\'en afecta aquellas determinaciones de abundancias basadas en otros indicadores de O/H, que incluyen l\\'i{}neas excitadas colisionalmente, tales como los llamados $O_3N_2$, $N_2$, $S_{23}$, $Ar_3O_3$, y $S_3O_3$. Argumentamos que la controversia entre el m\\'etodo $T$(4363) y el m\\'etodo basado en modelos de fotoionizaci\\'on para determinar O/H se debe principalmente a la presencia de variaciones de temperatura dentro de las regiones \\ion{H}{2} observadas.}  \\addkeyword{Galaxies: Abundances} \\addkeyword{H~II regions: Abundances}   \\begin{document} ",
        "introduction": "\\label{sec:intro}    The advent of large telescopes is permitting to observe \\ion{H}{2} regions in galaxies many tens of megaparsecs away from us and emission line galaxies up to distances of $z\\sim 3$. But since the amount of photons that we obtain from faraway objects is small, often we have reliable information only for a few bright emission lines which has led to the idea of using different metallicity indicators based on at most a handful of bright emission lines.  The most popular metallicity indicator was introduced by Pagel et al. (1979, see also Edmunds \\& Pagel 1984) and is indistinctly known as Pagel's, or $R_{23}$, or $O_{23}$ indicator, where $O_{23}\\equiv I([$\\ion{O}{2}$]\\lambda 3727 + [$\\ion{O}{3}$]\\lambda\\lambda 4959, 5007)/ I({\\rm H}\\beta)$. The $O_{23}$ indicator has been calibrated with the O/H values based on two different methods: by using photoionization models, that we will call $PIM$ calibrations or $PIM$ method, and by using observational determinations of the O/H abundances based on the electron temperature derived from the $I(4363)/I(5007)$ [$\\ion{O}{3}$] ratio together with  the $I(3727)/I({\\rm H}\\beta)$ and the $I(5007)/I({\\rm H}\\beta)$ line ratios, the so called $T(4363)$ method.   There are significant differences between the calibrations of Pagel's indicator based on models (e. g. McCall et al. 1985; Dopita \\& Evans 1986; McGaugh 1991; Kobulnicky \\& Kewley 2004) and the calibrations based on observations and $T(4363)$ (e. g. Torres-Peimbert, Peimbert, \\& Fierro 1989; Pilyugin 2000, 2003; Castellanos, D\\'{\\i}az, \\& Terlevich 2002; Pilyugin, \\& Thuan 2005). The differences in the O/H values are in the $0.2$ - $0.6$ dex range and could be due mainly to the presence of temperature inhomogeneities over the observed volume (e. g. Campbell 1988; Torres-Peimbert et al. 1989; McGaugh 1991; Roy et al. 1996; Luridiana et al. 1999; Kobulnicky, Kennicutt, \\& Pizagno 1999: Kobulnicky, \\& Kewley 2004). These differences need to be sorted out if we want to obtain absolute accuracies in O/H of the order of 0.1 dex or better. We will call these differences the calibration controversy.  In this paper we present a qualitatively different calibration of the $O_{23}$ indicator that is based on the intensity ratio of \\ion{O}{2} recombination lines to  \\ion{H}{1}  recombination lines, that we will call the \\ion{O}{2}$_{RL}$ method. As of now the \\ion{O}{2}$_{RL}$ calibration has been established only for objects in the high metallicity branch of the $O_{23}$ versus O/H relation and for log $O_{23}$ $>$ 0.5.  Preliminary discussions of the \\ion{O}{2}$_{RL}$ calibration were presented by Peimbert \\& Peimbert (2003, 2005). ",
        "conclusions": "\\label{sec:conc}  The \\ion{O}{2}$_{RL}$ method supports the suggestion that the controversy produced by the relatively high O/H values predicted by the $PIM$ calibrations and the relatively low O/H values predicted by the  $T(4363)$ calibrations are mainly due to temperature variations.  The best way to calibrate the $O_{23}$ indicator is to use the \\ion{O}{2}$_{RL}$ method to obtain the O/H values because it is independent of the temperature structure.  The use of $T$(4363) values to derive O/H, under the assumption of constant temperature, provides a lower limit to the O/H abundance ratios.  Since the nebular lines are less sensitive to $T_0$ and $t^2$ than the auroral lines, the model calibrations that adjust the nebular lines are closer to the \\ion{O}{2}$_{RL}$ calibration than those derived using the observed $T$(4363) values.  For a given object the \\ion{O}{2} recombination lines provide gaseous abundances that are about 0.2 to 0.3 dex higher than those derived from collisionally excited lines and $T$(4363) under the assumption that $t^2$ = 0.00.  By using the  \\ion{O}{2} recombination lines to derive O/H abundances in Galactic \\ion{H}{2} regions together with state of the art Galactic chemical evolution models we obtain an excellent agreement with the O/H solar value derived by Asplund et al.(2005). We also find an excellent agreement between the \\ion{H}{2} region O/H abundances derived from recombination lines and those derived from young G dwarfs, this comparison is independent of Galactic chemical evolution effects since both types of objects correspond to the present abundances in the ISM of the solar vicinity.  With the \\ion{O}{2}$_{RL}$ calibration and observations of emission line galaxies at different $z$ values it will be possible to study the chemical evolution of the universe as a whole. All the other O/H indicators, like those known as $O_3N_2$, $N_2H\\alpha$, $S_{23}$, $Ar_3O_3$, and $S_3O_3$, depend on an absolute calibration and we recommend to calibrate them using the \\ion{O}{2}$_{RL}$ method.   Models of the chemical evolution of N/O versus O/H as a function of time for different types of galaxies are required to calibrate the $N_2$ and $O_3N_2$ indicators.  We need more high resolution observations of \\ion{O}{2} recombination lines in the local universe to refine and extend the \\ion{O}{2}$_{RL}$ calibration. We need objects with log $O_{23}$ smaller than 0.5 in the high metallicity branch. We also need  objects of different $P$ values at a given $O_{23}$ value in the high metallicity branch."
    },
    "0608/astro-ph0608395_arXiv.txt": {
        "abstract": "{To investigate the formation of sub-stellar objects, we observed a sample of ultracool dwarfs members of the Upper Scorpius OB association.} {The properties of disks, such as their composition, life-time, frequency, and how it compares to the accretor frequency, provide important clues on the mechanisms responsible for the formation of stellar and sub-stellar objects.} {We report the results of mid-IR observations with VISIR at the VLT of 10 ultracool dwarfs members of the nearby Upper Scorpius OB association in four filters ranging between 8.59 (PAH1) to 12.8~$\\mu$m (Ne~II), and one brown dwarf with Spitzer between 3.6 and 24~$\\mu$m.} {Seven of our targets are detected in at least one of the bands, and we derive upper limits on the fluxes of the remaining 4. These results combined with previous studies from the literature lead to an improved disk frequency of 50$\\pm$12\\%. This frequency is significantly higher than that of accretors (16.3\\%$\\pm$6.2\\%). Only one object showing mid-IR excess also has H$\\alpha$ emission at a level indicating that it must be accreting. Four of the detected targets are multiple system candidates.} {The observed disk frequency for sub-stellar objects in the Upper Scorpius association is similar to that of stars, consistent with a common formation scenario. It is also similar to the disk fractions observed in younger clusters, suggesting that the disk lifetimes might be longer for ultracool dwarfs than for higher-mass stars. } ",
        "introduction": "Over the last ten years, a large number of brown dwarfs (BD) and very low mass stars have been discovered in the field, in star forming regions and in clusters. One of the most debated questions is whether sub-stellar objects form the same way as stars, via the fragmentation and collapse of a molecular cloud, or as aborted stellar embryos ejected from their parent proto-stellar cluster. A comparison of stellar and sub-stellar disk properties directly addresses this question and provides crucial informations and straightforward constraints on the models of formation. The first evidences showing that brown dwarfs have disks have been given several years ago \\citep{1999ApJ...525..440L,1998A&A...335..522C} and since then substantial fractions of young brown dwarfs in star forming regions have been reported to have circumstellar material \\citep[see e.g.][]{2005ApJ...635L..93L,2005ApJ...620L..51L,2005ApJ...631L..69L,2005ApJ...626..498M, 2004A&A...427..245S,2003AJ....126.1515J,2003ApJ...585..372L,2001ApJ...558L..51M, 2001A&A...376L..22N} and to undergo an accretion phase \\citep[e.g ][]{2002ApJ...578L.141J, 2004A&A...424..603N,2004ApJ...604..284B,2005ApJ...626..498M,2005ApJ...625..906M}. From their near-IR observations,  \\citet{2003AJ....126.1515J} estimated an initial frequency of disk around BD and very low mass stars of 50\\%$\\pm$25\\% in various star forming regions. Using Spitzer to study the population of the IC348 young cluster, \\citet{2006AJ....131.1574L} estimated disk fractions of 47\\% for K6--M2 stars, and 28\\% for M2--M6 stars, while \\citet{2005ApJ...631L..69L} reported 42\\% for members later than M6. Although these results seem to show that BD originate like stars more than like ejected or evaporated embryos, our understanding of the formation processes is far from complete. It is still unclear whether one mechanism eventually prevails over the other, and how the contribution of each mechanism to the overall final population depends on the environment and initial conditions.  The Upper Scorpius OB association (hereafter USco) is an ideal place in which to study the properties of ultracool dwarf disks. It is close and young enough \\citep[145~pc and 5~Myr,][]{1999AJ....117..354D} that even its sub-stellar members have relatively bright luminosities, and the sample of known sub-stellar members is one of the largest for that type of association. Moreover, by comparing the disk properties of the 5~Myr USco to those of younger associations such as IC348 \\citep{2006AJ....131.1574L,2005ApJ...631L..69L} or Cha I \\citep{2005ApJ...631L..69L}, it is possible to study the disk evolution and lifetimes. We have obtained mid-infrared (mid-IR) photometry of 10 ultracool dwarf members of the association using VISIR on the VLT, and retrieved mid-IR photometry of an additional USco brown dwarf from the Spitzer public archive. In section \\ref{observations}, we describe the observations, data reduction and analysis, and in section \\ref{analysis}, we discuss the derived disk frequency and compare to previous studies. ",
        "conclusions": "We report mid-IR observations of 11 ultracool dwarfs members of the Upper Scorpius association. Seven of them are detected in at least one of the mid-IR images, and 3 show a clear excess indicating the presence of circumstellar material. The relatively high frequency of disks (50$\\pm$12\\%) found among Upper Scorpius ultracool dwarfs (5~Myr) is comparable to those reported by e.g \\citet{2003AJ....126.1515J, 2005ApJ...631L..69L} for a variety of younger associations (1--2~Myr), possibly indicating that the timescales of disk decay are longer for ultracool dwarfs than for solar-mass stars. Although the sample we present is too small to draw any statistically meaningful conclusion, there seems to be no significant difference between the disk frequency of single and multiple ultracool dwarfs. Since these relatively wide systems have most likely formed like stars (they are indeed very unlikely to have survived a gravitational ejection), this similarity between single and multiple objects might be an additional clue that most ultracool dwarfs must form like stars rather than ejected embryos. Further studies on larger samples of single and multiple systems are required to confirm these observations. Finally, only one of the objects displaying mid-IR excess also shows evidence of accretion, indicating that the frequency of transition disks in Upper Scorpius must be signiticantly higher than that of accretion disks."
    },
    "0608/astro-ph0608110_arXiv.txt": {
        "abstract": "We derive a new and user friendly simple analytic approximation for determining the effect of intergalactic absorption on the \\gray spectra of TeV blazars the energy range $0.2 ~{\\rm TeV} < E_{\\gamma} < 2$ TeV and the redshift range $0.05 <  z < 0.4$.   In these  ranges, the form  of the absorption    coefficient    $\\tau(E_{\\gamma})$    is    approximately logarithmic.  The  effect of  this  energy  dependence  is to  steepen intrinsic  source  spectra such  that  a  source  with an  approximate power-law  spectral index  $\\Gamma_{s}$ is  converted to  one  with an observed spectral index $\\Gamma_{o} \\simeq \\Gamma_{s} + \\Delta \\Gamma (z)$ where $\\Delta  \\Gamma (z)$  is a linear function  of $z$ in the redshift range 0.05 --  0.4.  We apply this approximation to the spectra of 7 TeV blazars. ",
        "introduction": "Stecker, Malkan \\& Scully (2006)  (SMS) have used recent {\\it Spitzer} observations (Le  Floc'h \\etal 2005, Perez-Gonzalez  \\etal 2005) along with other data on  galaxy luminosity functions and redshift evolution in order  to make  a detailed evaluation  of the  intergalactic photon density as a function of both energy  and redshift for $0 < z < 6$ for photon energies from .003 eV to the Lyman limit cutoff at 13.6 eV in a $\\Lambda$CDM universe with $\\Omega_{\\Lambda}  = 0.7$ and $\\Omega_{m} = 0.3$. They  then used their calculated  intergalactic photon densities to calculate  the optical  depth of the  universe, $\\tau$,  for \\grays having energies from 4 GeV to  100 TeV emitted by sources at redshifts from  ~0  to  5.  They  also  gave  a  parametric fit  with  numerical coefficients for $\\tau(E_{\\gamma},z)$.  As  an   example  of  the   application  of  the   detailed  numerical determination  of  the optical  depth,  they  calculated the  absorbed spectrum of  the blazar PKS  2155-304 at $z  = 0.117$ and  compared it with  the  spectrum observed  by  the  H.E.S.S.   air Cherenkov  \\gray telescope  array.  It  was  noted  that a  steepening  in an  $E^{-2}$ power-law differential  photon source spectrum for  PKS 2155-304 would be  steepened by  approximately  one  power to  $E^{-3}$  as would  be produced by a \\gray opacity with a logarithmic energy dependence.  The purpose of this letter is to generalize this result by determining fits for approximating  $\\tau(E_{\\gamma}, z)$ by logarithmic functions in  order to predict  power-law steepenings  in assumed  blazar source spectra of the form $E^{-\\Gamma_{s}}$  to observed spectra of the form $E^{-\\Gamma_{o}}$ where $\\Gamma_{o} \\simeq \\Gamma_{s} + \\Delta \\Gamma$ and $\\Delta \\Gamma$ is determined to be a linear function of $z$. ",
        "conclusions": "We have  derived a new and user friendly simple  analytic approximation for determining the effect of intergalactic absorption on  the spectra of \\gray sources in the energy  range 0.2 TeV $ <  E_{\\gamma} <$ 2 TeV and the redshift range $0.05 < z < 0.4$, where the absorption is primarily from interactions with intergalactic photons in the optical to near infrared wavelength range.  In these energy and redshift ranges, the form of the absorption coefficient $\\tau(E_{\\gamma})$ is approximately logarithmic. The effect of this  energy dependence is  to steepen  the intrinsic  source spectrum such  that  a source  with  an  approximate  power-law spectral  index $\\Gamma_{s}$  is converted  to  one with  an  observed spectral  index $\\Gamma_{o}  \\simeq \\Gamma_{s} +  \\Delta \\Gamma$  in the  energy range 0.2 -- 2 TeV,  where $\\Delta \\Gamma (z)$  is a linear function  of $z$ in the redshift  range 0.05 -- 0.4.  We  have applied this  approximation to the  spectra of  seven TeV  blazars.  These  power-law approximations, both observational  and theoretical, are  only useful over  the limited energy  range indicated,  {\\it viz.}  0.2--2 TeV.   The  actual spectra should exhibit  some curvature. At higher energies  the spectra should cut off either because of a  natural upper limit to the source spectra or, more  likely in the case  of high frequency peaked  BL Lac objects (HBLs),  because  of the increase  in  optical depth  with energy owing to interactions with the much more numerous intergalactic far infrared photons. A full numerical treatment of absorption can be made using the exact form of  the optical depth as a function  of redshift and energy as calculated in  the detailed work of SMS. Examples of  more  detailed  treatments  of  effects  of  intergalactic  \\gray\\ absorption are given for Mkn 501  and Mkn 421 by Konopelko \\etal (2003) and for PKS 2155-304  by SMS. However,  because of the limited energy range of empirical TeV spectral determinations, and because of  the  significant   statistical  and  systematic  uncertainties  in the empirical  determinations  of spectral  indices  of  TeV sources,  our approximate linear relations  can be quite useful in  making rapid and simple analyses of TeV source spectra."
    },
    "0608/astro-ph0608326_arXiv.txt": {
        "abstract": "Based  on  a  new  approach  on modeling  the  magnetically  dominated outflows  from AGNs  (Li et  al. 2006),  we study  the  propagation of magnetic tower jets in gravitationally stratified atmospheres (such as a galaxy  cluster environment)  in large scales  ($>$ tens of  kpc) by performing  three-dimensional  magnetohydrodynamic (MHD)  simulations. We present  the detailed  analysis of the  MHD waves,  the cylindrical radial force balance,  and the collimation of magnetic  tower jets. As magnetic energy is injected into  a small central volume over a finite amount of time, the magnetic fields expand down the background density gradient, forming a collimated jet and an expanded ``lobe'' due to the gradually decreasing background density and pressure. Both the jet and lobes  are magnetically  dominated.   In addition,  the injection  and expansion produce  a hydrodynamic shock  wave that is moving  ahead of and enclosing the magnetic tower  jet. This shock can eventually break the hydrostatic equilibrium  in the ambient medium and  cause a global gravitational  contraction.    This  contraction  produces   a  strong compression  at the  head of  the magnetic  tower front  and  helps to collimate radially to produce a slender-shaped jet.  At the outer edge of the jet, the magnetic pressure is balanced by the background (modified) gas pressure, without any significant contribution from the hoop stress.   On the other hand,  along the central axis  of the jet, hoop stress is  the dominant force in shaping  the central collimation of the poloidal  current.  The system, which possesses  a highly wound helical  magnetic  configuration,  never  quite reaches  a  force-free equilibrium  state though the  evolution becomes  much slower  at late stages.    The  simulations   were  performed   without   any  initial perturbations  so the  overall  structures of  the  jet remain  mostly axisymmetric. ",
        "introduction": "A number of astronomical systems have been discovered to eject tightly collimated and hyper-sonic plasma  beams and large amounts of magnetic fields  into the interstellar,  intracluster and  intergalactic medium from the  central objects  during their initial/final  (often violent) stages.  Magnetohydrodynamic  (MHD) mechanisms are  frequently invoked to  model the  launching, acceleration  and collimation  of  jets from Young Stellar  Objects (YSOs), X-ray binaries  (XRBs), Active Galactic Nuclei  (AGNs),  Microquasars,  and  Quasars (QSOs)  \\citep[see,  {\\it e.g.},][and references therein]{M01}.  Theory of  magnetically driven outflows in  the electromagnetic regime has been proposed by  \\citet[]{B76} and \\citet[]{L76} and subsequently applied  to  rotating black  holes  \\citep[]{BZ77}  and to  magnetized accretion  disks   \\citep[]{BP82}.   By  definition,   these  outflows initially are dominated by electromagnetic forces close to the central engine.   In  these  and  subsequent  models  of  magnetically  driven outflows (jets/winds), the plasma velocity passes successively through the  hydrodynamic  (HD)   sonic,  slow-magnetosonic,  Alfv\\'enic,  and fast-magnetosonic critical surfaces.  The  first  attempt  to  investigate  the  nonlinear  (time-dependent) behavior  of magnetically  driven  outflows from  accretion disks  was performed by \\citet[]{US85}.  The  differential rotation in the system (central star/black  hole and the  accretion disk) creates  a magnetic coil  that simultaneously  expels and  pinches some  of  the infalling material.  The buildup of  the azimuthal (toroidal) field component in the  accretion disk  is released  along  the poloidal  field lines  as large-amplitude   torsional   Alfv\\'en   waves  (``sweeping   magnetic twist'').   After  their  pioneering   work,  a  number  of  numerical simulations  to study  the MHD  jets have  been done  \\citep[see, {\\it e.g.},][and  references   therein]{F98}.   An  underlying  large-scale poloidal field  for producing the  magnetically driven jets  is almost universally  assumed in  many theoretical/numerical  models.  However, the origin and  existence of such a galactic  magnetic field are still poorly understood.  In   contrast   with  the   large-scale   field  models,   Lynden-Bell \\citep[]{LB94,  L96, L03,  L06} examined  the expansion  of  the local force-free magnetic loops anchored to  the star and the accretion disk by using  the semi-analytic approach.  Twisted magnetic  fluxes due to the disk rotation make the magnetic  loops unstable and splay out at a semi-angle 60\\degr  \\ from  the rotational axis  of the  disk.  Global magnetostatic  solutions  of  magnetic  towers with  external  thermal pressure were also computed by \\citet[]{Li01} using the Grad-Shafranov equation  in axisymmetry \\citep[see  also,][]{L02, LR03,  UM06}.  Full MHD numerical  simulations of magnetic  towers have been  performed in two-dimension  (axisymmetric) \\citep[]{R98, T99,  U00, K02,  K04a} and three-dimension \\citep[]{K04b}.  Magnetic  towers are also observed in the laboratory experiments \\citep[]{HB02, L05}.  This  paper describes  the  nonlinear dynamics  of propagating  magnetic tower jets in large scales  ($>$ tens of kpc) based on three-dimensional MHD  simulations.  We  follow closely  the approach  described in  Li et al. (2006; hereafter Paper I). Different from Paper I, which studied the dynamics of magnetic field evolution  in a uniform background medium, we present  results  on  the  injection  and the  subsequent  expansion  of magnetic  fields  in  a  stably  stratified background  medium  that  is described  by  an  iso-thermal  King  model  \\citep[]{K62}.   Since  the simulated  magnetic structures  traverse  several scale  heights of  the background medium, we  regard that our simulations can  be compared with the radio sources inside the galaxy cluster core regions. Due to limited numerical  dynamic range, however,  the injection  region (see  Paper I) assumed in this  paper will be large  (a few kpc).  Our goal  here is to provide the  detailed analysis of the  magnetic tower jets,  in terms of its  MHD wave  structures,  its cylindrical  radial  force balance,  and collimation.  The paper is organized as follows: In \\S 2, we outline the model  and numerical  methods.   In  \\S 3,  we  describe the  simulation results.  Discussions and conclusions are given in \\S 4 and \\S 5. ",
        "conclusions": "By performing 3-D  MHD simulations we have investigated  in detail the nonlinear dynamics of magnetic tower jets, which propagate through the stably  stratified   background  that  is   initially  in  hydrostatic equilibrium.   Our simulations,  based  on the  approach developed  in \\citet{Li06}, confirm a number of the global characteristics developed in  Lynden-Bell  (1996, 2003)  and  Li  et  al. (2001).   The  results presented here give additional  details for a dynamically evolving jet in a  stratified background.   The magnetic tower  is made  of helical magnetic fields, with poloidal  flux and poloidal current concentrated around  the central axis.  The ``returning''  portion of  the poloidal flux and  current is  distributed on the  outer shell of  the magnetic tower. Together they form a self-contained system with magnetic fields and associated  currents, being confined  by the ambient  pressure and gravity.  The  overall morphology  exhibits  a  confined magnetic  jet ``body'' with  an expanded ``lobe''.  The confinement of  the ``body'' comes  jointly  from  the  external  pressure  and  the  gravity.  The formation of the lobe is due  to the expansion of magnetic fields in a decreasing background pressure.  A  hydrodynamic shock  wave  initiated by  the  injection of  magnetic energy/flux propagates ahead  of the magnetic tower and  can break the hydrostatic  equilibrium  of  the  ambient medium,  causing  a  global gravitational contraction.   As a result, a  strong compression occurs in  the axial  direction  between  the magnetic  tower  front and  the reverse slow-mode MHD shock wave front that follows.  Furthermore, the magnetic tower  jets are deformed radially into  a slender-shaped body due to the inward-directed flow of the ambient (non-magnetized) gas.  The  lobe is  magnetically dominated  and  is likely  filled with  the toroidal magnetic  fields generated  by the central  poloidal current. At the edge  of the magnetic tower jet,  the outward-directed magnetic pressure  gradient  force  is  balanced  by  the  inward-directed  gas pressure gradient force, so the radial width of the magnetic tower can be determined jointly by the magnitude of the poloidal current and the magnitude  of the  external gas  pressure.  The  highly  wound helical magnetic  field in  the magnetic  tower never  reaches  the force-free equilibrium precisely, but obtains radial force-balance, including the gas pressure gradient inside the magnetic tower.  The stability  of the  magnetic tower jets  will be considered  in our forthcoming papers."
    },
    "0608/astro-ph0608056_arXiv.txt": {
        "abstract": "{} { In order to estimate the contribution of low surface brightness (LSB) galaxies to the local ($\\rm{z\\le 0.1}$) galaxy number density, we performed an optical search for LSB candidates in a 15.5 deg${^{2}}$  part of the region covered by the 65 deg${^{2}}$ blind Arecibo HI Strip Survey (AHISS). } { Object detection and galaxy profile fitting were done with analytical algorithms. The detection efficiency and the selection effects were evaluated using large samples of artificial galaxies.} { Our final catalogue is diameter-limited and contains 306 galaxies with diameters $>18\\arcs$ at the limiting surface brightness of $\\mu_{B,lim}=25.2\\pm0.31\\,$\\magarcsec.  Of these 306 galaxies, 148 were not catalogued previously. Our results indicate that low surface brightness galaxies contribute at least to 30\\,\\% to the local galaxy number density.} { Without additional distance information, choosing the limiting diameter and the surface brightness at which the diameter is measured is crucial. Depending on these choices, diameter-limited optical catalogues are either biased against LSB galaxies, or contaminated with cosmologically dimmed high surface brightness galaxies, which affects the implied surface brightness distribution. The comparison to the AHISS showed that although optical surveys detect more galaxies per deg$^{2}$ than HI surveys, their drawback is the need for spectroscopic follow up observations to derive distances. Blind HI surveys have no diameter limits, but tend to miss gas-poor galaxies and all galaxies which lie outside their redshift limits. HI and optical surveys thus provide complementary information and sample different parts of the LSB galaxy population. } ",
        "introduction": "The volume density of low surface brightness (LSB) galaxies has been underestimated for a long time, as they are quite underrepresented in many earlier optical catalogues due to strong selection effects against their detection \\citep[a detailed description of these effects can be found in e.g.,][]{impey-1997-review}. These selection effects resulted in the Freeman law \\citep{freeman-1970} indicating that spiral galaxies have a typical inclination corrected central disc surface brightness with a relatively small dispersion (${\\rm\\mu_{B,corr}=(21.65\\pm0.3)\\,mag\\,arcsec^{-2}}$). \\cite{disney-1976} suggested that there could be a large population of galaxies which was undiscovered by most of the optical surveys due to selection effects. Subsequent studies showed that there is indeed a large population of galaxies with central surface brightnesses much fainter than the Freeman value \\citep[e.g.,][]{impey-1988, schombert-1988}. More recent estimates were able to show that LSB galaxies account for a significant fraction of the total galaxy numbers \\citep[see e.g.,][]{mcgaugh-1995, mcgaugh-1996, impey-1997-review, bothun-1997, oneil-bothun-2000}. \\cite{minchin-2004} estimate the contribution (by numbers) of LSBs to gas-rich ($M_{HI}>10^8\\,M_{\\odot}$) galaxies to be 50-60 per cent (using two different methods for the estimate). Following \\cite{minchin-2004}, LSBs contribute approximately 30$\\pm$10\\,\\% to the neutral hydrogen density.\\\\ However, the classification of galaxies into high and low surface brightness galaxies (HSBs \\& LSBs) is not a strict separation. In this paper, we will classify all galaxies with a central, inclination corrected blue surface brightness of ${\\rm\\mu_{B,corr}>22.5\\,mag\\,arcsec^{-2}}$ (i.e., $\\ge 3\\sigma\\,$ fainter than the Freeman value) as low surface brightness galaxies.\\\\ As stated before, optical surveys show a severe selection effect regarding the detection of LSB galaxies. As this optical selection bias does not apply to blind HI surveys, they are considered a good alternative in the search for LSB galaxies. Moreover, they allow a different probe of the galaxy population. However, HI surveys will also show some kinds of selection effects. Although LSBs are often regarded as gas-rich, having total HI masses which are comparable to HSB galaxies \\citep[e.g.,][]{deblok-1996b, burkholder-2001, oneil-2004a}, blind HI surveys will miss gas poor galaxies. As there is a trend to low HI column densities for LSB galaxies \\citep[][]{deblok-1996b, minchin-2003}, very deep HI surveys are needed to be able to easily detect LSBs. For column densities below $\\rm{n_{HI}\\leq10^{19.7}cm^{-2}}$, ionisation of the neutral hydrogen may become important \\citep[][]{sprayberry-1998} and the amount of HI, and therefore also the detectability in HI, will decrease rapidly. And HI surveys have a limit in bandwidth, and thus in radial velocity range. Thus, it is clear that both HI and optical surveys each in their own way are biased to some extent against the detection of LSB galaxies and that both kind of surveys will result in a different sampling of the LSB population\\\\ A blind optical follow-up observation of a region of the sky which was initially observed in the 21 cm line allows one to make a direct comparison between an optically and an HI selected sample. We have therefore made a blind optical follow-up observation of a part of the region covered by the blind Arecibo HI Strip Survey \\citep[AHISS,][]{zwaan-1997}.\\\\ The paper is organised as follows. Section 2 presents our data and the data reduction. In section 3, the object detection, the catalogue handling and the selection of the sample are described. Additionally, this section deals with the galaxy profile fitting which was done using Galfit \\citep[][]{peng-2002}. We estimate the detection efficiency and the bias due to our selection criteria in section 4, using intensive studies on artificial galaxies. Section 5 deals with the comparisons to other optical surveys and to the AHISS itself. Finally, we present our conclusions in section 6. ",
        "conclusions": "We performed optical follow-up observations of a part of the AHISS, and detected optical counterparts of all HI detections. We estimated the detection efficiency of our survey using large samples of artificial galaxies. The comparison of our survey with two other optical surveys indicates that an automated search algorithm and an analytical size estimation may increase the number of galaxies per deg$^2$ compared to catalogues based on visual estimates. We show that diameter-limited catalogues created by automated routines are at least as complete as visually created ones.\\\\ Although our number density is higher than that of the AHISS (mainly due to the limited velocity range of AHISS), the fraction of LSBs in the AHISS is higher than in our survey (30-42\\,\\% in TBHD vs. 49\\,\\% in AHISS). We suggest that this is caused by our relatively large diameter limit, which reduces the contamination of the sample with cosmologically dimmed higher redshifted galaxies, but also rejects local LSBs. In order to reduce this bias against the selection of LSBs in optical surveys one needs redshifts for the whole sample. For a given observing time, optical surveys are despite all drawbacks best suited to detect a large number of galaxies. The information which optical surveys provide on detected galaxies is, however, fundamentally different from that of HI surveys. The radial velocity information inherent to HI surveys allows a direct estimation of the volume density and the rejection of high-z HSBs. However, they will miss the gas-poor LSBs and can only probe a limited velocity range. With current telescopes, only very massive galaxies ($M_{HI}>10^{10}\\,M_\\odot$) can be detected beyond 10\\,000 km\\,s$^{-2}$. For example, the 6$\\sigma$ HI mass limit of the ALFALFA survey will be $\\sim 9.5\\,\\times\\,10^9\\,M_{\\odot}$ at $cz=10\\,000\\,km/s$ \\cite[][]{giovanelli-2005}. Thus, HI and optical surveys are complementary, as are the differences between the samples of LSB galaxies they detect. Therefore, we especially need extremely deep surveys of both kinds to extend our - still limited - know\\-ledge about LSBs. That we have not yet reached \\emph{the end} can e.g., be seen from the very deep HBDP optical survey."
    },
    "0608/astro-ph0608260_arXiv.txt": {
        "abstract": "We determine stellar parameters for the M dwarf GJ~436 that hosts a Neptune-mass planet.  We employ primarily spectral modeling at low and high resolution, examining the agreement between model and observed optical spectra of five comparison stars of type, M0-M3.  Modeling high resolution optical spectra suffers from uncertainties in TiO transitions, affecting the predicted strengths of both atomic and molecular lines in M dwarfs.  The determination of T$_{\\rm eff}$, gravity, and metallicity from optical spectra remains at $\\sim$10\\%.  As molecules provide opacity both in lines and as an effective continuum, determing molecular transition parameters remains a challenge facing models such as the PHOENIX series, best verified with high resolution and spectrophotometric spectra.  Our analysis of GJ 436 yields an effective temperature of T$_{\\rm eff}$ = 3350 $\\pm$ 300 K and a mass of 0.44 \\msune.  New Doppler measurements for GJ~436 with a precision of 3 \\ms taken during 6 years improve the Keplerian model of the planet, giving a minimum mass, $M\\sin i$ = 0.0713 \\mjup = 22.6 \\mearth, period, $P$ = 2.6439 d, and $e$ = 0.16$\\pm$0.02. The noncircular orbit contrasts with the tidally circularized orbits of all close-in exoplanets, implying either ongoing pumping of eccentricity by a more distant companion, or a higher Q value for this low-mass planet.  The velocities indeed reveal a long term trend, indicating a possible distant companion. ",
        "introduction": "To date, radial velocity surveys have revealed three exoplanetary systems with M dwarf hosts: GJ 876, GJ 436, and GJ 581 \\citep{Rivera05, Butler04,Bonfils05b}.  Remarkably, all three of these M dwarfs host planets with minimum masses (\\msini) less than 0.1 \\mjupe, ranking them among the lowest mass exoplanets known. Microlensing surveys have strengthened the case for the preferntial occurence of such ``super-Earths'' around M dwarfs.  Current results suggest four M dwarf systems, with two harboring planets in the super-Earth mass range \\citep{Beaulieu06,Gould06, Bond04,Udalski05}.  Given detection biases against the discovery of low mass planets, all these findings imply that super-Earths are more common close companions to M dwarfs than gas giants \\citep{Endl06,Beaulieu06}. The frequency of longer period planets remains poorly constrained.  From a theoretical viewpoint, the growth of planets around low mass stars has been explored only recently.  Low mass stars likely form by accretion at the centers of protoplanetary disks having lower mass than those surrounding solar mass stars.  Such low mass disks may spawn low mass planets, both because of less available mass and because of shorter survival times for the disks (Laughlin et al 2004).  However, core accretion models of giant planet formation that include growth of a rocky core from dust particles followed by the gravitational accretion of gas suggest that giant planets may have enough time to form. Solutions to the time scale competition between planet growth and disk lifetime have been proposed by invoking both migration to move the planet to gas-rich areas and also by adopting lower opacities to allow faster accretion of the envelope onto the core (Hubickyj et al. 2004, Alibert et al. 2005).  Nonetheless, rocky cores that are starved of gas may instead become ice giants similar to Neptune and Uranus (Goldreich et al. 2004, Thommes et al. 2002, Ida \\& Lin 2005). Under the core accretion paradigm, it remains unknown how commonly planets fail to accrete gas after successful growth of a rock and ice core, thereby leaving behind a super-Earth planet. Most recently, \\citet{Boss06} provided a very different viable alternative to the core accretion scenario, showing that disk instability can also be invoked to explain current M dwarf observations.  Distinguishing between these competing models of planet formation requires more observational results.  To further understanding in this area, we present a follow-up study of GJ 436, an M dwarf recently discovered to host a Neptune-mass planet (\\msini = 0.066 \\mjup; Butler et al. 2004).  At present, little is known about this system.  As there is no other giant planet within a few AU of the host star, a larger planet cannot easily be invoked to explain a curtailed accretion of gas of the known planet, hereafter GJ~436b.  Furthermore, since the stellar luminosity and mass are not well known, the planet's minimum mass and predicted temperature are poorly constrained. To better characterize this system, we have obtained new Doppler measurements and attempted to deduce accurate stellar parameters. In \\S 2, we discuss photometric observations leading to estimates of the stellar mass and luminosity.  In \\S 3, we present an effort to determine the effective temperature, surface gravity, and metallicity of GJ 436 via low and high resolution spectral modeling. In \\S 4, new Doppler measurements are presented, leading to an improved orbit and minimum mass for the planet.  In \\S 5 and 6, we discuss the implications of our results. ",
        "conclusions": "The revised mass for the star GJ~436 of 0.44 \\msun and the revised orbital parameters and \\msini for the planet have tightened the constraints on the structure of this planetary system.  The minimum mass of the planet remains slightly greater than that of Neptune with \\msini = 22.6 \\mearth and orbiting with semimajor axis, $a$ = 0.0285 AU .  Two new results have emerged from the present analysis that render the system interesting and puzzling.  The eccentricity is definitely non-zero, with $e$ = 0.16$\\pm$0.02, the highest eccentricity for any exoplanet with an orbital period less than 4 d.  It has apparently avoided tidal circularization.  Moreover, the velocities exhibit a linear trend of 1.3 \\ms per year that appears to be real, indicating the presence of a more distant companion, the mass and orbit of which remain poorly constrained.  It is tempting to suppose that this outer companion is responsible for pumping the eccentricity of the inner planet.  Two scenarios seems possible.  In the first scenario, planet b resides in an eccentric orbit and the linear velocity slope is due to a companion so far away that it isn't exciting the eccentricity of planet b.  In this case, tidal theory would argue for a not-low value of Q and hence against a rocky planet.  Scaling the Earth's Q$\\sim$300 by the forcing period and scaling the Earth's radius by 22$^{1/3}$, one gets a tidal circularization time of 2 million years, not consistent with the age of the star that is certainly several billion years. It shows no signs of youth such as rapid rotation, or enhanced magnetic, chromospheric and coronal activity.  If one scales Neptune's $Q$ value, $Q\\sim 10^5$, by the forcing period and uses its radius, then the tidal circularization time is $2 \\times 10^9$ yr.  If the planet started on an eccentric orbit, then this timescale is plausible, as the orbit will have only partially circularized.  The $Q$ estimates for the ice giants in the solar system range from $10^4$ to a few $10^6$.  Such values provide an interesting constraint on $Q$ for a hot-Neptune.  In a second scenario, the planet b resides on an eccentric orbit, but the slope is due to a planet or binary star that comes close enough to pump the planet's eccentricity despite the tidal damping.  In this case, planet b can be terrestrial or an ice giant since it is continually being pumped.  But one wonders if such pumping is consistent with the velocity data. If we naively approximate the outer ``planet c'' to be on a circular orbit, then $m_c \\sin i \\sim$ 0.2 \\mjup (P/20 yr)$^{4/3}$ (slope / 1.25 m/s/yr), where a 20 yr minimum orbital time scale comes from $\\sim$4 times the duration of observations. Even if we take the duration of observations to be the minimum orbital period for planet c, then the ratio of semimajor axes is at least $\\sim$100.  With such large separation, the outer planet is unlikely to be effective at exciting an eccentricity.  To test this, we have performed numerical integrations in the secular octupole approximation (averaging over orbits and expanding in terms of ratio of semi-major axes, but not eccentricities, inclinations, or mass ratios).  We assume that this approximation gives a rough estimate of the mass of an outer planet in a coplanar system with $\\sin i$ = 1.  For an outer planet orbital period less than 14 years (and hence masses less than 0.12 \\mjup), the outer planet would need an eccentricity larger than 0.5 to be able to induce an eccentricity of 0.16 for planet b. Alternatively, an outer planet with orbital period of 25 years (and hence a mass of $\\sim$0.27 \\mjup) would need an eccentricity of only 0.2 to be able to induce the observed eccentricity of GJ~436b.  The timescale for the secular eccentricity perturbations is less than $10^6$ yr.  So this configuration would maintain the observed eccentricity of GJ~436b, regardless of the assumed composition and value of the tidal $Q$.  Given that we observe both an eccentricity and a slope, this scenario offers a reasonable explanation. Doppler observations during the upcoming years may reveal the mass and period of the outer planet, if it exists."
    },
    "0608/astro-ph0608274_arXiv.txt": {
        "abstract": "Extra-galactic radio sources are a significant contaminant in cosmic microwave background and Sunyaev-Zel'dovich effect experiments. Deep interferometric observations with the BIMA and OVRO arrays are used to characterize the spatial, spectral, and flux distributions of radio sources toward massive galaxy clusters at 28.5 GHz.  We compute counts of mJy source fluxes from 89 fields centered on known massive galaxy clusters and 8 non-cluster fields. We find that source counts in the inner regions of the cluster fields (within 0.5 arcmin of the cluster center) are a factor of $8.9^{+4.3}_{-2.8}$ times higher than counts in the outer regions of the cluster fields (radius greater than 0.5 arcmin). Counts in the outer regions of the cluster fields are in turn a factor of $3.3^{+4.1}_{-1.8}$ greater than those in the non-cluster fields. Counts in the non-cluster fields are consistent with extrapolations from the results of other surveys. We compute spectral indices of mJy sources in cluster fields between 1.4 and 28.5 GHz and find a mean spectral index of $\\alpha = 0.66$ with an rms dispersion of $0.36$, where flux $S \\propto \\nu^{-\\alpha}$. The distribution is skewed, with a median spectral index of $0.72$ and 25th and 75th percentiles of $0.51$ and $0.92$, respectively. This is steeper than the spectral indices of stronger field sources measured by other surveys. ",
        "introduction": "Extra-galactic radio sources are a significant contaminant in cosmic microwave background (CMB) and Sunyaev-Zel'dovich effect (SZE) experiments (e.g., see \\citet{holder2002, knox2004, tegmark2000}). Measurements of the CMB and SZE \\citep{sunyaev1970,sunyaev1972} have the potential to yield a wealth of cosmological information if foreground contaminants are well-understood. Models for the number counts of radio sources as a function of flux \\citep{dezotti2005, toffolatti1999, sokasian2001} have been derived from observations at lower frequencies and extrapolated to microwave frequencies.   Radio sources are often associated with the clusters of galaxies themselves. This is a potential source of bias for current and planned SZE surveys such as the Sunyaev-Zel'dovich Array (SZA)\\footnote{SZA website: http://astro.uchicago.edu/\\~sza}, the Arcminute Microkelvin Imager (AMI)\\footnote{AMI website: http://www.mrao.cam.ac.uk/telescopes/ami/}, the Atacama Pathfinder Experiment Sunyaev-Zel'dovich (APEX-SZ) survey\\footnote{APEX-SZ website: http://bolo.berkeley.edu/apexsz}, the South Pole Telescope (SPT)\\footnote{SPT website: http://spt.uchicago.edu/}, and the Atacama Cosmology Telescope (ACT)\\footnote{ACT website: http://www.hep.upenn.edu/act/}. To understand the impact on planned SZE cluster surveys, it is critical to characterize the spatial, spectral, and flux distribution of sources associated with clusters.  Contaminating radio emission from extra-galactic sources at frequencies  less than approximately 100 GHz is attributed to synchrotron radiation from active galactic nuclei (AGN) and star-forming galaxies. The AGN-powered radio galaxies dominate the source counts at high luminosities. At higher frequencies, the emission is attributed to dust emission from star-forming galaxies. Observations of radio sources toward galaxy clusters at very low frequencies ($\\lesssim 5$ GHz; e.g., \\citet{slee1983,slee1998, owen1996, ledlow1995, reddy2004}) show a strong central concentration of radio galaxies in clusters. The distribution of synchrotron-emitting star-forming galaxies is found to be less centrally-peaked in clusters \\citep{rizza2003} than that of AGN-powered radio galaxies. At our sensitivity level and observing frequency of 28.5 GHz, radio sources powered by AGN dominate and we refer to them henceforth as radio sources.  Several CMB experiments, such as WMAP \\citep{bennett2003}, DASI \\citep{kovac2002}, VSA \\citep{cleary2005}, and CBI \\citep{mason2003}, have measured microwave source counts as a function of flux for sources brighter than about 10 mJy. There have been two prior analyses of radio sources in SZE data taken with the Owens Valley Radio Observatory (OVRO)\\footnote{The OVRO mm-wave array is operated by Caltech with support from the National Science Foundation} and Berkeley-Illinois-Maryland Association (BIMA)\\footnote{The BIMA array is operated with support from the National Science Foundation} arrays. Using a sample of 56 fields centered on known massive galaxy clusters, \\citet{cooray1998a} computed counts and spectral indices of radio sources.  Using the outer regions of 41 cluster fields, \\citet{laroque2002} computed the normalization of source counts as a function of flux for faint sources in SZE data. The data used in \\citet{cooray1998a} and in \\citet{laroque2002} are subsets of the data presented in this paper.  In this paper we analyse faint ($\\sim$~mJy) radio sources found serendipitously toward massive galaxy clusters at 28.5 GHz from the OVRO/BIMA SZE imaging project.  Characterizing the spatial, spectral, and flux distribution at this relatively high frequency should help improve projections for radio source contamination in SZE and CMB experiments at frequencies of 30 GHz and higher.  We use 89 fields centered on known massive galaxy clusters and 8 non-cluster fields. The paper is organized as follows: Section 2 reviews the observations, data reduction, field selection, and the measured source fluxes. In Section 3, we compute spectral indices between 1.4 and 28.5 GHz using fluxes from our data and published 1.4 GHz surveys. In Section 4, we present source counts as a function of flux for cluster and non-cluster fields as well as the angular radial dependence of counts in cluster fields. We compare our results with those from other experiments and with theoretical models. ",
        "conclusions": "From deep interferometric observations at 28.5 GHz of unresolved radio sources toward 89 fields centered on massive galaxy clusters and 8 non-cluster fields, we find that differential source counts are greatly elevated in the centers of cluster fields. Counts are a factor of $8.9^{+4.3}_{-2.8}$ higher in central regions (r $\\leq 0.5$ arcmin) than in the outer regions (r $\\geq 0.5$ arcmin) of the cluster fields. Counts in our non-cluster fields are consistent with those expected from models and from extrapolations from other experiments. In addition, source counts in the outer regions of cluster fields are a factor of $3.3^{+4.1}_{-1.8}$ higher than counts in non-cluster fields.  Using the NVSS and other surveys, we find a mean spectral index for sources in cluster fields between 1.4 and $28.5\\,$GHz of $\\alpha = 0.66$ with rms dispersion of $0.36$, where flux $S \\propto \\nu^{-\\alpha}$. The distribution is skewed, with a median spectral index of $0.72$ [$0.51$ (25\\%), $0.92$ (75\\%)]. This is steeper than spectral indices of stronger ($\\sim$~20~mJy) field sources, and much steeper than those of much stronger ($\\sim$~Jy) field sources measured by other surveys. No significant differences are found between the distributions of the spectral indices of sources in the inner and outer regions of the clusters.  These results can be used for improving forecasts for radio source contamination of SZE and CMB experiments. The cluster fields used in this work were chosen to contain massive clusters and the sources were identified at $28.5\\,$GHz. We anticipate that the SZA and other instruments will be able to extend this work to less massive clusters and to sources identified at yet higher frequencies."
    },
    "0608/astro-ph0608042_arXiv.txt": {
        "abstract": "In this paper I briefly review recent progress in the field of ground-based gamma ray astrophysics. ",
        "introduction": "It is not my intention to review the historical development of ground-based gamma ray astrophysics or the experimental techniques employed in the field. Interested readers are referred to an excellent recent review on those subjects~\\cite{Weekes06}. Neither is it my intention to discuss every source that has been detected, because this is no longer practical. Instead, I will present a personal (thus biased and incomplete) perspective of the scientific impact of TeV observations on astronomy and physics. ",
        "conclusions": "TeV astrophysics has become a viable branch of modern astronomy. It provides a unique window to the extreme non-thermal side of the universe. Many classes of astronomical systems have been detected at TeV energies. The observations have not only shed new light on the properties of the systems themselves but on the physical processes operating in diverse astronomical settings. For instance, taken together, Blazars, microquasars, and gamma-ray bursts (though none has been detected at TeV energies yet~\\cite{Atkins05,MAGICGRB,Horan06}) may offer an excellent opportunity for us to make some tangible comparisons of the processes of particle acceleration and interaction in the jets of black holes over a vast range of physical scales (from microparsecs to megaparsecs)~\\cite{Cui05}. As the capability of TeV observatories improves, it is hopeful that more sources in the established classes and, more importantly, new classes of sources will be detected.  The field is of equally great interest to physicists, because it has made it possible to study some of the most important questions in physics at energies much beyond the capabilities of present and future particle accelerators. Independent of theoretical scenarios, TeV observations are capable of constraining the intrinsic spectrum of emitting particles and thus casting light on the nature of the particles and on the acceleration mechanisms. TeV observations have already begun to have a serious impact on modern cosmology. They have also provided insights into such fundamental issues as dark matter, primordial black holes, and Lorentz invariance. The constraints are expected to improve as the quality of data improves. In many cases, however, the challenge is to separate astronomical and physical origins of the TeV photons detected."
    },
    "0608/astro-ph0608568_arXiv.txt": {
        "abstract": "The supernova yields of r-process elements are obtained as a function of the mass of their progenitor stars from the abundance patterns of extremely metal-poor stars on the left-side [Ba/Mg]-[Mg/H] boundary with a procedure proposed by Tsujimoto and Shigeyama. The ejected masses of r-process elements associated with stars of progenitor mass $M_{ms}\\leq18M_{\\odot}$ are infertile sources and the SNe II with 20$M_{\\odot}\\leq M_{ms}\\leq 40M_{\\odot}$are the dominant source of r-process nucleosynthesis in the Galaxy. The ratio of these stars 20$M_{\\odot}\\leq M_{ms}\\leq40M_{\\odot}$ with compared to the all massive stars is about $\\sim$18\\%. In this paper, we present a simple model that describes a star's [r/Fe] in terms of the nucleosynthesis yields of r-process elements and the number of SN II explosions. Combined the r-process yields obtained by our procedure with the scatter model of the Galactic halo, the observed abundance patterns of the metal-poor stars can be well reproduced. ",
        "introduction": "\\label{intro} Elemental abundances in metal-poor Galactic halo stars are providing evidence of the earliest Galactic nucleosythesis history and clues about the identities of the first stellar generations, the progenitors (or predecessors) of the halo stars. The sample of identified field stars with [Fe/H]$\\leq$-2.5 has increased by more than an order of magnitude in the last decade. Because the elements in the atmospheres of these stars have been produced in a small number of nucleosynthetic events, abundance determinations can provide direct tests of model yields from different nuclear processes.  The elements heavier than the iron peak are made through neutron capture via two principal processes: the r-process and the s-process (Burbidge et al., 1957). The r-process (for rapid process) occurs when neutrons are added much more rapidly than the $\\beta$ decay times of the relevant nuclei. The site or sites of the r-process are not known, although suggestions include the $\\nu$-driven wind of Type II SNe (e.g., Woosly and Hoffman, 1992; Woosly et al., 1994) and the mergers of neutron stars (e.g., Lattimer and Schramm, 1974; Rosswog et al., 2000).  Particular attention to the Galactic evolution of elements produced by neutron-capture nucleosynthesis was given by Mathews et al., (1992), Pagel and Tautvaisiene (1997), and more recently by Travaglio et al. (1999). These authors adopted the standard approach to Galactic chemical evolution, assuming that stars form from a chemically homogeneous medium at a continuous rate. A more realistic model for the chemistry and dynamics of the gas is needed in order to investigate the earliest phases of halo evolution.  Recently, several chemical evolution models statistics on the build-up of chemical elements in the early Galaxy (e.g., Tsujimoto et al., 1999; Argast et al., 2000; Oey, 2000; Travaglio et al., 2001; Fields et al., 2002); Tsujimoto et al., (1999) provided an explanation for the spread of Eu observed in the oldest halo stars in the context of a model of supernova-induced star formation. Assuming that r-process nucleosynthesis sites would most likely be identified with Type II supernovae (SNe II), Mathews et al. (1992) suggested a mass range of $M_{ms}$=$7-8M_{\\odot}$ for the site, while Travaglio et al. (1999) supported a somewhat higher mass range, $M_{ms}$=$8-10M_{\\odot}$. Cescutti et al. (2005) concluded that the Eu should originate as an r-process element in stars with masses in the range $10-30M_{\\odot}$.  Recent observations and the analyses imply that the abundance pattern of an extremely metal-deficient star with [Fe/H]$\\leq$-2.5 may retain information of a preceding single supernova (SN) event or at most a few SNe (Mcwilliam et al., 1995; Ryan et al., 1996). Tsujimoto and Shigeyama (1998) have also shown that the mass of r-process elements ejected by each SN II as a function of progenitor mass at the main sequence ($M_{ms}$) can be derived from the observed [Ba/Mg]-[Mg/H] trend combined with the [Mg/H]-$M_{ms}$ relation in theoretical SN models. But some of extremely metal-deficient stars with [Fe/H]$\\leq$-2.5 could be formed out of gas enriched by several SNe (Fields et al., 2002).  In this paper, we assume that the stars on the left-side [Ba/Mg]-[Mg/H] boundary are made from individual supernova events. In Section 2 we discuss the r-process elements yields based on the model proposed by Tsujimoto and Shigeyama (1998). In Section 3 the r-process elements yields are used to explore the early stages of inhomogeneous chemical evolution in the Galaxy. Conclusions are given in Section 4. ",
        "conclusions": "\\label{sec:4} Assuming that the stars on the left-side [Ba/Mg]-[Mg/H] boundary retain the abundance pattern of a single supernova, we have calculated the r-process elements yields as a function of the initial stellar mass from theoretical SN models. The scatter of [r/Fe] in halo strongly suggests that SNe II associated with stars of progenitor mass $M_{ms}\\leq18M_{\\odot}$ are infertile sources for the production of r-process elements, which is in agreement with the result of Tsujimoto et al. (2000). We conclude that SNe II with $20M_{\\odot}\\leq M_{ms}\\leq 40M_{\\odot}$ are the dominant source of r-process nucleosynthesis in the Galaxy. The effect on the SNe II which produces the high r-process yield could be significant. Assuming a Salpeter IMF with $\\alpha$=-2.35, the ratio of these stars compared to the all massive stars is about $\\sim$18\\%, which is higher than the value of case A in Fields, Truran, and Cowan (2002). In fact, the mean value of [r/Fe] obtained in our work is also higher than the value of \\={R} in Fields et al. (2002).  We have presented an approach to understand the scatter in heavy r-process-to-iron ratio in metal-poor halo stars and used a stochastic description of progenitor mass SNe II that contribution to the heavy r-process and iron abundances in each halo star. The random star-to-star variations in nucleosynthetic ancestry lead to scatter in [r/Fe]. The models we present are all successful in reproducing the scatter in the available data, which go down to about [Fe/H]$\\approx$-3. In conclusion, the ratios of [r/Fe] with the metallicity are twofold. One is the abundance ratios for [Fe/H]$\\leq$-2.5 imprinted by the nucleosynthesis in a few supernovae on the timescale $\\sim10^{7}$ yr and the other for [Fe/H]$\\geq$-2 results from the mixing of the products from a whole site of the nucleosynthesis, taking place on the timescale longer than $10^{9}$ yr."
    },
    "0608/astro-ph0608332_arXiv.txt": {
        "abstract": "We present a sample of 1777 bright ($9<B<14$) metal-poor candidates selected from the Hamburg/ESO Survey (HES). Despite saturation effects present in the red portion of the HES objective-prism spectra, the data were recoverable and quantitative selection criteria could be applied to select the sample. Analyses of medium-resolution ($\\sim2$\\,{\\AA}) follow-up spectroscopy of the entire sample, obtained with several $2$ to $4$\\,m class telescopes, yielded 145 new metal-poor stars with metallicity $\\mbox{[Fe/H]}<-2.0$, of which 79 have $\\mbox{[Fe/H]}<-2.5$, and 17 have $\\mbox{[Fe/H]}<-3.0$. We also obtained C/Fe estimates for all these stars. From this, we find a frequency of C-enhanced ($\\mbox{[C/Fe]}>1.0$) metal-poor ($\\mbox{[Fe/H]}<-2.0$) giants of $9\\%\\pm2\\%$, which is lower than previously reported. However, the frequency raises to similar ($>20\\%$) and higher values with increasing distance from the Galactic plane. Although the numbers of stars at low metallicity are falling rapidly at the lowest metallicities, there is evidence that the fraction of carbon-enhanced metal-poor stars is increasing rapidly as a function of declining metallicity. For $\\sim60$ objects, high-resolution data have already been obtained; one of these, HE~1327$-$2326, is the new record holder for the most iron-deficient star known. ",
        "introduction": "The first systematic searches for metal-poor stars in the Galactic halo, by means of objective-prism surveys, began with \\citet{Bond1970, Bond1980, Bond1981} and \\citet{Bidelman_MacConnell}.  More recent surveys are the HK survey \\citep{BPSI,BPSII,beers_stromlo_symp} and the Hamburg/ESO survey (HES; \\citealt{hespaperIII, Christlieb:2003}). Over time, deeper surveys have became possible to reach further into the Galactic halo. The limiting magnitude of the HES is $B\\sim 17.5$\\,mag as opposed to the earlier HK survey ($B\\sim15.5$\\,mag).  Particularly the HK survey and HES have both shown that a variety of astrophysically interesting and unusual objects can be found in a large sample of metal-poor stars. Some of these exhibit large overabundances of heavy neutron-capture elements that are synthesized in the s-process (e.g., \\citealt{2001aokisprocess, 2001vaneck}), or in the r-process (e.g., CS~22892$-$052, \\citealt{Snedenetal:1996}, or CS~31082-001 \\citealt{Cayreletal:2001}). In the case of strong r-process enhancement, nucleo-chronometry becomes possible based on the abundance of e.g., Th and U. Bright metal-poor stars are particularly useful for this task since the weak uranium line at 3859\\,{\\AA} is only detectable in very high-quality spectra. The availability of U, together with Th, as the U/Th chronometer drastically reduces the theoretical uncertainties on the derived ages, as compared to the more easily detectable Th/Eu ratio \\citep{wanajo2002}. Hence, a lower limit for the age of the Universe can be derived independently of other measurements such as WMAP \\citep{WMAP}. Many \\mbox{C-enhanced} metal-poor objects (hereafter CEMP; $\\mbox{[C/Fe]}\\footnote{We use the common notation of $\\mbox{[A/B]} = \\log({N_{\\rm A}/N_{\\rm B}})_{\\star} -\\log({N_{\\rm A}/N_{\\rm A}})_{\\odot}$, for elements A and B.}>1.0$) were discovered as they occur with increased frequency amongst metal-poor objects (e.g., \\citealt{1999rossicarbon}). One of those is the very iron-deficient giant HE~0107$-$5240 ($\\mbox{[Fe/H]} =-5.2$)\\footnote{Employing the same Fe NLTe correction as used for HE~1327$-$2326} that was found in a sample of $\\sim2000$ HES metal-poor candidates \\citep{HE0107_Nature}. This star provided critical observational material to the theories of the formation of the first objects in the Universe.  For elemental abundance studies involving isotopes, such as $^{6}$Li in metal-poor turnoff stars, or very weak and/or blended lines, such as uranium in some strongly r-process enhanced objects, it is crucial to be able to obtain high-resolution ($R>40,000$) spectra with a very high signal-to-noise ratio. Bright metal-poor stars are thus ideal targets. With the increased light-collecting power of current 8-10\\,m telescopes it is possible to achieve a $S/N$ of $200$ or more for such objects within reasonable exposure times.   Apart from the faint stars, the HES also contains large numbers of bright ($B<14$) stars. However, those have previously not been investigated. The main technical reason is the partial saturatation of the photographic plates, resulting in partially satured spectra. Despite an increased interest to search for more fainter metal-poor stars in the extended survey volume of the HES, recovering those brighter stars is well worth the effort, because they offer a variety of advantages over fainter stars.  In general, for bright stars, the availability of various data sources, such as the photometric survey 2MASS \\citep{2MASS}, or astrometric catalogs, e.g., US Naval Observatory CCD Astrograph Catalog (UCAC2; \\citealt{UCAC}), or the Southern Proper Motion Catalog (SPM3.1; \\citealt{SPM}), provide useful additional information for the analysis of large samples of objects.  Hence, more detailed analyses, including kinematic studies that make use of the full space motions, of many bright stars can be carried out more quickly and easily.  The HES thus provides the opportunity to quickly identify large numbers of bright metal-poor stars in a systematic fashion. Such objects provide important information in the intermediate magnitude range between the early objective prism surveys and the HES. For example, a second star, HE~1327$-$2326, with an extremely low iron abundance of $\\mbox{[Fe/H]} =-5.4$ \\citep{HE1327_Nature, o_he1327,Aokihe1327} was already recently discovered in the sample of bright HES stars discussed in this paper.  Here we present the sample of bright metal-poor candidates selected from the HES. The candidate selection is described in \\S~\\ref{hes}, while the results of the follow-up observations of the sample are presented in \\S~\\ref{followup}. In \\S~\\ref{ana} we outline our data analysis and the abundance results of the metal-poor stars found in the sample. In \\S~\\ref{concl} we briefly conclude with an outlook for future work concerning the newly-discovered stars. ",
        "conclusions": "\\label{concl}  In this paper we have presented a sample of 1777 bright ($9<B<14$) metal-poor candidates selected from the Hamburg/ESO survey. Compared to the previously explored faint HES metal-poor stars, our sample contains stars whose objective-prism spectra suffer from saturation effects. The selection procedure has been described in detail, together with results from medium-resolution follow-up spectroscopy of the entire sample (Table~\\ref{res_mp}). From the medium-resolution observations, the metallicity $\\mbox{[Fe/H]}$ was determined by means of the \\citet{BeersCaKII} method involving the \\texttt{KP} index and a spectroscopically or photometrically derived $B-V$ color. Based on the 2MASS $J-K$ colors, we also employed the \\citet{rossi_carbon05} calibrations to obtain a third iron abundance estimate, as well as [C/Fe]. Where available, the iron abundances derived from the different methods were weighted and averaged. In total we have identified $16\\%$ of the sample stars to be of metal-poor nature (i.e., 286 objects with $\\mbox{[Fe/H]}<-1.0$). Of those, $174$ have $\\mbox{[Fe/H]}<-2.0$, $98$ have $\\mbox{[Fe/H]}<-2.5$, and $23$ have $\\mbox{[Fe/H]}<-3.0$. Amongst those objects with $\\mbox{[Fe/H]}<-2.0$, we found 13 carbon-rich objects with $\\mbox{[C/Fe]}>1.0$ . Apart from the metal-poor stars, a few examples of other groups were also serendipitously identified (through the SIMBAD data base), such as main-sequence A-type stars, emission-line objects, variable stars and one nova-like star.  Due to rediscoveries of several metal-poor stars from the HK survey and other studies, $\\sim30$ of the 286 metal-poor stars are already known as such in the literature. This leaves $\\sim250$ bright objects newly classified as metal-poor. The effective yields of our sample are of roughly the same order as found in the HK survey. These are significantly lower than reported for the faint HES stars. This difference can largely be understood in terms of the saturation effects and the subsequent, adversely affected, selection. We also investigated the frequency of C-enhanced objects amongst metal-poor stars with distance from the Galactic Plane. Two effects are likely to play a role in the production of C-rich objects in the Galaxy. First, the frequency of CEMP stars is higher for lower-metallicity samples, perhaps indicating that the carbon production becomes more efficient with decreasing [Fe/H]. We find this effect to exist independent of distance from the galactic plane. The second effect is a plateau of the frequency of CEMP objects out to $Z\\sim2$\\,kpc (regardless of the sample metallicity cut-off), complemented with a steep increase for larger distances. This may be due to a different mechanism for the production of carbon and iron in the different populations. For the halo population, at least, the mechanism does not seem to be the same for normal metal-poor and CEMP stars.  Using 4-8m class telescopes snapshot spectroscopy ($R \\sim 20,000$, $S/N > 30$ per pixel at $\\sim4100$\\,{\\AA}) of selected targets with sufficiently low metallicity is well underway. Our general procedure is as follows. If a spectrum indicates any interesting features, e.g., r-/s-process enrichment or a metallicity as low as $\\mbox{[Fe/H]}<-4.0$ we are obtaining higher resolution (R$>40,000$), higher signal-to-noise observations. These data are then used for a detailed abundance analysis. This has already been done for one bright star in particular, HE~1327$-$2326. Subaru/HDS data were obtained and revealed it to be the new record holder for the most iron-poor star known to date, with $\\mbox{[Fe/H]}=-5.4$ (Frebel et al. 2005, 2006; \\citealt{Aokihe1327}). Further papers of this series on bright metal-poor stars from the HES will include the high-resolution spectroscopic analyses of the metal-poor objects indentified in the present study as well as the candidate follow-up spectroscopy of the remaining, yet unexplored, 51 HES fields (Frebel et al. 2006, in preparation)."
    },
    "0608/astro-ph0608618_arXiv.txt": {
        "abstract": "We present a period-luminosity-amplitude analysis of 5899 red giant and binary stars in the Large Magellanic Cloud, using publicly available observations of the MACHO project. For each star, we determined new periods, which were double-checked in order to exclude aliases and false periods. The period-luminosity relations confirm the existence of a short-period, small-amplitude P--L sequence at periods shortward of Seq. A. We point out that the widely accepted sequence of eclipsing binaries between  Seqs. C and D, known as Seq. E, does not exist. The correct position for  Seq. E is at periods a factor of two greater, and the few stars genuinely lying  between Seq. C and D are under-luminous Mira variables, presumably  enshrouded in dust. The true Seq. E overlaps with the sequence of Long Secondary Periods (Seq. D) and their P--L relation is well described by a simple model assuming Roche geometry. The amplitudes of LSPs have properties that are different from both the pulsations and the ellipsoidal variations, but they are more similar to the former than the latter, arguing for pulsation rather than binarity as the origin of the LSP phenomenon. ",
        "introduction": "The multiplicity of red giant period-luminosity (P--L) relations has been a major discovery on the road to interpreting complex light variations of these stars. Following the two seminal papers by \\citet{woo99} and \\citet{woo00}, a picture has emerged that can be summarized as follows: large-amplitude Mira stars pulsate in the fundamental mode, whereas smaller-amplitude semiregulars are often multimode pulsators, in which various overtone modes can be excited \\citep[see also][]{bed98}. Besides the pulsating P--L sequences (Seq. A, B and C, as labeled by \\citet{woo99}), two other sequences were suggested: Seq. E with red giants in eclipsing binaries and Seq. D with stars that have long secondary periods (LSPs). The latter pose a great mystery and the nature of their slow variations is still not understood, with several different mechanisms proposed \\citep{oli03,woo04}.  The basic picture of multiple P--L relations has been confirmed by many independent studies, mostly based on $K$-band magnitudes. It has emerged that the original five sequences have further details, including a break at the tip of the Red Giant Branch (RGB), which is due to the existence of distinct RGB pulsators that are mixed with the more evolved AGB variables \\citep[e.g.][]{ita02,kis03,kis04,ita04a,ita04b,sos04a,fra05}. Almost all authors have accepted the existence of the distinct sequences of red giant binaries (Seq. E) and LSP stars (Seq. D). The only exception was \\citet{sos04b}, who showed that Seqs. E and D seem to merge at a specific luminosity (as measured by the Wesenheit index) and suggested that this may imply the binary origin of LSPs.  Here we report on a combined analysis of MACHO observations of eclipsing binaries and red giants in the Large Magellanic Cloud, which shed new light on these stars and on the LSP phenomenon. ",
        "conclusions": "The combined P--L plot in Fig.\\ \\ref{pl} shows the well-known complex structure of distinct sequences. Besides sequences A, B, C, D and E, we also detect the existence of the faintly visible new short-period P--L sequence (Seq. A$^\\prime$) on the left-hand side boundary of the diagram (labeled as a$_{4}$ and b$_{4}$ in \\citet{sos04a} and $P_4$ in \\citet{kis06}). The eclipsing stars (pluses) seem to merge with Seq. D rather than forming Seq. E as adopted in the literature \\citep{woo99,woo00,kis03,kis04,ita04a,ita04b,nod04,fra05}.  To clarify this issue, we re-checked periods for: (i) stars on Seq. E in Fig. 1 of \\citet{woo00}, for which the identifiers and basic data were kindly provided by Peter Wood; (ii) stars on Seq. E in our Fig.\\ \\ref{pl} (black dots). In both cases, it turned out that for most of the objects, the given periods were half of the true ones, as one might expect from a Fourier analysis of eclipsing binary light curves. We have carefully double-checked all individual light curves on Seq. E and D and corrected the periods (see Fig.\\ \\ref{example}). The final P--L plot is shown in the right panel of Fig.\\ \\ref{pl}.  As a result of the period correction, the sequence between C and D, which is known as Seq. E in the literature, has completely disappeared. A few stars remain in the gap but practically none of them are eclipsing binaries. The majority turn out to be Mira stars that are very red ($J-K>2$ mag) and are presumably carbon-rich Miras that are dimmed by circumstellar dust clouds (bottom panel of Fig.\\ \\ref{example}). We propose to retain the label E for the sequence of ellipsoidal variables in its corrected position (right panel of Fig.\\ \\ref{pl}).  \\begin{figure} \\epsscale{0.9} \\plotone{f2.eps} \\caption[]{\\label{example} Example of phase diagrams before (left panels) and  after (right panel) period correction. The first and the second star from the top show eclipsing binary and ellipsoidal variable light curve, while the third  star has a typical LSP variation. The bottom panel shows the light curve of a Mira star that lies in the gap between Seqs. C and D and which faded $\\sim$2 mags during 8 years of MACHO observations.} \\end{figure}   In the literature, only \\citet{sos04b} plotted the eclipsing binary sequence at the correct (doubled) period. However, they did not discuss this issue  and one can still find more recent studies where Seq. E  was shown at the wrong period. \\citet{sos04b} have, however, shown that Roche geometry gives a good fit to the OGLE ellipsoidal variables. We have also checked this on the MACHO sample. We calculated the theoretical orbital periods of systems at mass ratios of 1, where the components fill their Roche lobes. For this, we used evolutionary models of \\citet{cas03} and applied equations of Section 4 in \\citet{sos04b}. (Note, however, that the definition of $f(q)$ in their Eq. 1 actually gave $f(q)^{-1}$ as the filling factor.) For the calculations we took the evolutionary tracks of the 0.85 ${\\rm M}_{\\odot}$ and 2.5 ${\\rm M}_{\\odot}$ models, since theses masses represent the mass limits of stars that evolve through the RGB and AGB. The $K$ magnitudes of the models were determined from $T_{\\rm eff}$ vs. $(V-K)$ calibrations, combined from \\citet{hou00} and \\citet{kuc06}.  In the right panel of Fig.\\ \\ref{pl} we show these two limits (the shorter period line belongs to the higher mass). For any other mass ratios, the orbital periods shift towards smaller values. This simple approach with Roche geometry describes remarkably well the observed period-luminosity relation of ellipsoidal variables in the MACHO sample, in agreement with the study of \\citet{sos04b}. However, the OGLE sample contains a larger fraction of ellipsoidals and most of them have longer periods than are predicted by the models. Those stars do not entirely fill the Roche lobe \\citep{sos04b}. It is also worth mentioning that the low fraction of ellipsoidal/eclipsing RGB stars in the MACHO data (1.5\\%) was used by \\citet[]{woo04} as an argument against the binary origin of LSPs. However, the OGLE statistics clearly showed that there are at least 10 times more such RGB stars, and so that argument is no longer valid.  Since the lower mass model fits Seq. D quite well, the question arises: how similar are the ellipsoidal and LSP variables? To assess this, we examined the amplitudes of these stars. \\citet{sos04b} mentioned that amplitudes of LSPs are positively correlated with the brightness of the star. Compared to OGLE data, the MACHO observations have the advantage of giving information on the color variations. We examined the amplitudes in the MACHO blue and red bands of the binary and LSP stars, and also included Seq. C stars, which allow a comparison with stars that we know to be pulsating.   The amplitudes were measured by fitting smooth spline functions to the phased light curves. The resulting blue peak-to-peak amplitudes are shown as a function of $K$ magnitude in Fig.\\ \\ref{amlu}, where several features are apparent. Firstly, the amplitudes of ellipsoidal variables (circles in the upper panel) do not show any correlation with luminosity, as expected for variation caused by the geometry of a binary system. Secondly, the amplitudes of LSP variables (crosses in the upper panel) increase with luminosity, with their distribution forming a striking triangular envelope (the correlation coefficient is $r\\sim0.51$). It is also very interesting that the points below the well-defined upper envelope seem to show a flat distribution. Whatever the cause of the LSP phenomenon, there appears to be a maximum possible  amplitude at each luminosity, with an apparently uniform distribution of  amplitudes beneath this maximum value.  \\begin{figure} \\epsscale{0.9} \\plotone{f3.eps} \\caption[]{\\label{amlu} Luminosity-amplitude relation of stars on Seq. C, D and E. {\\it Top panel:} Black crosses show the luminosity-amplitude relation for Seq. D stars, while turquoise circles the same for Seq. E stars, which contain  ellipsoidal variables. The difference of the two sequence is very clear, e.g. amplitudes of Seq. D stars show moderately positive correlation with the  luminosity. {\\it Bottom panel:}  Luminosity-amplitude relation for Seq. C stars, which are fundamental mode pulsators. This relation is significantly different from those of  Seqs. D and E stars.} \\end{figure}   In comparison, stars on the pulsating P--L sequence C (lower panel in Fig.\\ \\ref{amlu}), show a very different distribution. To quantify the difference between the amplitude-luminosity distribution of Seqs. C and D, we performed the multivariate $\\varepsilon$-test of \\citet{sze04}, applied for two dimensions. In this test we measure the effects of random permutations of the initial distributions via changes in the ``information energy'' of the two distributions. In our case, 1000 random permutations showed that the difference between the C and D samples is highly significant. Therefore, we conclude that the physical mechanism causing the LSP phenomenon must be different from both the ellipsoidal variations of Seq. E and the radial fundamental-mode pulsations of Seq. C. Note, however, that for Seqs. A and B, there is a similar positive correlation between amplitude and luminosity (see the upper three panels in figs. 4 in \\citet{kis03} and \\citet{kis04}) than for Seq. D.  \\begin{figure} \\epsscale{0.9} \\plotone{f4.eps} \\caption[]{\\label{pamp} Blue-to-red amplitude ratios as function of period for stars on Seqs. B (grey triangles), C (orange pluses), D (black crosses) and E (purple circles). Grey  triangles show those short period pulsations of Seq. D stars, which fall to  Seq. B. The two clearly pulsating sequences, Seqs. B and C are the continuation  of each other. Ellipsoidal variables (Seq. E) have lower amplitude ratio than  pulsating stars. Amplitude ratios of Seq. D stars are more similar to those of pulsating stars.} \\end{figure}   At the same time, comparing blue and red amplitudes for individual stars revealed further interesting information. In Fig.\\ \\ref{pamp} we plot the blue-to-red amplitude ratios as function of period for Seqs. B, C, D and E stars. For the pulsating objects the median of the ratio is 1.40, indicating strong color, thus temperature, changes during the pulsations. For ellipsoidal/eclipsing binaries, the median ratio is 1.13, while the LSPs have a median  ratio of 1.29, being more similar to the pulsating stars \\citep[see also][]{hub03}. This behavior agrees with the findings of \\citet{woo04} who, based on color-amplitude variations of single objects, argued for the pulsational origin of LSPs. The overall statistics of more than 700 LSP variables favors this argument over the binary hypothesis, also agreeing with \\citet{hin02}, who concluded that the long-period velocity changes in their observed stars probably result from some kind of pulsation."
    },
    "0608/astro-ph0608348_arXiv.txt": {
        "abstract": "We establish that global solar $p$-mode frequencies can be measured with sufficient  precision on time scales as short as nine days to detect activity-related shifts. Using ten years of GONG data, we report that mode-mass and error-weighted frequency shifts derived from nine days are significantly correlated with the strength of solar activity and are consistent with long-duration measurements from GONG and the SOHO/MDI instrument. The analysis of the year-wise distribution of the frequency shifts with change in  activity indices shows that both the linear-regression slopes and the magnitude of the correlation varies from year to year and they are well correlated with each other.  The study also indicates that the magnetic indices behave differently in the rising and falling phases of the activity cycle. For the short-duration nine-day observations, we report  a higher sensitivity to activity. ",
        "introduction": "\\label{Introduction} It is now well established that the global oscillation frequencies of the Sun change in phase with the solar activity cycle, see for example Jain and Bhatnagar (2003)  for intermediate-degree modes, Jim{\\'e}nez-Reyes \\emph{et al.} (2004) and Howe \\emph{et al.} (2006) for low-degree modes, and references therein.  However, there is still no consensus regarding the physical mechanism that gives rise to these changes (Kuhn, 2001).  One deficiency is the lack of measurements for modes at high frequency and high degree (Rhodes, Reiter, and Schou, 2003) and little information about the correlation on short time scales.  Since the {\\it p}-mode frequency changes are thought to be associated with individual active regions that come and go continuously (Hindman \\etal, 2000), one would anticipate that the frequencies also change continuously on any time scale. Due to the  finite lifetime of the modes, the correlation between frequency and activity may depend on the length of the observing run.  Thus, the use of mode frequencies and other parameters derived from short time series, during which the solar activity varies less, may help understand the underlying mechanism of these variations.  Rhodes, Reiter, and Schou (2003) used  a few sets of three-day time series in the computation and fitting of  intermediate-degree {\\it p}-mode frequencies and widths.  The analysis, using Michelson Doppler Imager (MDI) and Mount Wilson Observatory (MWO) data, resulted in a higher sensitivity as measured from the slope of the linear regression between the frequency shifts and activity differences, compared to the slopes from longer time series.  However, the result was not confirmed when Global Oscillation Network Group (GONG) data were used (Rose \\etal, 2003).   In the domain of low-degree modes, Chaplin \\etal{} (2001) studied global modes on different time scales, the shortest of which was 27 days. Their results suggest that the sensitivity to changes in the Kitt Peak magnetic field measurements may be higher in the rising phase of the cycle. Salabert \\etal\\ (2002) studied high-frequency modes ($\\nu >$ 3.7 mHz) with a time series of eight days using IRIS$^{++}$ data and confirmed that the frequency shift becomes negative above 4.5~mHz up to the cutoff frequency of 5.5~mHz. A similar trend for intermediate-degree frequencies was reported earlier by Ronan, Cadora, and LaBonte (1994) and Jefferies (1998).   Here, we study short temporal variations in the frequency shifts and their degree of correlation with solar activity over a complete solar cycle. To do so, we use  GONG data for the period 1995\\todash{}2005, which covers partially the descending phase of the cycle 22 and the ascending and descending part of the current solar cycle 23. The analysis is carried out by calculating {\\it p}-mode frequencies on a time scale as short as nine days. ",
        "conclusions": "Figure~\\ref{fig3} shows the mean frequency shifts as a function of the epoch (left panel) and the change in activity (right panel) for two activity indices; one magnetic (MPSI) and the other non-magnetic ($F_{10}$). It is reassuring to see that the mean frequency shifts derived from nine-days clearly show a similar solar cycle variation  to those derived from longer time series. The only difference is the rapid fluctuations seen in each of the plots which arises due to the short interval over which the mean frequency shifts and mean activity differences are calculated. The solid lines in the two right-hand panels represent the straight lines obtained by using Equation~(3) to give the best linear fits to the two sets of data points; the corresponding slopes and intercepts of both fits are tabulated in the last rows of Tables~1 and 2.  A significant correlation between the frequency shifts  and change in activity indices is evident except near the activity-minimum phase of the cycle around 1996\\todash{}97 reflecting deviations from the assumed linear dependence.  We examine this more critically in the next Section.  \\subsection{Year-wise Variations} In order to study how closely the frequencies follow the activity levels, we study year-wise variations. The distribution of the number of points in each year can be seen in Table~1. As before, the mean frequency shift for each year (also for the entire data set) is calculated using Equation~(2) but the reference frequency is now considered to be the middle point in each set.  Similarly, the middle point in each activity set was considered as the reference activity for the linear regression analysis. For example, for the year 1995, the reference point for calculating the frequency shift and change in activity was taken as the 14th point.  The results of the linear least-square fits and correlation coefficients for $F_{10}$ are summarized in Table~1. For comparison, we also tabulate the slope and correlation coefficient obtained from the entire data set in the last row of the Table~1. We find significant year-to-year variation in the calculated slopes and the correlation coefficients.  For a more detailed investigation of the variation in the slopes, we show mean-frequency shifts as a function of the activity differences ($\\delta F_{10}$)  for four  representative  years in Figure~\\ref{fig4}.  The solid lines in the four panels represent the linear fits to the pairs of frequency differences and 10.7-cm flux differences in the four different years, while the dashed line which is repeated in all four panels, is the linear fit to all 424 datasets. This linear fit can be considered  as the representative fit for the current solar cycle. It is  interesting to note that the solid and dashed lines for the year 1997 have similar orientation implying close agreement in the values of the slopes indicating similar mean shifts per unit activity. Surprisingly, the best correlation between the frequency shifts  and 10.7 cm radio flux also exists during 1997 which represents the beginning of the rising phase of the current solar cycle.  Results for MPSI obtained with the similar analysis  are given in Table~2 and Figure~\\ref{fig5}.  For MPSI, a significant correlation is obtained for  1997, 1998, and 2001 and a poor correlation for  1996, 2003, and 2005, and it closely agrees with $F_{10}$.  Similar results are found for all other activity indices considered in this paper.  Specifically, we note that the correlation coefficients and slopes obtained from the analysis of the entire data sets are similar to those of 1997 or slightly higher and do not represent average values of the year-wise variations.  This can be more clearly visualized in Figure~\\ref{fig6} where we have shown the linear fits between the frequency shifts and activity differences using the slopes and intercepts  given in Tables~1 and 2. Numerically, the average slopes of the yearly regression fits are 1.26 nHz/sfu for $F_{10}$ and 52.57 nHz/Gauss for MPSI. These mean slopes are each approximately a  factor of two smaller than are the slopes of the overall fits to the corresponding collection of 424 datasets.  \\begin{figure}             % \\begin{center} \\psfig{file=fig4.ps,angle=90,width=\\linewidth} \\caption{Mean-frequency shifts plotted as a function of  $\\delta F_{10}$ for four selected years as marked on each figure.  The solid line denotes the linear fit to the activity measurements for that year. The error bars signify the 1$\\sigma$ uncertainties in the mean-frequencies.  The dashed line denotes the linear fit to all 424 data points.} \\label{fig4} \\end{center} \\end{figure} \\begin{figure}              % \\begin{center} \\psfig{file=fig5.ps,angle=90,width=\\linewidth}            % \\caption{Same as Figure~\\ref{fig4} but for $\\delta$MPSI}. \\label{fig5} \\end{center} \\end{figure}  In a similar analysis, Woodard \\etal\\ (1991) have studied the variation of oscillation frequencies during 1986, 1988, and 1989. They report different correlation coefficients for annual variations (except for the magnetic field during 1988 and 1989) which are also different when all the  data were taken together. Our result qualitatively agrees with these findings although there are differences in detail. However, the discrepancy can not be investigated further since the two studies do not have common activity indices.  Since the linear regression analysis clearly shows the  yearly variations in the slopes and the correlation coefficients, we investigate a possible correlation between these two quantities. This is carried out by plotting the annual variation of the correlation coefficients along with the slopes for four out of the five activity indicators considered in this paper (Figure ~\\ref{fig7}). The similarity between the two curves stands out in each case signifying strong correlations, although some differences can be seen in detail. The correlation is further confirmed by calculating the linear and rank correlation coefficients between the two quantities (Table~3). The maximum correlation is obtained for \\ion{Mg}{II}, followed by MPSI while MWSI has the lowest correlation coefficient.  The relationship between the annual linear-regression slopes and yearly averaged solar activity is shown in Figure~\\ref{fig8}.  The striking feature is the strong temporal variability of the regression slopes which do not appear to be correlated with activity indices on the time scale of the solar cycle.  On closer examination, there appears to be a good correlation between 1995--1998 and 2001--2005.  This could indicate that there are two components of the activity which contribute to the variability of the slopes.  The major contribution comes from the overall weak magnetic field, and hence during the ascending and descending phase  of the cycle we have a strong correlation with activity.  But during the maximum phase of the cycle, when the field is dominated by the strong field, due to the increase in the number and strength of the active regions, the slopes do not follow the activity measurements.  \\begin{landscape} \\begin{center}       % \\begin{table} \\caption{Year-wise distribution of correlation statistics between  frequency shifts (\\,$\\delta\\nu$\\,) and $\\delta F_{10}$. Shown are the number of data sets (\\,$N$\\,), epoch covered, the mean shift per unit change in activity (\\,$a$\\,) in nHz/sfu, the intercept (\\,$b$\\,) in nHz, Pearson's linear coefficient (\\,$P_{\\rm p}$\\,) Spearman rank correlation coefficient (\\,$r_{\\rm s}$\\,) and its two-sided significance (\\,$P_{\\rm s}$\\,). The last row denotes the values corresponding to all 424 data sets. } \\begin{tabular}{rr@{ \\todash {}}rr@{ $ \\pm $ }rr@{ $ \\pm $ }rccl} \\hline $N$&\\multicolumn{2}{c}{Epoch}&\\multicolumn{2}{c}{$a$}&\\multicolumn{2}{c}{$b$}&$P_{\\rm p}$&$r_{\\rm s}$&$P_{\\rm s}$\\\\ \\hline\\noalign{\\smallskip} 27&7 May 1995 & 4 January 1996&1.25 & 0.56&-6.33&2.70&0.40&0.40&3.9 $\\times\\ 10^{-2}$\\\\ 41&5 January 1996&  7 January 1997&1.56 & 0.53&3.42&3.45&0.42&0.42&5.9 $\\times\\ 10^{-3}$\\\\ 40&8 January 1997&  2 January 1998&2.73 & 0.19&10.62&2.68&0.92&0.92&6.1 $\\times\\ 10^{-17}$\\\\ 41&3 January 1998&  6 January 1999&2.14 & 0.23&5.88&5.06&0.83&0.79&1.1 $\\times\\ 10^{-9}$\\\\ 40&7 January 1999&  1 January 2000&1.02 & 0.20&-0.83&10.40&0.64&0.64&9.4 $\\times\\ 10^{-6}$\\\\ 41&2 January 2000&  4 January 2001&0.69 & 0.15&-43.47&4.17&0.58&0.50&9.8 $\\times\\ 10^{-4}$\\\\ 41&5 January 2001&  8 January 2002&1.59 & 0.15&-4.74&9.30&0.87&0.88&6.4 $\\times\\ 10^{-14}$\\\\ 40&9 January 2002&  3 January 2003&1.22 & 0.18&-4.32&7.09&0.75&0.76&1.5 $\\times\\ 10^{-8}$\\\\ 41&4 January 2003&  7 January 2004&0.46 & 0.17&-11.38&4.21&0.39&0.37&1.9 $\\times\\ 10^{-2}$\\\\ 40&8 January 2004&  1 January 2005&0.70 & 0.15&-3.75&3.80&0.60&0.66&3.4 $\\times\\ 10^{-6}$\\\\ 32&2 January 2005&  16 October 2005&0.48 & 0.18&1.13&2.79&0.43&0.33&6.5 $\\times\\ 10^{-2}$\\\\ 424&7  May 1995&  16 October 2005&2.51 & 0.04&32.98&5.54&0.94&0.95&0.0\\\\ \\noalign{\\smallskip}\\hline \\end{tabular} \\end{table} \\end{center} \\end{landscape} \\begin{landscape} \\begin{center}      % \\begin{table} \\caption{Year-wise distribution of correlation statistics between frequency shifts (\\,$\\delta\\nu$\\,) and $\\delta{\\rm{MPSI}}$. The notations have the same meaning as in Table~1. The units of $a$ are nHz/Gauss.} \\begin{tabular}{rr@{ \\todash {}}rr@{ $ \\pm $ }rr@{ $ \\pm $ }rccl} \\hline $N$&\\multicolumn{2}{c}{Epoch}&\\multicolumn{2}{c}{$a$}&\\multicolumn{2}{c}{$b$}&$P_{\\rm p}$&$r_{\\rm s}$&$P_{\\rm s}$\\\\ \\hline\\noalign{\\smallskip} 27&7 May 1995 & 4 January 1996 &51.45 & 24.30&-4.65&2.53&0.39&0.36&6.2 $\\times\\ 10^{-2}$\\\\ 41&5 January 1996 & 7 January 1997&34.71 & 20.30&1.49&4.15&0.26&0.24&1.3 $\\times\\ 10^{-1}$\\\\ 40&8 January 1997 & 2 January 1998&113.35 & 11.31&11.79&3.64&0.85&0.87&2.1 $\\times\\ 10^{-13}$\\\\ 41&3 January 1998 & 6 January 1999&104.59 & 11.56&40.20&7.09&0.82&0.80&3.4 $\\times\\ 10^{-10}$\\\\ 40&7 January 1999 & 1 January 2000&55.57& 11.51&0.06&11.04&0.61&0.54&3.5 $\\times\\ 10^{-4}$\\\\ 41&2 January 2000 & 4 January 2001&28.10 &  10.05&-48.39&5.16&0.41&0.32&4.0 $\\times\\ 10^{-2}$\\\\ 41&5 January 2001 & 8 January 2002&84.78 &  8.04&2.79&8.96&0.86&0.86&5.3 $\\times\\ 10^{-13}$\\\\ 40&9 January 2002 & 3 January 2003&61.10&  6.88&-5.32&5.86&0.82&0.78&2.3 $\\times\\ 10^{-9}$\\\\ 41&4 January 2003 & 7 January 2004&14.63 &  9.39&-11.23&4.77&0.24&0.27&8.9 $\\times\\ 10^{-2}$\\\\ 40&8 January 2004 & 1 January 2005&24.29 &  8.52&-0.06&4.47&0.42&0.36&2.4 $\\times\\ 10^{-2}$\\\\ 32&2 January 2005 & 16 October 2005&5.68 & 8.46&3.65&3.10&0.12&0.11&5.5 $\\times\\ 10^{-1}$\\\\ 424&7  May 1995 & 16 October 2005 &118.25 & 2.10&7.70&5.03&0.94&0.94&0.0\\\\ \\noalign{\\smallskip}\\hline \\end{tabular} \\end{table} \\end{center} \\end{landscape}    \\begin{figure}          % \\begin{center} \\vskip -1.5in \\psfig{file=fig6_new.ps,angle=90,width=\\linewidth} \\caption{The panels shows the linear fit between the year-wise frequency shifts and activity differences and are constructed using the slopes and intercept values  given in Table~1 (left panel) and Table~2 (right panel). } \\label{fig6} \\end{center} \\end{figure}     \\begin{figure}[h]           % \\begin{center} \\psfig{file=fig7.ps,angle=90,width=\\linewidth} \\caption{Year-wise variation of the linear-regression slope (diamonds) and the Pearson's correlation coefficients (pluses) for different activity indices. A strong correlation between the two is clearly visible. The error bars signify the $1\\sigma$ errors in the determination of the slopes.}\\label{fig7} \\end{center} \\end{figure}     \\begin{center} % \\begin{table} \\caption{The correlation statistics between slopes and correlation coefficients for different activity indicators. Shown are the Pearson's linear coefficient (\\,$P_{\\rm p}$\\,), the rank correlation (\\,$r_{\\rm s}$\\,), and the two-sided significance (\\,$P_{\\rm s}$\\,).} \\begin{tabular*}{0.60\\textwidth}% {@{\\extracolsep{\\fill}}rccr} \\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{Activity}&$P_{\\rm p}$&$r_{\\rm s}$&\\multicolumn{1}{c}{$P_{\\rm s}$}\\\\ \\hline $F_{10}$ &0.72&0.72&1.1 $\\times$ $10^{-2}$\\\\ MPSI&0.90&0.89&2.3 $\\times$ $10^{-4}$\\\\ MWSI&0.57&0.58&6.0 $\\times$ $10^{-2}$\\\\ Mg\\,{\\sc ii}&0.96&0.91&1.1 $\\times$ $10^{-4}$\\\\ $R_{\\rm I}$&0.81&0.81&2.6 $\\times$ $10^{-3}$\\\\ \\noalign{\\smallskip}\\hline \\end{tabular*} \\end{table} \\end{center}  \\begin{figure}           % \\begin{center} \\psfig{file=fig8.ps,angle=90,width=\\linewidth} \\caption{Temporal dependence of the linear-regression slopes (diamonds) and corresponding  yearly-averaged solar activity (stars). The error bars signify the 1$\\sigma$ errors in the determination of the slopes.} \\label{fig8} \\end{center} \\end{figure}    We finally note that the year-wise variation of the oscillation frequencies is distinctly different than the analysis involving the entire data sets emphasizing that both the long and short--term variations are different in nature. We also find that the long-term variation is not a simple average of the short-term variations. Therefore, we conclude that studies of changes of oscillation frequencies  which consider all the data sets together may not accurately represent the variations on a scale which is smaller than the length of the entire data sets.  \\subsection{Variation in Rising and Falling Phases of Activity Cycle 23} From the discussion given in Section~4.1, it is evident that there is no significant difference in the correlation pattern between the magnetic (Table~2) and non-magnetic (Table~1) activity indices.  But the analysis  of Chaplin \\etal\\ (2001) involving low-degree modes suggested that the sensitivity to changes in the disk-averaged line-of-sight magnetic field component may be higher during the rising phase of the cycle.  We perform a similar analysis by separating our data set into rising and falling phases of activity cycle 23. Since May 1996 and April 2000 mark the minimum and maximum of cycle 23 (see SGD), we choose the period 1 May 1996 \\todash{}15 May 2000  consisting of 1476 days  as the rising phase.  Avoiding the secondary maximum of this cycle,  we choose 02 October 2001 \\todash{}16 October 2005  again consisting of 1476 days as the falling phase so that both halves have a similar minimum to maximum range in activity indices. Further, to be consistent with the analysis of annual variations, $\\delta I$ is calculated by taking the middle data point in each set as the reference. The results are summarized in Table~4 where we have shown the slopes, the rank correlation coefficients, and the differences in the slopes  between descending and ascending phases normalized by their  combined uncertainty ($\\sigma_{\\delta a}$); a negative $\\sigma_{\\delta a}$  indicating a higher value for the rising phase.  The differences for magnetic indices, MPSI and MWSI, with $\\sigma_{\\delta a}~>$~4 suggest significant differences in behaviour of the magnetic flux emerging over the ascending and descending phases.  The value of 2.79 for the 10.7 cm radio flux implies behaviour similar to the magnetic indices but not as significant as the magnetic flux, while {Mg}{\\sc II} with a value closer to zero implies no phase dependence.  The sunspot number, the only index with a positive $\\sigma_{\\delta a}$  exhibits an opposite behaviour compared to other indices, but since its value is less than 2, the difference can not be considered to be of significance.  This result broadly agrees with the findings of Chaplin \\etal{} (2001) where a phase-dependent behaviour   for the Kitt Peak magnetic index was reported; but it is not clear why other magnetic indices {\\it e.g.} MPSI and MWSI did not show a similar difference. Our results, on the other hand, shows phase-dependent behaviour for all of the magnetic indices and  suggest that the frequency shifts are caused by the magnetic fields at or very near the solar surface.  Similarly, the rank correlations obtained between the frequency shifts and activity indicators also manifest a phase-dependent pattern. In general, the rank correlation coefficients are lower in the descending phase of the activity cycle. The maximum change is seen in the case of MWSI but we should point out that there are significant gaps in the measurement of MWSI index which may contribute to this discrepancy.  Finally, we speculate that this phase-dependent behaviour of the frequency shifts with magnetic flux gives rise to the  hysteresis pattern seen only in case of the magnetic activity indices (Jim{\\'e}nez-Reyes \\etal, 1998 and  Tripathy \\etal, 2001).  \\begin{landscape} \\begin{center}     % \\begin{table} \\caption{Results of linear fits to the change in activity indices in the rising and falling phases of cycle 23. Each phase consists of about four years of data. Shown are the mean shift per unit change in activity (\\,$a$\\,), the intercept (\\,$b$\\,), Spearman rank correlation coefficient (\\,$r_{\\rm s}$\\,).  The last column (\\,$\\sigma_{\\delta a}$\\,) gives the difference in the slopes between ascending and descending phases normalized by the combined uncertainty. We have omitted the two-sided significance as these are all zeros indicating high significance except for MWSI  with a value of $10^{-22}$ in the descending phase.} \\begin{tabular}{lr@{ $\\pm$ }rr@{ $\\pm$ }rcr@{ $\\pm$ }rr@{ $\\pm$ }rcr} \\hline\\noalign{\\smallskip} &\\multicolumn{5}{c}{Ascending phase}&\\multicolumn{5}{c}{Descending phase}&\\\\ {Activity}&\\multicolumn{2}{c}{$a$}&\\multicolumn{2}{c}{$b$}&$r_{\\rm s}$&\\multicolumn{2}{c}{$a$}&\\multicolumn{2}{c}{$b$} &$r_{\\rm s}$&$\\sigma_{\\delta a}$\\\\ index&\\multicolumn{2}{c}{(nHz/index)}&\\multicolumn{2}{c}{(nHz)}&&\\multicolumn{2}{c}{(nHz/index)}&\\multicolumn{2}{c}{(nHz)}&\\\\ \\hline\\noalign{\\smallskip} $F_{10}$&2.54 & 0.08&180.07 & 3.12&0.95&2.32 & 0.08&174.44 & 3.66&0.90&$-$2.79\\\\ {MPSI}&129.72 & 3.67&142.63 & 3.01&0.94&109.08 & 3.54&202.59 & 3.28&0.88&$-$5.72\\\\ {MWSI}&358.29 & 26.12&167.21 & 5.93&0.86&248.41& 24.48&184.73 & 7.23&0.67&$-$4.34\\\\ {Mg\\,{\\sc ii}}&21476~~~~&516~~~~ &162.86 & 2.55&0.95&21326~~~~&638~~~ &199.47 & 3.07&0.90&$-$0.26\\\\ $R_I$&2.30 & 0.08&190.29 & 3.82&0.92&2.49 & 0.12&184.51 & 4.64&0.84&1.84 \\\\ \\noalign{\\smallskip}\\hline \\end{tabular} \\end{table} \\end{center} \\end{landscape}  \\begin{center}     % \\begin{table}[ht] \\caption{Results of linear fits to differences in activity indices and correlation statistics for  frequency shifts obtained from time series of different length. Shown are the length of the time (\\,$t$\\,), the mean shift per unit change in activity (\\,$a$\\,), the intercept (\\,$b$\\,), Pearson's linear coefficient (\\,$P_{\\rm p}$\\,), Spearman rank correlation (\\,$r_{\\rm s}$\\,). The two-sided significance (\\,$P_{\\rm s}$\\,) is smaller than $10^{-17}$ and is not shown.} \\begin{tabular*}{0.80\\textwidth}% {@{\\extracolsep{\\fill}}crrrcc} \\hline\\noalign{\\smallskip} Activity&\\multicolumn{1}{c}{$t$}&\\multicolumn{1}{c}{$a$} &\\multicolumn{1}{c}{$b$}&$P_{\\rm p}$&$r_{\\rm s}$\\\\ index&\\multicolumn{1}{c}{(days)}&\\multicolumn{1}{c}{(nHz/index)}&\\multicolumn{1}{c}{(nHz)}&&\\\\ \\hline $F_{10}$&9 &2.44 $\\pm$ 0.04&55.72 $\\pm$ 3.24&0.94&0.95\\\\ &72&1.75 $\\pm$ 0.04&15.37 $\\pm$ 2.69&0.99&0.99\\\\ &108&1.77 $\\pm$ 0.02&14.57 $\\pm$ 1.82&0.99&0.99\\\\ \\hline MPSI&9&114.90 $\\pm$ 2.05&47.41 $\\pm$ 3.27& 0.94&0.94\\\\ &72&79.45 $\\pm$ 2.23&15.49 $\\pm$ 3.57&0.98&0.98\\\\ &108&80.41 $\\pm$ 1.34&11.93 $\\pm$ 2.14&0.99&0.98\\\\ \\hline MWSI&9&427.6 $\\pm$ 22.7&123.93 $\\pm$ 5.63 &0.68&0.80\\\\ &72&481.7 $\\pm$ 40.5&34.94 $\\pm$ 8.96&0.87&0.89\\\\ &108&504.6 $\\pm$ 25.9&29.57 $\\pm$ 5.65&0.89&0.91\\\\ \\hline {Mg\\,{\\sc ii}}&9&19621~ $\\pm$ 344~~&23.38 $\\pm$ 3.56&0.94&0.94\\\\ &72&13601~ $\\pm$ 494~~&0.26 $\\pm$ 5.07&0.97&0.97\\\\ &108&13330~ $\\pm$ 312~~&1.47 $\\pm$ 3.20&0.97&0.97\\\\ \\hline $R_{\\rm I}$&9&2.40 $\\pm$ 0.06&65.77 $\\pm$ 3.96&0.90&0.92\\\\ &72&1.81 $\\pm$ 0.07&12.80 $\\pm$ 5.21&0.97&0.95\\\\ &108&1.85 $\\pm$ 0.04&8.72 $\\pm$ 2.97&0.97&0.97\\\\ \\noalign{\\smallskip}\\hline \\end{tabular*} \\end{table} \\end{center}  \\subsection{Comparison With Standard GONG and MDI Data} It is expected that with longer time series, the determination of the frequencies will be more precise as the uncertainty is inversely proportional to the square root of the length of the time series for resolved modes -- those with lifetimes less than the length of the time series (Libbrecht, 1992).  Additionally, small fluctuations in activity would be averaged out during the longer time period. Thus, it appears that the correlation between the frequency shift and activity will depend on the length of the observing run.  The analysis involving low-degree modes (Chaplin \\etal, 2001) on different time scales, from 216 days to 27 days, confirms that variations in rank correlation are not inconsistent with the change in accuracy expected for different observing intervals. They also report a small rank-correlation coefficient of the order of 0.67 on the time scale of 27 days which is independent of all of the activity indices.  If we extrapolate this result to a temporal interval of nine days, we expect to find  a very small correlation between the activity indices and frequency shifts.  We compare our results obtained from three different time samples, nine days and 108 days from GONG, and 72 days from MDI in Table~5. The 108-day GONG data covering the period between 7 May 1995 and 16 October 2005 consists of 104 sets with an overlapping period of 36 days between the two adjacent data sets. The MDI data consists of 46 intervals spanning the period from 1 May 1996 to  16 October 2005. The reference frequency for each data set is chosen in such a way that it corresponds to a minimum value of the activity. From Table~5, we find that all three correlations are highly significant with slight variations in the degree of correlation probably due to the different length of the time series.  For the nine-day frequencies, we find the smallest correlation coefficients but they are significantly higher than the values corresponding to analysis of low-degree modes on time scale of 27 days (Chaplin \\etal, 2001).  Since localized  active regions have much more power at high degrees,  the low and intermediate-degree modes have different responses to activity indices. We also note that for a data set of given length, no significant difference  is found between correlations among different activity indices except in the case of MWSI. This is probably due to the large number of gaps in the data which resulted in a poor average when the average was taken over nine days and had no significant effect for averages taken over longer periods.  Similar evidence for different behaviour between low and intermediate-degree modes comes from the comparison of the sensitivity values. We find a higher sensitivity for the short-duration frequencies and confirm the findings of Rhodes, Reiter, and Schou (2003). However, we do not find any explanation for the results of Rose \\etal{} (2003) who could not confirm the higher value of sensitivity using GONG$^{++}$ data. It is possible that during the high-activity period of 2001, the variation in sensitivity is too small to be detected. This result is again different from that of Chaplin \\etal{} (2001), who found similar values for all time intervals considered in their study. Finally, it is worth noting that only in the case of the MWSI index, the slope from the nine-day time series is smaller than the slopes from either the 72- or 108-day time series and may again be attributed to the problem of data gaps as discussed earlier.  \\begin{figure}           % \\begin{center} \\psfig{file=fig9a.ps,angle=90,width=12.cm,clip=} \\psfig{file=fig9b.ps,angle=90,width=12.cm,clip=} \\caption{The Spearman rank correlation coefficients, as a function of frequency, between the weighted frequency shifts and $\\delta{\\rm{MPSI}}$ for all modes present in nine-day (top panel) and 108-day frequency determinations (bottom panel).} \\label{indicor} \\end{center} \\end{figure} \\begin{figure}  % \\begin{center} \\psfig{file=fig10a_bw.ps,angle=90,width=12.cm,clip=} \\psfig{file=fig10b_bw.ps,angle=90,width=12.cm,clip=} \\caption{Correlation maps between the frequency shifts of individual modes and $\\delta{\\rm{MPSI}}$ as a function of  degree (\\,$\\ell$\\,) and frequency (\\,$\\nu$\\,). The coefficients obtained from Spearman's rank correlation analysis are binned with a bin size of eight in $\\ell$ and $300~\\mu$Hz in $\\nu$. The top and bottom panels correspond to frequencies measured from nine days and 108 days duration, respectively. } \\label{cormap} \\end{center} \\end{figure} \\subsection{Correlation of Individual Modes} In order to investigate the differences that exist between the low and intermediate-- degree modes and between  different lengths of the observing run, we carry out the same fitting and correlation analysis for each individual mode using only the GONG data. The Spearman rank correlation coefficients for $\\delta{\\rm{MPSI}}$ are shown in Figure~\\ref{indicor}.  It is evident that the correlation is a function of frequency and shows a three-part structure: a rise in the low-frequency range, a plateau in the five-minute band, and a decrease in the high-frequency zone.  In the case of nine days (top panel of Figure~\\ref{indicor}), the extent of the plateau region is small indicating that a high value of correlation exists only for a small mode set.  On  average, it appears that the nine-day correlation coefficients are small  compared to the 108-day coefficients confirming that the correlation is a function of the length of the observation. We also note negative coefficients indicating anti-correlation at low ($\\nu <$ 1700 $\\mu$Hz) and high ($\\nu >$ 4200 $\\mu$Hz) ends of the frequency range. Similar anti-correlation has been reported by  Ronan, Cadora, and LaBonte (1994), Harvey (1995), Jefferies (1998), and Rhodes, Reiter, and Schou (2003).  However, the frequency range where the shift becomes negative is different in our study than those reported earlier. One possible explanation of the discrepancy could be the lack of sufficient modes at higher frequencies in this analysis. Further work is needed to understand the source of the difference and is beyond the scope of the present paper.  To better illustrate the dependence of the correlation on mode characteristics, we make rank correlation maps as a function of  $\\ell$ and $\\nu$ for all modes in nine-day and 108-day data sets with a bin size of eight in $\\ell$ and 300~$\\mu$Hz in $\\nu$ (Figure~\\ref{cormap}).  In addition to the frequency dependence of the correlation seen in Figure~\\ref{indicor}, these maps show a decreasing correlation with decreasing $\\ell$ for the nine-day sets.  This decrease is markedly smaller for the 108-day analysis. The effect may arise from the increase of the mode lifetime with decreasing $\\ell$. Longer-lived modes will be observed less often in short time series due to the greater temporal interval between excitation events. The decrease of the signal-to-noise ratio at low frequencies and low degrees may also play a role. Since the value of correlation decreases at low-degree and low frequency, it seems possible that, for low-degree modes ($\\ell$ $\\le$ 3), it may drop to the values observed by Chaplin \\etal\\ (2001). Figure~\\ref{cormap} emphasizes that the correlation between frequency shifts and activity indices are complex and have different behaviours for different mode sets and observing periods."
    },
    "0608/astro-ph0608381_arXiv.txt": {
        "abstract": "We visually and quantitatively determine the host galaxy morphologies of 94 intermediate redshift ($0.2 \\le z < 1.2$) active galactic nuclei (AGN), selected using {\\it Chandra} X-ray and {\\it Spitzer} mid-infrared data in the Extended Groth Strip. Using recently developed morphology measures, the second-order moment of the brightest $20\\%$ of a galaxy's flux ($M_{20}$) and the Gini coefficient, we find that X-ray--selected AGN mostly reside in E/S0/Sa galaxies ($53^{+11}_{-10}\\%$), while IR--selected AGN show no clear preference for host morphology. X-ray--selected AGN hosts are members of close pairs more often than the field population by a factor of $3.3\\pm1.4$, but most of these pair members appear to be undisturbed early-type galaxies and do not tend to show direct evidence of gravitational perturbations or interactions. Thus, the activation mechanism for AGN activity remains unknown, even for pair members. ",
        "introduction": "Galaxy interactions and major mergers are currently the most popular explanation for energetically active galactic centers (e.g., Jogee 2005), commonly termed active galactic nuclei (AGN). Both the frequency of disturbed morphologies and the presence of AGN correlate with the host galaxy's IR luminosity (e.g., Sanders \\& Mirabel 1996). However, using measures such as concentration, rotational asymmetry ({\\it C}, {\\it A}; Conselice 2003), and near--neighbor frequency, some studies failed to find a correlation between X-ray--selected AGN and galaxy interactions (e.g., Grogin et al. 2003, 2005, hereafter G03, G05). Thus, galaxy interactions may correlate more strongly with IR--luminous AGN or may be less apparent for X-ray-luminous AGN. Recent galaxy merger simulations also suggest a connection between disturbed morphologies, AGN, and hard X-ray spectra (e.g., DiMatteo, Springel, \\& Hernquist 2005, hereafter DSH05; Hopkins et al.\\ 2005a, 2005b, hereafter H05a, H05b).  The Gini coefficient, which measures the distribution of flux amongst a galaxy's pixels ({\\it G}; Abraham 2003), and the second-order moment of the brightest $20\\%$ of a galaxy's flux ($M_{20}$; Lotz, Primack, \\& Madau 2004, hereafter LPM04), have been shown to be robust to high $z$ for determining the morphologies of potentially disturbed objects. In particular, $M_{20}$ is more sensitive to merger signatures than other measures of concentration. Morphologically undisturbed galaxies form a well-defined sequence, correlated with Hubble type, in the $G-M_{20}$ plane; merging and interacting galaxies are separated from this sequence by having a high {\\it G} for a given $M_{20}$ (LPM04).  In this work, we classify X-ray and IR--selected AGN host galaxy morphologies using {\\it G}, $M_{20}$, {\\it C}, {\\it A}, and visual inspection of {\\it Hubble Space Telescope} Advanced Camera for Surveys (ACS) images in the Extended Groth Strip (EGS; Davis et al.\\ 2006, hereafter D06; Lotz et al.\\ 2006). We also use Deep Extragalactic Evolutionary Probe 2 (DEEP2) spectroscopic velocities to determine the frequency of close kinematic galaxy pairs (Lin et al.\\ 2004, 2006). Throughout, we use $H_{0} = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda} = 0.7$, and $\\Omega_{M} = 0.3$. ",
        "conclusions": "The relatively high fraction of X-ray--selected AGN in kinematic pairs supports models that predict that galaxy interactions are responsible for the activation of some galactic nuclei (e.g., DSH05, H05ab). However, most AGN are found in undisturbed, isolated systems, implying that interactions may not be the main method by which AGN are turned on or that the timescales for interactions and AGN activity differ significantly from what is suggested by simulations.  X-ray and IR AGN selection methods select AGN with different host galaxy morphologies. A two-dimensional K--S test (FF87) is inconsistent with the X-ray-- and IR--selected AGN hosts being drawn from the same population ($p=5.4\\times10^{-3}$). IR selection via a power-law misses many AGN hosted by E/S0/Sa galaxies; although all X-ray--selected AGN with $I < 23.5$ and $0.2 \\le z < 1.2$ are detected by IRAC, most do not have power-law SEDs (see also Barmby et al.\\ (2006) who found that only $40\\%$ of {\\it Chandra} sources have IRAC power laws). X-ray selection misses many AGN hosted by Sc/d/Irr galaxies. Using the AGN with $I < 23.5$ and $0.2 \\le z < 1.2$ that are selected by both X-ray and IR methods, $f_{8 \\mu{\\rm m}}/f_{0.5-7 {\\rm keV}} = 2.4\\pm0.5$ (consistent with Elvis et al.\\ 1994, Table 15); this predicts detectable X-ray emission from the IR--selected AGN, yet many of these objects are not detected by {\\it Chandra}. The disparity between the morphologies of X-ray-- and IR--selected AGN suggests that the methods may select AGN at different stages or orientations with varied levels of gas, that galaxy light from the hosts of some X-ray--selected AGN may overwhelm the IR emissions (so that the measured IR emissions do not follow a power-law), or that star-forming field galaxies are contaminating the IR--selected AGN sample.  The observation of soft X-ray spectra from four X-ray--selected AGN hosted by morphologically interacting galaxies may contradict predictions from H05b that soft X-rays from AGN hosted by interacting galaxies should be highly attenuated. However, most of our X-ray--selected AGN are less luminous than the QSO simulated by Hopkins (2006ab). Though it is not surprising that many X-ray--selected AGN hosted by E/S0/Sa galaxies have soft spectra, soft X-ray--selected AGN hosted by interacting galaxies are unexpected. Dissipationless mergers would explain unobscured nuclei but would be unlikely to provide enough gas to trigger AGN. Thus, the observed interactions were probably not responsible for these AGN."
    },
    "0608/astro-ph0608662_arXiv.txt": {
        "abstract": "Fast heuristically weighted, or pseudo-$C_\\ell$, estimators are a frequently used method for estimating power spectra in CMB surveys with large numbers of pixels. Recently, Challinor \\& Chon showed that the E-B mixing in these estimators can become a dominant contaminant at low noise levels, ultimately limiting the gravity wave signal which can be detected on a finite patch of sky. We define a modified version of the estimators which eliminates E-B mixing and is near-optimal at all noise levels. ",
        "introduction": "Measurements of CMB polarization, reported by four experiments to date \\cite{DASI1,CBI1,CAPMAP,Boom03}, are rapidly improving in sensitivity. Future ground-based surveys will be optimized for sensitivity to the B-mode power spectrum, which offers a unique window into the physics of the early universe via the gravity wave signal \\cite{SZgw,KKSgw}, and can help break parameter degeneracies via the lensing signal \\cite{Stompor:1998zj,me2}.  On a practical level, two general methods for estimating power spectra from noisy maps have been used in experiments to date: maximum likelihood \\cite{BJK}, and pseudo-$C_\\ell$ estimators \\cite{WHGCl,MASTER,HGpolCl}. The strengths and weaknesses of these methods are complementary. Maximum likelihood power spectrum estimatation is optimal (in the sense that the Cramer-Rao inequality is saturated) but computationally expensive, requiring $\\bigoh(\\Npix^3)$ time and $\\bigoh(\\Npix^2)$ memory, and rapidly becomes infeasible for large maps. In contrast, pseudo-$C_\\ell$ estimators are very fast, suboptimal in principle, but frequently near-optimal in practice.  For this reason, they are currently the method of choice for surveys with large pixel counts.  In the context of CMB polarization, pseudo-$C_\\ell$ estimators have an additional drawback, as studied by \\citet{ChalChon}: the estimated BB power spectrum acquires a nonzero contribution from E-modes in the map. The pseudo-$C_\\ell$ construction includes a debiasing step (Eq.~(\\ref{eq:debias}) below) which removes the E-B mixing in the mean; however, the variance of the BB estimators still depends on the level of EE power. This is analagous to treating noise by subtracting the noise bias from each estimator; even though the bias from noise is removed in the mean, it still makes a contribution to the estimator variance. This contribution can dominate if the instrumental noise is sufficiently small; e.g. in \\cite{ChalChon} it is shown that for surveys with $\\fsky \\sim 0.01$, it limits the gravity wave signal which can be detected using pseudo-$C_\\ell$ estimators to $T/S\\sim 0.05$.  We will describe a modification to the pseudo-$C_\\ell$ construction, defining {\\em pure pseudo-$C_\\ell$ estimators} which separate E and B in the strongest possible sense. On a finite patch of sky, the BB estimator completely filters out E-modes, acquiring contributions only from the B-mode signal and noise. We will show that these estimators significantly improve power spectrum errors for noise levels $\\simle 20$ $\\mu$K-arcmin, and impose no limit on the gravity wave signal which can be detected. ",
        "conclusions": "For CMB surveys with many pixels, pseudo-$C_\\ell$ power spectrum estimation is currently the ``industry standard'', but has a shortcoming for future high-sensitivity polarization experiments: the EE signal acts as an extra source of noise on the BB power spectrum, becoming a significant contaminant at noise levels $\\simle 20$ $\\mu$K-arcmin. We have defined pure pseudo-$C_\\ell$ estimators, which eliminate this shortcoming while preserving the fast ($\\bigoh(\\ellmax^3)$) evaluation of the estimator and transfer matrix. These estimators completely filter out the E-mode component of the map; the estimated BB power receives contributions only from the B-mode part of the signal and from noise.  A property of the pure pseudo-$C_\\ell$ estimator is that the statistical weight is given by a combination of the weight function $W(x)$ and its first two derivatives. This imposes no restriction on the way the statistical weight is distributed (one can always choose $W(x)$ to give any desired statistical weight while satisfying the boundary conditions) but makes optimizing the weight function for a given noise distribution less intuitive. The optimization problem will be studied systematically in \\cite{ZS}, in the context of both pure and ordinary pseudo-$C_\\ell$ estimators.  We thank Niayesh Afshordi, Wayne Hu, Dragan Huterer, Bruce Winstein, and Matias Zaldarriaga for collaborative work and stimulating discussions. This work was supported by the Kavli Institure for Cosmological Physics through the grant NSF PHY-0114422."
    },
    "0608/astro-ph0608454_arXiv.txt": {
        "abstract": "The SuperWASP Cameras are wide-field imaging systems sited at the Observatorio del Roque de los Muchachos on the island of La Palma in the Canary Islands, and the Sutherland Station of the South African Astronomical Observatory.  Each instrument has a field of view of some $~$482 square degrees with an angular scale of 13.7 arcsec per pixel, and is capable of delivering photometry with accuracy better than 1\\% for objects having $V \\sim 7.0 - 11.5$. Lower quality data for objects brighter than $V \\sim 15.0$ are stored in the project archive. The systems, while designed to monitor fields with high cadence, are capable of surveying the entire visible sky every 40 minutes. Depending on the observational strategy, the data rate can be up to 100\\,GB per night. We have produced a robust, largely automatic reduction pipeline and advanced archive which are used to serve the data products to the consortium members.  The main science aim of these systems is to search for bright transiting exo-planets systems suitable for spectroscopic followup observations. The first 6 month season of SuperWASP-North observations produced lightcurves of $\\sim$6.7 million objects with 12.9 billion data points. ",
        "introduction": "In recent years, interest has grown in relatively small aperture and inexpensive wide-field imaging systems, essentially composed of large CCDs mounted directly to high-quality wide-angle camera optics.  The first prominent success of such an instrument was the spectacular discovery of the neutral sodium tail of comet Hale-Bopp \\citep{cremonese}, with a temporary purpose-built camera system (CoCam).  Since then, similar cameras have resulted in imaging of a gamma-ray burst during the burst period \\citep{akerlof}, and the first detection of the transits of an extra-solar planet in front of its parent star, HD\\,209458 \\citep{charbonneau}.  Such instruments are ideal for projects requiring photometry of bright but rare objects \\citep[see ][]{pinfield05}.    \\subsection{Extra-solar planetary transits}  The first extra-solar (exo-) planets were discovered in 1992 by pulsar timing experiments \\citep{wolszczan}. Whilst this technique is sensitive to the detection of terrestrial-sized planets, its limited applicability has restricted its use to just a few objects. \\citet{mayor} discovered the first exo-planet, 51\\,Peg, from optical radial velocity studies, and since that time the field has been dominated by this technique. One of the surprises of these surveys \\citep[e.g. ][]{marcy05} is the existence of a significant population of solar-type stars accompanied by relatively rapidly orbiting Jupiter-sized planets.  However, spectral measurements alone cannot determine unambiguously the true mass or radius of the planet as the orbital inclination is unknown.  These surveys have discovered Jupiter-sized objects in orbits out to 3\\,AU around 6\\% of the nearby Sun-like stars surveyed.  Of these, some 30\\% are {\\it Hot Jupiters} situated in $\\sim$4-day $0.05$\\,AU orbits, where the equilibrium temperature is $\\sim$ 1500\\,K.  About 10\\% of Hot Jupiters in randomly inclined orbits will transit their host star. Therefore, in random Galactic fields, roughly 1 in every 1000 solar-type stars should exhibit transits lasting roughly 2 hours with a period of a few days.  Given that the radius of Jupiter is $\\simeq$ 0.1\\,R$_{\\odot}$, these transits should result in a dimming of the parent star by $\\simeq$ 0.01 mag. The first transits were discovered in late 1999 \\citep{charbonneau}.  The $V = 7.7$ star HD\\,209458 is dimmed by 0.016\\,mag for 2 hours every 3.5 days, hence both proving the existence of the planet detected by a radial velocity search, and resulting in a precise measurement of its radius, mass and bulk density.  Because of the multiplexing advantage of imaging, this technique promises to be the fastest way of detecting exo-planets, and could over the next few years dictate which candidates are followed up by radial velocity studies (rather than {\\it vice-versa} as at present).  Initially, groups trying to find transits of exo-planets reported disappointing results. For example the Vulcan Project \\citep{vulcan2002} searched some 6000 stars, finding only 7 transit like variables. Followup observations of these showed them all to be stellar in origin.  More recent surveys have have had more success with published transit detections reported by the OGLE survey (Udalski et al., 2002a/b), TrES-1 \\citep{alonso}, HD\\,189733 \\citep{bouchy} and most recently, XO-1 \\citep{mccullough}. Part of the reason for the apparent lack of transits stems from the difficulty in obtaining photometry of sufficient numbers of solar and late type main sequence stars. To increase the numbers of stars sampled wide-field surveys have often concentrated on low galactic latitude fields. However, while the number of observable stars is undoubtedly increased, we suspect the stellar population in such surveys is dominated by more distant K giants.  \\citet{brown2003} showed that the number of {\\it bona fide} exo-planet transits (as opposed to stellar impostors) is consistent with expected numbers of binary and multiple stellar and exo-planet systems.   \\citet{horne2003} lists some 23 photometric transit projects either in operation or under construction at that time. While many of these are pencil-beam surveys from traditional telescopes, a number, no doubt encouraged by the relative cheapness of the equipment, are employing novel wide field cameras.   \\subsection{The WASP Consortium}  The Wide Angle Search for Planets (WASP) Consortium was established in 2000 by a group of primarily UK-based astronomers with common scientific interests.  In order to reduce the development cycle time and be on sky rapidly, we use commercially available hardware and hence limit development work as much as possible.  Our ethos is therefore quite distinct from other apparently similar projects e.g. the HAT Project \\citep{bakos}.  The WASP Consortium's first venture was the production of the WASP0 camera. Our experience with these types of systems stems from the CoCAM series of cameras at the Isaac Newton Group of Telescopes from 1996 -- 1998, one of which was responsible for the discovery of the so-called Sodium Tail in Comet Hale-Bopp 1995 \\citep{cremonese}. WASP0 is composed completely of commercial parts, and utilizes a Nikon 200\\,mm, f2.8 telephoto lens coupled to an Apogee AP10 CCD detector. WASP0 was used in 2000 and 2001 in La Palma and Kryoneri (Greece) respectively, and has been shown to easily detect the extrasolar transit of HD\\,209458b, amongst other variables (Kane et al., 2004, 2005a/b).  On the strength of the WASP0 success the Consortium was able to raise sufficient funding for a more ambitious project -- the multi-detector SuperWASP cameras. The limited development required is reflected in the aggressive project timescale: for the La Palma instrument funding was approved in March 2002 and first light achieved in November 2003, while for the South African Astronomical Observatory (hereafter SAAO) sited instrument funding was secured in April 2004 and first light in December 2005.  In this paper, we describe the SuperWASP facilities at the Observatorio del Roque de los Muchachos on La Palma (SuperWASP-N) and the recently commissioned system at the Sutherland Station of the SAAO (SuperWASP-S). Along with the hardware and data acquisition system, we outline the SuperWASP reduction pipeline and archiving system for the data products. ",
        "conclusions": "The original SuperWASP Camera was constructed at Observatory of the Roque de los Muchachos on la Palma in 2003. Since then we have also constructed a clone facility at SAAO which has recently seen first light.  Each system has a total field of view of $~$482 square degrees and is capable of monitoring the sky down to $\\sim 15$th magnitude every $\\sim 40$ minutes. For the La Palma facility, in the first season of operation the instrument was used solely to monitor fields with high cadence ($\\sim$10 minutes). In subsequent seasons the development of the dynamic scheduler will allow all sky monitoring interspersed with the high cadence exo-planet fields (we expect complete sky coverage once per night).  The reduction pipeline is a continued development from the WASP0 instrument now generalized and expanded to cater for both facilities. The rms plots show that the precision at the bright end of our magnitude range reaches 3-4 milli-magnitudes over an extended period and stars brighter than $\\sim$11.5 magnitude have errors less than 1\\%.  An average night's data of 50\\,Gb, produces some 6$\\times10^6$ records in the archive.  We have developed an archive system capable of storing tens to hundreds of terabytes of raw and processed data, the results of high-level analyses of those data, but doing so in a manner that provides rapid access to individual lightcurves and images. The central catalog is designed to remain flexible and responsive to evolving user requirements despite the burgeoning data volume. The choice of HTTP as an interface protocol allows the end-user a great deal of flexibility in the choice of access clients, and facilitates truly distributed, collaborative working within the Consortium.  From the start of operations in 2006 at both facilities we have begun an all-sky survey mode. As the exo-planet program will continue we have integrated these different survey modes (with only a small loss to the exo-planet survey efficiency) enabling once a night coverage of the entire visible sky.  To do this effectively we will eventually deploy a reduction pipeline at the sites themselves. By compromising performance slightly we will be able to use this to grade the quality of the night, directly feeding back to the dynamic scheduler in order to modify the observing queue, hence matching the science to the prevailing weather conditions. In principle this could also provide a real time alert system and if sufficient software resources become available we could use this to arrange followup observations on other nearby facilities.  The SuperWASP Cameras are now surveying large parts of the sky primarily in search of bright, transiting exo-planets. In the near future we will present large numbers of high quality light curves suitable for followup and expect to at least double the numbers of known transiting exo-planets. Interests of the Consortium members also cover other ``secondary'' science areas and publications in these areas will be forthcoming. For non-members of the consortium we also offer the opportunity to present ideas to the Project Steering Committee and hence achieve a Guest Observer status in this area."
    },
    "0608/gr-qc0608139_arXiv.txt": {
        "abstract": "The multiyear problem of a two-body system consisting of a Reissner-Nordstr\\\"om black hole and a charged massive particle at rest is here solved by an exact perturbative solution of the full Einstein-Maxwell system of equations. The expressions of the metric and of the electromagnetic field, including the effects of the electromagnetically induced gravitational perturbation and of the gravitationally induced electromagnetic perturbation, are presented in closed analytic formulas. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608038_arXiv.txt": {
        "abstract": "We present mid-infrared spectra of T Tauri stars in the Taurus star-forming region obtained with the {\\it Spitzer} Infrared Spectrograph (IRS). For the first time, the 5--36 $\\mu$m spectra of a large sample of T Tauri stars belonging to the same star-forming region is studied, revealing details of the mid-infrared excess due to dust in circumstellar disks. We analyze common features and differences in the mid-IR spectra based on disk structure, dust grain properties, and the presence of companions. Our analysis encompasses spectral energy distributions from the optical to the far-infrared, a morphological sequence based on the IRS spectra, and spectral indices in IRS wave bands representative of continuum emission. By comparing the observed spectra to a grid of accretion disk models, we infer some basic disk properties for our sample of T Tauri stars, and find additional evidence for dust settling. ",
        "introduction": "The Taurus-Auriga star-forming region is a relatively nearby \\citep[140 pc;][]{kenyon94b,bertout99} molecular cloud complex, which is characterized by lower extinction and more isolated, low-mass star formation than other star-forming regions at comparable distances, like Ophiuchus \\citep[][hereafter KH95]{kenyon95}. It is therefore fairly well-studied from wavelengths ranging from the X-rays to the radio and used to test theoretical predictions of the star formation process. Given the young age of the Taurus complex (1--2 Myr), many of its pre-main-sequence stars are still surrounded by material left over from their formation process and accrete gas and dust from circumstellar disks.  Dust grains in these disks are heated by irradiation from the star and from the accretion shocks at the stellar surface, as well as by local viscous dissipation \\citep{dalessio98}, and emit excess emission, i.e.\\ at flux levels above those expected from a stellar photosphere, at infrared wavelengths. Most disks have flared disk surfaces, i.e. the disk thickness increases more steeply than just linearly with radius; this flaring allows disks to absorb more radiation from the central source than a flat, thin disk, and thus to become hotter \\citep[see, e.g.,][]{kenyon87}.  Over time, accretion rates decrease, and, while the disk material is depleted, some of the dust and gas could be built into planets. Young, low-mass ($\\lesssim 2 M_{\\odot}$) stars surrounded by these accretion disks are referred to as classical T Tauri stars (CTTSs). Once accretion has virtually stopped and disks are either passive, i.e. only reprocessing stellar radiation, or already dissipated, a young star is called a weak-lined T Tauri star (WTTS). The distinction between CTTSs and WTTSs is based on the equivalent width of their H$\\alpha$ emission line (EW[H$\\alpha$]): CTTSs have broad and asymmetric H$\\alpha$ line profiles, which are generated in the magnetospheric accretion flows \\citep{muzerolle03a}, resulting in an EW[H$\\alpha$] larger than 10 {\\AA}, while WTTSs have an EW[H$\\alpha$] of less than 10 {\\AA}; in WTTSs, strong surface magnetic fields provide the energy for strong X-ray, UV, and H$\\alpha$ emission.  Based on an evolutionary sequence derived from the shape of the spectral energy distribution (SED) in the infrared \\citep{lada84, lada87}, young, low-mass stars are divided into three classes. Class I objects have rising SEDs over the infrared spectral range, while Class II objects have either flat or somewhat decreasing SEDs; Class III objects have little infrared excess, and their SEDs can be understood as reddened stellar photospheres for wavelengths $\\lesssim$ 10 $\\mu$m. Thus, Class II objects are characterized by the presence of mid-infrared excess emission due to dust in circumstellar disks, while Class III objects have, for the most part, already dissipated their disks, and therefore their infrared flux is close to photospheric levels. Classical T Tauri stars fall into the Class II category; weak-lined T Tauri stars can be considered Class II or Class III objects, depending on the presence of disks around them. When a young star has almost completely dissipated its circumstellar material, it is usually classified as a Class III object. Therefore, all Class III objects are WTTSs, but not all WTTSs are Class III objects, as defined by their infrared SEDs.  Mid-infrared spectra of T Tauri stars (mostly CTTSs) reveal the details of the thermal emission by dust grains in the inner regions of circumstellar disks, from a few tenths to several AU from the central object. By deriving the structure and composition of the dust in protoplanetary disks, we can study the conditions and processes leading to planetary formation. In particular, dust growth and settling are seen as key steps in the formation of larger bodies.  The {\\it} Spitzer Space Telescope \\citep{werner04} enables us, for the first time, to obtain spectra from 5 to 36 $\\mu$m of T Tauri stars covering a variety of flux densities, from tens of mJy to tens of Jy. As part of a larger Infrared Spectrograph\\footnote{The IRS was a collaborative venture between Cornell University and Ball Aerospace Corporation funded by NASA through the Jet Propulsion Laboratory and the Ames Research Center.} \\citep[IRS;][]{houck04} guaranteed-time observing program, we observed 150 young stellar objects in the Taurus-Auriga star-forming region, ranging in evolutionary state from the protostellar, or Class I, stage, when a forming star is still enshrouded by an envelope, to the Class II and III stages, when disk material is accreted and dissipated.  Our Taurus sample is largely based on the sample studied by KH95, who compiled and analyzed photometric data of known young stellar objects in the Taurus-Auriga star-forming region. Here we present 85 IRS spectra of Class II objects and 26 spectra of Class III objects; the remaining spectra of the Class I objects in our sample will be presented in a subsequent paper. We analyze the spectral energy distributions and IRS spectra of Class II and III objects in terms of common features and peculiarities arising from their disk structure, the properties of dust grains embedded within, and their surrounding environment. The large-scale environment has little impact on objects in Taurus, unlike young stars in OB associations or young clusters, where the cluster members are subject to the influence of high-energy radiation and powerful winds from the OB stars.  In $\\S$ \\ref{obs_data_reduction} we introduce our {\\it Spitzer} IRS observations and data reduction; in $\\S$ \\ref{sample} we present spectral energy distributions, from the optical to the far-infrared, of the Class II and III objects in our sample, as well as a morphological sequence of the IRS spectra of our 85 Class II objects; in $\\S$ \\ref{analysis} we compare our data with a set of accretion disk models to infer some disk properties, present a new median of mid-infrared fluxes of Class II objects in Taurus, and estimate upper limits for dust masses around our Class III objects. We discuss some of our results in $\\S$ \\ref{discuss}, and give our conclusions in $\\S$ \\ref{conclude}. ",
        "conclusions": "\\label{conclude}  We presented the mid-infrared spectra of 85 Class II and 26 Class III objects in Taurus, obtained with the Infrared Spectrograph on board the {\\it Spitzer Space Telescope}. This large sample of pre-main-sequence stars of similar age, belonging to the same star-forming region, allows us to draw the following conclusions about circumstellar disk evolution:  ${\\bullet}$ All CTTSs and some WTTSs display mid-IR excesses, which are generated in the inner (a few tenths to a few AU) parts of circumstellar disks. Accretion signatures like the H$\\alpha$ equivalent line width do not necessarily indicate the presence or absence of disks, but they usually reveal whether the innermost disk regions are still present. While accreting T Tauri stars have dusty disks extending from the dust sublimation radius outward and thus generate excess emission starting at near-IR wavelengths, WTTSs often have inner disk holes and thus lack near-IR excess emission. Since there are few WTTSs with disks, the transition period between an active (i.e., accreting) and a passive (i.e., purely reprocessing) disk stage seems short-lived.  ${\\bullet}$ The details of the mid-IR excess, and thus the distribution and properties of the dust in circumstellar disks, do not seem to be correlated with the multiplicity of the young star system. A sub-arcsecond ($\\lesssim$ 140 AU at the distance of Taurus) companion does not necessarily influence the inner disk; even spectroscopic binaries, like DQ Tau, seem to have substantial inner disks. We note that our IRS spectra cannot separate the contribution of close multiple systems, and therefore, in some cases, they show the combined mid-IR spectrum of circumprimary, circumsecondary, and circumbinary disk. The mid-IR spectra of multiple systems do not have any peculiarities that would set them apart from single stars; this could be explained if most disks in these systems are not influenced by the presence of companions and either the circumprimary disk clearly dominates the system or all components have similar disks, or if a circumbinary disk mimicked the behavior of a circumstellar disk around a single star.  ${\\bullet}$ Almost all Class II objects show 10- and 20-$\\mu$m silicate emission features in their mid-IR spectra, generated in the optically thin disk surface layer. There is a wide variety of feature strengths and crystallinity, independent of the spectral type and multiplicity of the objects. However, objects with weak 10-$\\mu$m silicate features usually have a decreasing SED in the mid-IR range, which could be an indication of dust growth and settling. In objects where the 10- and 20-$\\mu$m features are suppressed, dust grains must exceed radii of about 4 $\\mu$m.  ${\\bullet}$ After comparing the spectral indices $d \\log(\\lambda  F_{\\lambda})/d \\log (\\lambda)$ between 6 and 13 $\\mu$m and between 13 and 25 $\\mu$m of our Class II objects with accretion disk models that include dust settling, we infer that our data are consistent with these models. Within the assumptions of the adopted models, most T Tauri stars seem to have experienced substantial dust settling (and likely also dust growth), with depletions of factors of about 100--1000 of the standard dust-to-gas mass ratio in the disk atmosphere. Thus, we are able to estimate the amount of dust settling for a large sample of T Tauri stars, using the high sensitivity and spectral coverage of the IRS. This result will aid in the refinement of models predicting the evolution of dust in protoplanetary disks.  ${\\bullet}$ Most WTTSs in our sample are Class III objects, which have nearly photospheric fluxes in the infrared, indicating a dispersal of dust in their circumstellar disks. The upper limit for the mass of small, warm dust grains in these disks lies around 10$^{-4}$ lunar masses, which is similar to the dust mass of low-mass debris disks. In some cases nearby companions ($\\sim$ 10\\arcsec), which seem to have an infrared excess, contaminated the LL part of our Class III sources. It is interesting that objects of roughly the same age, formed in close proximity to each other, do evolve on different time scales. This result, as well as the diversity in mid-IR spectra observed for our Class II objects, indicates that age does not seem to be the dominating factor determining disk evolution, but rather the initial conditions of the star-forming core, like its mass and angular momentum.  While there likely is an evolutionary sequence from a classical T Tauri star with an accretion disk to a weak-lined T Tauri star with a passive disk to a T Tauri star with little or no disk left, there does not seem to be an absolute, evolutionary timescale regulating when the transitions from one stage to the next occur. Therefore, we can witness circumstellar disk evolution in young star-forming regions such as Taurus and infer the major steps in the process by analyzing a large sample of pre-main-sequence stars; this paper is a contribution to this effort by presenting and analyzing unprecedented data from the {\\it Spitzer Space Telescope}."
    },
    "0608/astro-ph0608512_arXiv.txt": {
        "abstract": "We investigate the luminosity functions of Lyman-break galaxies (LBG) at $z\\sim 4$ and 5 based on the optical imaging data obtained in the Subaru Deep Field (SDF) Project, a program conducted by Subaru Observatory to carry out a deep and wide survey of distant galaxies. Three samples of LBGs in a contiguous 875 arcmin$^2$ area are constructed. One consists of LBGs at $z\\sim 4$ down to $i'=26.85$ selected with the \\br\\ vs \\ri\\ diagram ($BRi'$-LBGs). The other two consist of LBGs at $z\\sim 5$ down to $z'=26.05$ selected with two kinds of two-color diagrams: \\vi\\ vs \\iz \\  ($Vi'z'$-LBGs) and \\ri\\ vs \\iz\\ ($Ri'z'$-LBGs). The number detected is $3,808$ for $BRi'$-LBGs, 539 for $Vi'z'$-LBGs, and 240 for $Ri'z'$-LBGs. The adopted selection criteria are proved to be fairly reliable by the spectroscopic observation of 63 LBG candidates, among which only 2 are found to be foreground objects. We estimate the fraction of contamination and the completeness for these three samples by Monte Carlo simulations, and derive the luminosity functions of the LBGs at rest-frame ultraviolet wavelengths down to $M_\\mathrm{UV}=-19.2$ at $z\\sim 4$ and $M_\\mathrm{UV}=-20.3$ at $z\\sim 5$. We find clear evolution of the luminosity function over the redshift range of $0\\le z\\le 6$, which is accounted for by a sole change in the characteristic magnitude, $M^\\ast$. The cosmic star formation rate (SFR) density at $z\\sim 4$ and $z\\sim 5$ is measured from the luminosity functions. The measurements of the integrated SFR density at these redshifts are largely improved, since the luminosity functions are derived down to very faint magnitudes. We examine the evolution of the cosmic SFR density and its luminosity dependence over $0\\le z\\lesssim 6$. The SFR density contributed from brighter galaxies is found to change more drastically with cosmic time. The contribution from galaxies brighter than $M^\\ast_\\mathrm{z=3}-0.5$ has a sharp peak around $z=3$ -- 4, while that from galaxies fainter than $M^\\ast_\\mathrm{z=3}-0.5$ evolves relatively mildly with a broad peak at earlier epoch. Combining the observed SFR density with the standard Cold Dark Matter model, we compute the cosmic SFR per unit baryon mass in dark haloes, i.e., the specific SFR. The specific SFR is found to scale with redshift as $(1+z)^3$ up to $z\\sim 4$, implying that the efficiency of star formation is on average higher at higher redshift in proportion to the cooling rate within dark haloes, while this is not simply the case at $z\\gtrsim 4$. ",
        "introduction": "\\label{sec:intro}  When and how galaxies formed is one of the primary questions in astronomy today. Observations of young galaxies at high redshifts are a straightforward approach to this problem. Over the past decade, it has become feasible to undertake large surveys of galaxies at high redshifts. What made this possible is the progress in observing technology including telescopes and detectors, and also sophistication of selection methods to locate high-redshift galaxies. Probably the most efficient method is the Lyman-break technique. It is a simple photometric technique based on the continuum features in rest-frame ultraviolet spectra redshifted into optical bandpasses, and requires only optical imaging in a few bandpasses. Pioneered by \\citet{guhathakurta1990}, this method has been successfully used to find many young, star-forming galaxies beyond $z\\sim 2$ \\citep[e.g.,][]{steidel1992, steidel1999, lehnert2003, iwata2003, ouchi2004, dickinson2004, sawicki2006}. \\citet{steidel2003} made spectroscopic observation for about 1000 photometrically selected $z\\sim 3$ galaxies and verified the usefulness of this method. High-redshift galaxies selected by this method are called Lyman-break galaxies (LBGs).  Since their first discovery in the 1990s, various properties of LBGs have been extensively studied. One of the most fundamental measurements is the luminosity function (LF) at rest-frame ultraviolet wavelengths (i.e., the number density of galaxies as a function of ultraviolet luminosity). Besides providing information on the number density of galaxies, the LF can be used to probe the star formation activity in the universe, because ultraviolet luminosity is sensitive to star formation. Thus the ultraviolet luminosity density derived by integrating the LF is related to the star formation rate (SFR) density in the universe. By obtaining the LF at various redshifts and examining its evolution, one can gain insights into the formation history of galaxies and the star formation history in the universe.  \\citet{steidel1999} derived the LF at $z\\sim 3$ from their large LBG sample, finding that their data are well fitted by a Schechter function with a low-luminosity slope $\\alpha$ of $-1.6$ down to $L\\simeq 0.1L^\\ast$. They extended their LBG search to $z\\sim 4$, to detect no significant evolution of the LF at bright magnitudes ($M<-21.1$) from $z\\sim 4$ to $z\\sim 3$. More recent observations have also derived the LF at $z\\sim 4$ and $z\\sim 5$ \\citep{iwata2003, ouchi2004, gabasch2004, sawicki2006}. The frontier of LBG searches now extends beyond $z\\sim 5$. \\citep{yan2003, stanway2003, bunker2004, bouwens2004, bouwens2005, shimasaku2005}.  However, it is not easy to construct a large sample of LBGs beyond $z\\sim 4$ covering a wide range of absolute magnitude because of their apparent faintness and low surface number density. In the first place, most of the studies to date are limited to bright LBGs. Consequently, the shape of the faint end of the LF has not been well determined. Thus, the estimation of the cosmic SFR density has a considerably large uncertainty (about a factor of 3 -- 10) due to the long extrapolation of the observed LFs to faint magnitudes \\citep[see][]{ouchi2004}. Exploring LBGs down to faint luminosities is necessary in order to determine the overall shape of the LF, and to measure the cosmic SFR density accurately. Although some surveys are extremely deep, they are restricted to very small areas (e.g. $\\sim 4$ arcmin$^2$ for the Hubble Deep Field, and $\\sim 10$ arcmin$^2$ for the Hubble Ultra Deep Field). Surveys based on such a small area probably suffer from cosmic variance, i.e., inhomogeneities in the spatial distribution of LBGs. Although deep and wide surveys of LBGs at $z\\sim 4$ have been made by \\citet{gabasch2004} and \\citet{sawicki2006} very recently, large samples for $z\\gtrsim 4$ are still very limited. It is therefore crucial to construct a new LBG sample from a survey of a similar depth and width and derive the LF and the cosmic SFR independently.  In this paper, we present a detailed study of LBGs at $z\\sim 4$ -- 5 based on the deep and wide-field images obtained in the Subaru Deep Field (SDF) Project \\citep{kashikawa2004}. The SDF Project is a program conducted by Subaru Observatory to carry out a deep galaxy survey over a blank field as large as $\\simeq 0.2$ square degree. Exploiting these unique data, we make the largest samples of LBGs at $z\\sim 4$ and 5, and derive the LFs down to very faint magnitudes; $M_\\mathrm{UV}=-19.2$ at $z\\sim 4$ and $M_\\mathrm{UV}=-20.3$ at $z\\sim 5$. This extends the LBG search by \\citet{ouchi2004} based on shallower and smaller area data of the same field some $\\sim 0.5$ mag further and $1.5$ times wider, leading to a $\\sim 3$ times increase in number of LBGs detected. These samples enable us to examine the behavior of the LF more accurately over a wide magnitude range and obtain more reliable measurements of the cosmic SFR density in the early universe. The LFs of LBGs at $z \\sim 4$ and 5 we derive are found to differ from those at the same redshift ranges obtained in some of the previous surveys.  Recently, the LF at ultraviolet wavelengths of present-day galaxies ($z\\sim 0$) has been derived very accurately from a large galaxy survey by Galaxy Evolution Explorer (GALEX) \\citep{wyder2005}. \\citet{arnouts2005} have derived the LFs at $0.2\\le z\\le 1.2$ also based on GALEX observations. Measurements at $z\\sim 6$ are now available \\citep{bouwens2005, shimasaku2005}. Combining the results at $z\\sim 4$ and 5 from this work with those at lower redshifts and $z\\sim 6$ from the literature, we investigate the evolutionary history of galaxies in terms of the star formation activity over $0\\le z\\lesssim 6$.  The outline of this paper is as follows. Section \\ref{sec:data} describes the data of the SDF Project used in this study. Section \\ref{sec:lbg} describes the selection of LBGs at $z\\sim 4$ and 5 from the photometric data. The contamination by interlopers and the completeness of the samples are also estimated. In Section  \\ref{sec:lf}, the rest-frame UV luminosity functions of LBGs at $z\\sim 4$ and 5 are derived, and compared with those by other authors. In Section \\ref{sec:dis} we discuss the evolution of the cosmic star formation activity over the redshift range $0\\le z\\le 6$. The evolution of the luminosity function and the cosmic star formation rate density is examined. Section \\ref{sec:sum} summarizes this study.  Throughout this paper, the photometric system is based on AB magnitude \\citep{oke1983}. The cosmology adopted is a flat universe with $\\mathrm{\\Omega}_m = 0.3$, $\\mathrm{\\Omega}_\\mathrm{\\Lambda} = 0.7$, and a Hubble constant of $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$. These values are consistent with those obtained from the latest CMB observations \\citep{spergel2003}. ",
        "conclusions": "  \\renewcommand{\\labelenumi}{(\\roman{enumi})} \\begin{enumerate}  \\item Clear evolution of the luminosity function is detected. The characteristic magnitude $M^\\ast$ changes strongly over the redshift range of $0\\le z\\le 6$. It brightens by about 1 mag from $z\\sim 6$ to $z\\sim 4$, remains unchanged to $z\\sim 3$ and then fades by about 3 mag from $z\\sim 3$ to $z\\sim 0$. On the other hand, the normalization factor, $\\phi^\\ast$, and the faint-end slope, $\\alpha$, are almost constant over $0.5\\lesssim z\\lesssim 6$.  \\item The measurements of the SFR density at $z\\sim 4$ and $z\\sim 5$ are greatly improved, since the luminosity functions are derived down to very faint magnitudes. Hence, the evolutionary behavior of the SFR density is more constrained. The cosmic SFR density does not drop from $z\\sim 3$ to $z\\sim 4$, and the decrease from $z\\sim 4$ to $z\\sim 5$, if any, cannot be larger than a factor of about five.  \\item The SFR density contributed from brighter galaxies changes more drastically with cosmic time. The contribution from galaxies brighter than $M^\\ast_{z=3}-0.5$ has a sharp peak around $z=3$ -- 4, while that from galaxies fainter than $M^\\ast_{z=3}-0.5$ evolves relatively mildly with a broad peak at earlier epoch. This luminosity-dependent evolution of the cosmic SFR density suggests that the galaxy population that contributes to the total SFR density in the universe varies with time, reflecting the evolution of the luminosity function.  \\item The specific SFR is found to scale with redshift as $(1+z)^3$ up to $z\\sim 4$, implying that the efficiency of star formation is on average higher at higher redshift in proportion to the cooling rate within dark haloes. It seems that this is not simply the case at $z\\gtrsim 4$. \\end{enumerate}"
    },
    "0608/hep-ph0608104_arXiv.txt": {
        "abstract": "{We present a new mechanism for baryogenesis, which links the baryon asymmetry of the universe to the dark matter density. The mechanism arises naturally in the Pentagon model of TeV scale physics.  In that context, it forces a re-evaluation of some of the assumptions of the model, and we detail the changes that are required in order to fit observations.}  \\received{} \\accepted{} \\preprint{hep-ph/{0608104}\\\\\\\\RUNHETC-06-20 \\\\SCIPP-06-10} \\begin{document} ",
        "introduction": "With a few exceptions\\cite{exceptions} most models of the early universe make no attempt to connect the observed baryon and dark matter densities.   Dark matter is usually assumed to consist of neutralinos or axions, and there is no connection between the properties of these particles, and the baryon asymmetry.   In this paper we will introduce a new class of models in which the solutions to these two problems are directly related.   These models are motivated by the Pentagon model of TeV scale physics\\cite{remodel}, where the mechanism we will discuss arises naturally.   Indeed, it is forced on us if we insist on finding a dark matter candidate within the low energy model itself.  The basic new structure in this class of models is an approximate symmetry, with current $J_{PB}^{\\mu}$ .   This symmetry is explicitly broken and an asymmetry $\\epsilon$ for $J_{PB}^0$ is generated in the very early universe.   The leading symmetry violating operator at low energy has dimension $D$ and comes from physics at a scale $M_b$.    The symmetry is also spontaneously broken at a scale $f$.    Finally, there is a coupling in the low energy effective Lagrangian of the form $${{1\\over \\Lambda^2}} J_{PB}^{\\mu} B_{\\mu} ,$$ where $B_{\\mu}$ is the ordinary baryon number current.  The asymmetry in $J_{PB}$ then acts as an effective chemical potential for ordinary baryon number.   Electroweak sphaleron processes will then generate a baryon asymmetry (and a compensating lepton asymmetry).    This is a form of spontaneous baryogenesis, a mechanism invented by Cohen and Kaplan\\cite{ck}.  At a lower energy scale, the effects of spontaneous and explicit violation of $J_{PB}$ set in.   These convert the energy stored initially in the asymmetry into a condensate of pseudo-Goldstone bosons, which can be the dark matter.   We can use the parameters of the model to fit both the baryon asymmetry, and dark matter density, in a variety of ways.  However, in the version of the model that arises from the Pentagon, all parameters but $\\epsilon$ and $M_b$ are determined in terms of other experimental quantities.   We find that, given the assumption in \\cite{remodel} that $M_b$ is at least as large as the neutrino seesaw scale ($5 \\times 10^{14}$ GeV)the baryon asymmetry is too large, or the dark matter density too small. We describe modifications of the model which remove this problem in section 4. ",
        "conclusions": "The Pentagon model suggests a novel form of spontaneous baryogenesis, which can tie together the dark matter density and the baryon asymmetry.   While the general idea works quite well, it does not work in the Pentagon model, unless we contemplate a scale of $ 10^8  \\rightarrow 10^{10} $ GeV for the leading irrelevant operator that violates penta-baryon number.   It remains to be seen whether one can invent a high energy extension of the model to generate this operator without generating standard model operators already ruled out by experiment.  If we stick with $10^{15}$ GeV as the scale for all irrelevant corrections to the Lagrangian, then we must abandon the penton theory of dark matter.   The penton is a light PNGB which might be produced in the laboratory, but its cosmological abundance is negligible.   It could still participate in the generation of the baryon asymmetry we observe.  The most plausible candidate for dark matter compatible with the Pentagon would then be a QCD axion.   This would also solve the strong CP problem."
    },
    "0608/astro-ph0608607_arXiv.txt": {
        "abstract": "We investigate the spins and shapes of over a million dark matter haloes identified at $z=0$ in the Millennium simulation. Our sample spans halo masses ranging from dwarf galaxies to rich galaxy clusters. The very large dynamic range of this \\lcdm cosmological simulation enables the distribution of spins and shapes and their variation with halo mass and environment to be characterised with unprecedented precision. We compare results for haloes identified using three different algorithms, and investigate (and remove) biases in the estimate of angular momentum introduced by both the algorithm itself and by numerical effects. We introduce a novel halo definition called the TREE halo, based on the branches of the halo merger trees, which is more appropriate for comparison with real astronomical objects than the traditional ``friends-of-friends'' and ``spherical overdensity'' algorithms.  We find that for this many objects, the traditional lognormal function is no longer an adequate description of the distribution, $P(\\lambda)$, of the dimensionless spin parameter $\\lambda$, and we provide a different function that gives a better fit for TREE and spherical overdensity haloes.  The variation of spin with halo mass is weak but detectable, although the trend depends strongly on the halo definition used.  For the entire population of haloes, we find median values of $\\lambmed=0.0367$--$0.0429$, depending on the definition of a halo.  The haloes exhibit a range of shapes, with a preference for prolateness over oblateness.  More massive haloes tend to be less spherical and more prolate.  We find that the more spherical haloes have less coherent rotation in the median, and those closest to spherical have a spin independent of mass ($\\lambmed \\approx 0.033$).  The most massive have a spin independent of shape ($\\lambmed \\approx 0.032$).  The majority of haloes have their angular momentum vector aligned with their minor axis and perpendicular to their major axis.  We find a general trend for higher spin haloes to be more clustered, with a stronger effect for more massive haloes. For galaxy cluster haloes, this can be larger than a factor of $\\sim 2$. ",
        "introduction": "The formation of galaxies is intimately linked to the acquisition and distribution of angular momentum. In the current cosmological paradigm, the inflationary \\lcdm model, cosmic structures grow hierarchically. Dark matter haloes form by the dissipationless gravitational collapse of material associated with peaks in the primordial density fluctuation field, growing, through mergers and smooth accretion, into objects with a wide range of masses at the present day. Galaxies form when baryons cool and condense near the centre of these haloes (\\citealt{1978MNRAS.183..341W}, \\citealt{1991ApJ...379...52W}).  They undergo mergers and tidal interactions along with their haloes, giving rise to the rich spectrum of galaxy types and environments that we see today.  Understanding the generation and evolution of the angular momentum of dark matter haloes is a prerequisite for understanding the angular momentum and morphology of galaxies. For example, the distribution of halo spins is a basic input to models of galaxy formation (e.g. \\citealt{1998MNRAS.295..319M}, \\citealt{1998ApJ...507..601V}, \\citealt{2000MNRAS.319..168C}).  The early evolution of the angular momentum of a density perturbation is adequately described by the linear tidal torque theory (see e.g. \\citealt{1949pca..conf..195H}, \\citealt{1969ApJ...155..393P}, \\citealt{1970Ap......6..320D}, \\citealt{1984ApJ...286...38W}, \\citealt{1996MNRAS.282..436C}, \\citealt{2000ApJ...532L...5L}). However, as work such as that of \\cite{1984ApJ...286...38W} and \\cite{2002MNRAS.332..325P} has shown, the non-linear effects inherent in the formation of large-scale structure lead to large quantitative disagreements between the predictions of the tidal torque theory and the angular momenta found in $N$-body simulations of dark matter haloes.  $N$-body simulations provide the way to progress beyond the linear regime.  As computing power has improved, so has the scale and resolution of simulations.  Very early numerical studies of the angular momentum of ``proto-galaxies'' were performed by \\cite{1971A&A....11..377P} (with $\\Np\\sim 100$ particles) and \\cite{1979MNRAS.186..133E} ($\\Np=1000$), and led the way to the analysis of the spins and shapes of CDM haloes in more sophisticated simulations (\\citealt{1985ApJ...292..371D, 1987ApJ...319..575B, 1988ApJ...327..507F}; all with $\\Np=32\\,768$). \\cite{1992ApJ...399..405W} used a much larger simulation ($\\Np\\sim 10^6$) and focused particularly on the details of the distributions of halo spins and shapes, and their relationship through the alignment of the halo angular momentum vector.  \\cite{1996MNRAS.281..716C} also investigated the shapes and spins of dark matter haloes, in addition to various other aspects of halo structure.  An early study by \\cite{1987ApJ...319..575B} examined the relationship between spin and the spatial clustering of haloes, as measured by the two-point correlation function. Later, \\cite{1999MNRAS.302..111L} examined the environmental dependence of halo properties, and found no correlations with halo spin. More recently, \\cite{2002A&A...395....1F} have carried out a mark correlation function analysis to investigate how spin varies with halo pair-separation.  They found that neighbouring cluster pairs tend to have higher spins than the average.  Recent years have seen a large amount of work on the analysis of haloes in \\lcdm simulations.  Halo shapes and their variation with mass were investigated by \\cite{2002sgdh.conf..109B}, \\cite{2005ApJ...629..781K}, \\cite{2006ApJ...646..815S} and \\cite{2006MNRAS.367.1781A}. In agreement with previous studies, halo spin was found to vary little, if at all, with halo mass. The relationship between halo shape and spin was investigated by \\cite{2005ApJ...627..647B}, \\cite{2005ApJ...634...51A}, \\cite{2006ApJ...646..815S} and \\cite{2006EAS....20...25G}.  While this paper was being completed, independent analysis of halo properties, investigating halo concentrations, spins and shapes in a series of simulations, was posted by \\cite{2006astro.ph..8157M}.  Using the the 10-billion particle Millennium simulation of the evolution of dark matter in the \\lcdm cosmology (\\citealt{2005Natur.435..629S}), we re-examine, in this paper, some of the shape and spin properties of dark matter haloes previously considered. Our analysis improves upon earlier work because the millions of haloes that formed in this simulation provide unprecedented statistical power. This allows us, for example, to quantify the distribution of halo spins and the relationship between spin, halo mass and shape with a precision that has not hitherto been possible.  Unlike previous work, we consider different ways to identify haloes in the simulations; it turns out that the details of halo definition and selection can have a strong impact on the results. Finally we investigate how halo clustering depends on spin and shape for haloes of different masses.  This paper is structured as follows.  Section \\ref{s:dmh} provides a description of the Millennium simulation itself, and the various halo properties we shall be investigating. Section \\ref{s:halocats} describes the construction of the catalogues whose haloes we investigate, including the group-finding algorithms and halo selection criteria.  These we use to remove haloes whose properties are unreliable or biased, due to both numerical effects and the group-finding algorithms themselves.  The main results of this paper are presented in \\S\\ref{s:results}, where we describe the distribution of halo spins as a function of mass and shape, and examine its effect on halo clustering.  Our conclusions are presented in \\S\\ref{s:conc}. Finally, the Appendix shows various examples of haloes that illustrate the effects of the group-finders. ",
        "conclusions": "\\label{s:conc} The huge size and high resolution of the Millennium simulation \\citep{2005Natur.435..629S} makes it possible to determine the properties of dark matter haloes in the \\lcdm cosmology with unprecedented statistical power. In this paper, we have concentrated on the spins and shapes of dark matter haloes, ranging in mass from those of dwarf galaxies to those of rich clusters. We have investigated the distribution of the spin parameter, $\\lambda$, its dependence on mass, the distribution of shapes, the relationship between shape and spin, and the environmental dependence of spin and shape. We provide accurate fitting formulae for several quantities of interest.  While many of the properties we have investigated here have been studied in earlier simulations going back over twenty years, a novel aspect of our work is the analysis and comparison of haloes identified in different ways. Alongside the traditional ``friends-of-friends'' (FOF) algorithm of \\cite{1985ApJ...292..371D} and the ``spherical overdensity'' (SO) algorithm of \\cite{1994MNRAS.271..676L}, we have introduced a new halo definition, the TREE haloes, which are perhaps the most appropriate when carrying out comparisons of the simulation results with galaxy and cluster data. The TREE haloes are defined as branches of the halo merger trees, in which special care has been taken to identify physical haloes by separating objects that are artificially and transiently linked together.  Each groupfinder results in more than $15$ million haloes at $z=0$.  The TREE halo catalogue was further `cleaned' in two ways. Firstly, to remove any remaining spurious objects, we applied a cut in the instantaneous virial ratio, ($-0.5 \\le \\frac{2T}{U}+1 \\le 0.5$). Secondly, to remove objects whose angular momentum is biased due to particle discreteness, we considered only haloes with more than $300$ particles, as indicated by convergence tests. Our final cleaned halo catalogues consist of $>10^6$ haloes at $z=0$.  We find that the distribution of the dimensionless spin parameter, $P(\\lambda)$, is poorly fit by a lognormal when this many objects are considered.  The function given by Eqn.~\\ref{e:arjfn} provides a much better description of the data.  Although the distribution of $\\lambda(\\Mh)$ is fairly broad, there is a clear trend of the median spin with halo mass, with more massive haloes spinning more slowly. However, the strength and shape of the trend is significantly different for different halo definitions.  The cubic polynomial of Eqn.~\\ref{e:poly} provides a very good fit to the median spin of the TREEclean halo catalogue, over a factor of $\\sim 10^3$ in halo mass.  We analysed the shapes of the haloes, and found, as in previous studies, that there is a broad distribution of shapes with a slight preference for prolateness over oblateness.  More massive haloes are less spherical and more prolate in the median, although the data span a large fraction of the available shape parameter space. We fit broken lines to the trends with $\\log\\Mh$ of the median sphericity axis ratio $s=c/a$ and the median triaxiality parameter $\\mathcal{T}$ (Eqn.~\\ref{e:broken}).  Both these quantities exhibit a change of behaviour at a galactic mass scale, $\\Mh\\approx 2\\times 10^{12}\\munit$, where the gradient of the fit changes sign, with haloes becoming increasingly aspherical and more prolate with increasing mass.  The rounder haloes have less coherent rotation, with a median spin that is independent of mass ($\\lambmed(s\\ga0.9)\\approx 0.033$). The most massive haloes have a median spin that is independent of sphericity ($\\lambmed(s)\\approx 0.032$).  However, there is significantly less correlation between the nature of halo triaxiality (prolateness {\\it vs} oblateness) and the spin parameter. Although the haloes are far from being rotationally supported, there is a strong preference for the spin vector to be aligned parallel to the halo minor axis and to be perpendicular to the major axis. The tendency for the spin to be perpendicular to the intermediate axis is significantly weakened by the prevalence of prolate shapes for which there is a near degeneracy between the intermediate and minor axes.  We find a clear signal that the spins and shapes of haloes are sensitive to the cosmological environment: more rapidly rotating haloes of a given mass are more strongly clustered. The strength of this effect increases with halo mass. It is weak for subgalactic and galactic haloes, but can be larger than a factor of $\\sim 2$ for galaxy cluster haloes.  A similar effect is seen when examining halo shapes: more spherical haloes are more strongly clustered, with a greater signal at higher masses.  Our result adds further evidence to the recent finding by \\cite{2005MNRAS.363L..66G}, also from analysis of the Millennium Run, that the internal properties of haloes depend not only upon their mass but also upon the environment in which they form.  The huge number of haloes in the \\lcdm Millennium simulation enables us to characterise the distribution of halo spins, and their relation to halo mass, shape and clustering, with unprecedented precision. However, we have also shown the significance of a careful halo definition.  The properties of haloes defined and identified in different ways are noticeably different, and it is important to make the appropriate choice for a given application.  For comparisons with real data, we recommend using the new class of ``TREE'' haloes which we have investigated in this work."
    },
    "0608/astro-ph0608094_arXiv.txt": {
        "abstract": "The progress over the last years in modelling the atmospheres and winds of PN central stars is reviewed. We discuss the effect of the inclusion of the blanketing by millions of metal lines in NLTE on the diagnostics of photospheric and stellar wind lines, which can be used to determine stellar parameters such as effective temperature, gravity, radius, mass loss rate and distance. We also refer to recent work on the winds of massive O-type stars, which indicates that their winds are possibly inhomogeneous and clumped. We investigate implications from this work on the spectral diagnostics of PN central stars and introduce a method to determine wind clumping factors from the relative strengths of H$_{\\alpha}$ and HeII 4686. Based on new results we discuss the wind properties of CSPN. ",
        "introduction": "Since many decades model atmospheres have been a fundamental tool to understand the physical nature of the PN Central Stars and the ionization and emission of their surrounding nebulae. After the pioneering work by Aller (1948), which revealed the importance and enormous potential of CSPN spectroscopy, and Heap (1977, and references therein), which provided the first quantitative spectral analyses based on model atmospheres, the field was advanced by M\\'endez et al. (1982), Kudritzki and M\\'endez (1987), M\\'endez et al. (1988). In this work, high quality spectra obtained with new 4m-class telescopes and very efficient spectrographs and detectors were analyzed in detail using a new generation of hydrostatic, planeparallel NLTE model atmospheres to determine effective temperatures, gravities and helium abundances. This work demonstrated nicely that O-type CSPN form an evolutionary sequence in the (log g, log T$_{eff}$)-plane and that the gravities and temperatures determined spectroscopically could be used to estimate stellar masses, radii, luminosities and distances by comparison with the prediction of post-AGB evolution and the core mass - luminosity relationship.  While this new concept seemed compelling, there were clear indications of quantitative deficiencies. The masses determined seemed systematically larger than White Dwarf masses and some of the objects (such as NGC 2392) had unrealistically high masses. The model atmospheres used, though in NLTE and certainly state-of-the art, did not include the opacities of metal lines and they also neglected the effects of stellar winds and spherical extension, which were suspected to be substantial, in particular for cooler objects, where the gravities are lower, and more massive objects, which are closer to the Eddington limit.  Indeed, Perinotto (1987) and Kudritzki and M\\'endez (1987) stressed the importance of stellar winds not only for the evolution of CSPN but also for their diagnostics. With a new generation of ``unified model atmospheres'', which included the effects of stellar winds and spherical extension, many CSPN were re-analysed and, indeed, somewhat lower masses were found (Kudritzki and M\\'endez, 1992, Kudritzki et al., 1997). In addition, the mechanical momenta of the stellar winds determined were in rough agreement with general scaling relations obtained from the theory of radiation driven winds and compared to the momenta of massive O-stars, however the scatter around this relationship was large.  The major remaining model atmosphere deficiency at this stage was the neglect of metal line opacity, which - if included - needed to be calculated in NLTE, certainly a formidable problem. This problem has been overcome in recent years and a wide variety of very efficient model atmosphere codes does exist now taking into account effects of millions of metal lines in NLTE and the velocity fields of stellar winds together with spherical extension. These new codes allow for detailed studies of the UV spectra and a re-analysis of the optical spectrum now with the inclusion of line-blanketing. This is the subject of the review presented here. We will focus on low gravity, relatively cool O-type CSPN. WR-type objects and objects of higher gravity are discussed in other reviews of these proceedings. ",
        "conclusions": "\\label{sec:disfu}  Comparing our results to Pauldrach et al. (2004) we find reasonable agreement for the effective temperatures for all cases (except He 2-108, which we find to be 17 percent cooler). This is very satisfying given the fact that different spectroscopic techniques were used for the determination of T$_{eff}$. However, in all but one case (He 2-131) the gravities obtained through our fitting of the Balmer lines are substantially lower. If one corrects for the temperature dependence of the Balmer lines (see Fig.~\\ref{fig1}) and the slightly different temperatures obtained, our gravities are on average 0.3 dex smaller. The reason is clearly the completely different approach used to determine gravities, as explained in section 2. Pauldrach et al. rely on the hydrodynamic simulation of stellar winds, whereas we have have used the classical spectroscopic concept of Balmer line fitting. This will need further investigation.  For the determination of masses, radii, luminosities, and distances we have used again the classical approach of using the post-AGB core mass - luminosity relationship. Compared to Kudritzki et al. (1997) our masses are generally smaller in agreement with what we expect from section 3, however some of them (IC 418, NGC 2392) are still uncomfortably high.  The mass loss rates determined are uncertain because of the effects of stellar wind clumping. However, if we take into account that mass-loss rates determined with H$_{\\alpha}$ scale with the stellar radius adopted as $\\dot{M} \\propto R^{3/2}$ (Kudritzki and Puls, 2000) and compare with Pauldrach et al. we find agreement within a factor of two except for IC 418. Fig.~\\ref{fig4} shows the CSPN wind momenta of our study compared to massive O-stars. There are two ways to interprete this plot. One is that CSPN form a convincing extension of the wind momentum - luminosity relationship of massive O-stars towards lower luminosities. Another one is that within the luminosity range of CSPN alone there is no clear relationship between wind momentum and luminosity (see also Tinkler and Lamers, 2002). Whether this is because of the uncertainties of mass-loss rate diagnostics or of the luminosity determinations, or both, will need further investigation.  It is interesting to compare this result with stellar wind models obtained with the theory of line driven winds. This is also done in Fig.~\\ref{fig4}, where we use the method by Kudritzki (2002) to calculate wind momenta for the CSPN of Tab.~\\ref{tab1}. With a few exceptions the theoretical momenta are in the right ballpark.  Several steps need to be undertaken for future work. The first is to extend the work presented here and to re-analyse the UV-spectra of the sample and to compare in detail to see whether or not the results obtained from the two spectral windows are compatible, in this way addressing the original point made by Pauldrach et al. (2004). Taking into account the effects of clumping through the diagnostics of the PV and similar lines will be crucial.  A very important issue is the compatibility of stellar wind hydrodynamics with the stellar parameters derived by our method. While the wind momenta seem to agree within the observational errors, the terminal velocities calculated are too small in many cases for the parameters obtained by us. Whether this is a deficiency of stellar wind hydrodynamics or of the Balmer line diagnostics used, remains to be investigated. We note that one modification to be made in the wind hydrodynamics is the inclusion of clumping factors. It will be interesting to see whether this will help to resolve the discrepancy."
    },
    "0608/astro-ph0608577_arXiv.txt": {
        "abstract": "% We discuss CO spectral line data of a volume-limited sample of 23 nearby ($z<0.03$) low luminosity radio galaxies, selected from the B2 catalogue. Most of such objects (16/23) have HST imaging.\\\\ Our aim is to establish the distribution of molecular gas masses in low luminosity radio galaxies, in comparison with other radio source samples, confirm the suggestion that the CO is in ordered rotation, determine its relation to the dust disks observed in these objects with HST and establish targets for future interferometric imaging. ",
        "introduction": "%  The fuelling of relativistic jets in radio-loud active galactic nuclei is still not fully understood.  It is clear that energy is released in the vicinity of a supermassive black hole, but whether the mechanism is direct electromagnetic extraction of rotational kinetic energy (e.g. Blandford \\& Znajek 1977; Koide et al. 2002) or more closely related to the process of accretion (e.g. Blandford \\& Payne 1982; Hujeirat et al. 2003) remains a matter of debate. The ratio of energy flux in jets to that radiated by an accretion disk varies by orders of magnitude between different classes of AGN. Here, we are concerned with low-luminosity (FRI) radio galaxies (Fanaroff \\& Riley 1974). These are known to contain very massive black holes, but show very little evidence for emission from accretion disks, their nuclear luminosities being a small fraction of that expected from accretion at the Eddington rate, $L/L_{\\rm Edd} < 10^{-4}$ (Marchesini, Celotti, \\& Ferrarese 2004). They form the parent population for nearby BL Lac objects and must therefore produce highly relativistic jets on sub-pc scales. There is direct evidence for relativistic motion on parsec scales in FR\\,I jets (Giovannini et al., 2001) and for smooth deceleration from relativistic to sub-relativistic speeds on scales of 1 -- 10\\,kpc (Laing et al., 1999). FR\\,I radio sources are located in fairly normal elliptical galaxies, invariably containing hot, X-ray emitting plasma (thought to confine the jets on large scales), but little ionized, line-emitting material at $\\approx10^4$\\,K. It has recently become clear that they may also contain substantial amounts of cool gas and dust (e.g. de Ruiter et al. 2002; Verdoes Kleijn \\& de Zeeuw 2005; Lim et al. 2000, 2003). Dust is observed in 53\\% of the B2 sample of nearby radio galaxies (mostly FRI; see below) and the dust mass is correlated with radio power (de Ruiter et al., 2002). There is also a connection between the dust-lane morphology (disk/irregular) and the presence of jets, and some tendency for dust lanes and jets to be orthogonal (de Ruiter et al. 2002; Verdoes Kleijn \\& de Zeeuw 2005). These associations argue that accretion of cool gas may indeed power the radio jets. The next step is to understand the dynamics of the cool gas. Lim et al. (2000) detected $^{12}$CO (1 $\\rightarrow$ 0) and (2 $\\rightarrow$ 1) emission from the FR\\,I radio galaxies 3C\\,31 and 264 with the IRAM 30m Telescope and established that the line profiles indicate disk rotation. Interferometric observations of 3C\\,31 by Okuda et al. (2005) showed that the CO coincides spatially with the dust disk observed by HST (Martel et al., 1999) and is in ordered rotation.  These authors suggest that the cool gas is in stable orbits. \\\\ In order to increase the number of FR\\,I radio galaxies with CO observations and so to improve our knowledge of their molecular gas properties, we are studying a volume-limited sample of 23 nearby ($z<0.03$) low luminosity radio galaxies, selected from the B2 catalogue (Colla et al., 1975). We notice that for 16 of such objects HST imaging is available (Capetti et al., 2000). The CO properties of this sample are compared to the 23 $z<0.03$ 3C radio galaxies studied by Lim et al. (2003) and to the 18 $z<0.0233$ (or $v<7000$ km/s) UGC galaxies with radio jets, studied by Leon et al. (2003). We notice that the three samples are partially overlapping.  \\begin{table}[t] \\caption{CO line measurements} \\footnotesize \\vspace{0.3truecm} \\centering \\begin{tabular}{crcrrcrr} \\hline \\multicolumn{1}{c}{Source} & \\multicolumn{3}{c}{$^{12}$CO(1$\\rightarrow$0)} & \\multicolumn{3}{c}{$^{12}$CO(2$\\rightarrow$1)} & \\multicolumn{1}{c}{$\\log{M_{H_2}}$}\\\\ \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{$\\Sigma T_a^* dv$}  & \\multicolumn{1}{c}{$\\Delta v_{\\tt FWHM}$} & \\multicolumn{1}{c}{$\\frac{T_{\\rm peak}}{T_{\\rm rms}}$} & \\multicolumn{1}{c}{$\\Sigma T_a^* dv$}  & \\multicolumn{1}{c}{$\\Delta v_{\\tt FWHM}$} & \\multicolumn{1}{c}{$\\frac{T_{\\rm peak}}{T_{\\rm rms}}$} & \\multicolumn{1}{c}{} \\\\ \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{K km/s}  & \\multicolumn{1}{c}{km/s} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{K km/s}  & \\multicolumn{1}{c}{km/s} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{M$\\odot$} \\\\ \\hline $0120+33$ & $<0.35$ & - & - & $<0.61$ & - & - & $<7.7$ \\\\ $0149+35$ & $1.34$ & 545 & 5.4 & $3.15$ & 551 & 7.7 & $8.3$ \\\\ $0258+35$ & $0.88$ & 382 & 6.6 & $0.21$ & 127 & 2.4 & $8.1$ \\\\ $0326+39$ & $<0.35$ & - & - & $0.64$ & 171 & 1.9 & $<8.0$ \\\\ $0331+39$ & $1.35$ & 807 & 3.6 & $0.90$ & 679 & 3.4 & $8.5$ \\\\ $1122+39$ & $6.84$ & 670 & 23.0 & $5.93$ & 672 & 16.7 & $8.2$ \\\\ $1217+29$ & $0.58$ & 346 & 3.3 & $<0.87$ & 335 & 3.1 & $6.1$ \\\\ $1321+31$ & $<0.35$ & - & - & $1.24$ & 297 & 2.7 & $<7.7$ \\\\ $1615+35$ & $<0.35$ & - & - & $0.52$ & 129 & 4.1 & $<8.2$ \\\\ \\hline \\multicolumn{8}{l}{$M_{H_2}$ derived from CO(1$\\rightarrow$0) measures and assuming $H_0=100$ km/s/Mpc} \\end{tabular} \\end{table} ",
        "conclusions": ""
    },
    "0608/astro-ph0608211_arXiv.txt": {
        "abstract": "We present initial results from a time-series BVI survey of two fields in \\gal\\ using the Advanced Camera for Surveys onboard the Hubble Space Telescope. This galaxy was selected because of its accurate maser-based distance, which is anticipated to have a total uncertainty of $\\sim3\\%$. The goal of the HST observations is to provide an absolute calibration of the Cepheid Distance Scale and to measure its dependence on chemical abundance (the so-called ''metallicity effect'').  We carried out observations of two fields at different galactocentric distances with a mean abundance difference of 0.5~dex. We discovered a total of \\ntot\\ Cepheids with periods ranging from 4 to 45 days (the duration of our observing window). We determine a Cepheid distance modulus for \\gal\\ (relative to the LMC) of $\\Delta\\mu_0=$\\dmua~mag. Given the published maser distance to the galaxy, this implies $\\mu_0(LMC)=\\mulmc$~mag or $D(LMC)=\\dlmc$~kpc. We measure a metallicity effect of $\\gamma=\\metz$~mag~dex$^{-1}$. We see no evidence for a variation in the slope of the Period-Luminosity relation as a function of abundance.  We estimate a Hubble Constant of $H_0=\\hnod$\\ksm\\ using a recent sample of 4 well-observed type Ia SNe and our new calibration of the Cepheid Distance Scale. It may soon be possible to measure the value of $H_0$ with a total uncertainty of 5\\%, with consequent improvement in the determination of the equation of state of dark energy. ",
        "introduction": "During the last 15 years, the {\\em Hubble Space Telescope} ({\\em HST}) has been used to discover $\\sim 10^3$ Cepheid variables in $\\sim 30$ galaxies with $D \\lesssim 25\\;$Mpc, mostly through $V$- and $I$-band observations carried out with the WFPC2 instrument. The distance moduli to these galaxies have been determined through the use of a fiducial Cepheid Period-Luminosity relation (P-L) based on observations of variables located in the Large Magellanic Cloud. Several secondary distance indicators (such as type Ia SNe, the Tully-Fisher relation, the Surface Brightness Fluctuation method) have been calibrated based on these Cepheid distances. As a result of these investigations, there is some agreement that H$_0$ is about 70\\ksm, perhaps with as little as 10\\% uncertainty \\citep{freedman01}. However, two significant sources of systematic error stand out.  First, the entire Cepheid Distance Scale is underpinned by the distance to the Large Magellanic Cloud (LMC). The distance to that galaxy is used to establish the absolute calibration of the Cepheid P-L relations, and its uncertainty dominates the calibration of any secondary distance indicator. The suitability of the LMC for this purpose is problematic, since independent estimates of its distance disagree by as much as 0.5 mag, or 25\\% \\citep{benedict02}. Additionally, the internal structure of the galaxy along the line of sight remains poorly understood \\citep{nikolaev04,vandermarel01}. Faced with this situation, most Cepheid-based determinations of $H_0$ have adopted $\\mu_{LMC}=18.5\\pm0.1$~mag, which corresponds to a distance of $D_{LMC}=50.1\\pm2.3$~kpc.  Second, the effect of metal abundance on the Cepheid P-L relation is controversial. Several independent methods for an observational determination have yielded a variety of results \\citep{sasselov97,kochanek97,kennicutt98,sakai04} with the {\\it opposite sign} to what has been predicted by some theoretical investigations \\citep{fiorentino02}, which also suggest a sensitivity to helium as well as metal content. Furthermore, the use of $V$ and $I$ photometry alone in previous HST surveys makes it difficult to disentangle the effects of reddening and metallicity and adds uncertainty to the determination of Cepheid distances.  We wish to establish a new Cepheid Distance Scale anchor galaxy, \\gal, for which accurate geometric estimates of distance are available. \\citet{herrnstein99} estimated its distance modulus to be $\\mumsr$~mag, and it is anticipated that \\citet{humphreys07} will reduce the total uncertainty of that estimate to $\\la 3\\%$. Our goal is even more compelling in light of the recent {\\it WMAP} results \\citep{spergel06} because many cosmological parameters depend sensitively on $H_0$ \\citep[e.g.][]{eisenstein04,tegmark04,hu05}. An accurate geometric distance to \\gal\\ can also be used to directly calibrate secondary distance indicators, such as the tip of the red giant branch (TRGB).  This paper contains the first results of our project: deep time-series $\\bvi$ photometry of two fields in \\gal\\ and the discovery and analysis of Cepheid variables. The paper is organized as follows: \\S 2 contains details of the observations, data reduction and photometry, and the search for variables; \\S 3 describes the selection criteria and the Cepheid samples; \\S 4 presents the determination of a Cepheid distance to \\gal, a measurement of the metallicity dependence of Cepheid-based distances, and a discussion of our results.  Throughout the paper, we denote random (statistical) uncertainties with a subscript {\\it r} and systematic uncertainties with a subscript {\\it s}, i.e., $\\pm0.10_r\\pm0.10_s$~mag. ",
        "conclusions": "The five main results presented in this paper are the following:  \\vspace*{3pt}  \\noindent {\\bf 1.}\\ We discovered $\\ntot$ Cepheid variables in two fields located within the galaxy \\gal, with accurately calibrated $\\bvi$ photometry in twelve epochs per band.  \\noindent {\\bf 2.}\\ We determined a relative distance modulus between \\gal\\ and the Large Magellanic Cloud, based on Cepheid variables, of $\\Delta\\mu_0=\\dmu$~mag.  \\noindent {\\bf 3.}\\ We determined a relative distance modulus between these two galaxies, based on the Tip of the Red Giant Branch method, of $\\Delta\\mu_{0,TRGB}=10.87\\pm0.04_r$~mag.  \\noindent {\\bf 4.}\\ We measured a metallicity dependence of the Cepheid distance scale of $\\gamma=\\metz$~mag dex$^{-1}$.  \\noindent {\\bf 5.}\\ Our observations are best fit with P-L relations that do not invoke changes in slope as a function of abundance."
    },
    "0608/astro-ph0608027_arXiv.txt": {
        "abstract": "We derive composite luminosity functions (LF) for galaxies in groups and examine the behaviour of the LF as a function of group luminosity (used as an indicator of group or halo mass). We consider both the entire galaxy population and split galaxies into red and blue (quiescent and star forming) samples, in order to examine possible mechanisms behind observed variations of galaxy properties with environment. We find evidence that $M^*$ brightens and $\\alpha$ steepens with group luminosity, until a threshold value where the LF parameters stabilize at those found in rich clusters. The effect is seen in the total LF and for the blue and red galaxies separately. The behaviour of the quiescent and star-forming samples is qualitatively consistent with variations resulting from interactions and mergers, where mergers build the bright end of the luminosity function at the same time as dwarf irregulars have their star formation quenched and evolve into dwarf ellipticals. These processes appear to take place preferentially in low luminosity groups and to be complete at a group luminosity of $-22.5$ in $B$, corresponding to a halo mass of order $10^{13.5}\\ M_{\\odot}$. ",
        "introduction": "There is considerable evidence that most galaxy properties, such as morphology \\citep{dressler80,postman84,allington93}, star formation rates \\citep{lewis02,gomez03, balogh04}, color \\citep{depropris04, kauffmann04} and even luminosity function parameters (Driver, Couch \\& Phillipps 1998; \\citealt{phillipps98, croton05}) correlate well with local density. It is likely that the main mechanism behind the observed correlations resides in rapid cessation of star formation in galaxies, once a threshold local density is reached (e.g. \\citealt{balogh04}). Observations of the color-magnitude relation in the neighborhood of clusters also point to a swift change in galaxy colors once the density in the surrounding structure reaches a critical value \\citep{kodama01,kodama05,tanaka05}. The environment where the transition in galaxy properties occurs, from those of typical field regions to those seen in dense clusters, is an important indicator of the evolutionary processes involved (see e.g. \\citealt{pimbblet06}, and references therein).  It is generally believed that most galaxy transformations (in morphology, star formation rates and color) take place in groups, since these are the environments where most galaxies reside and groups have sufficiently high densities, but sufficiently low velocity dispersions, that interactions and mergers may take place (e.g. \\citealt{ghigna98,mulchaey98}). Indeed, compact groups are seen to populate the infall regions of the rich clusters studied by \\cite{lewis02} and \\cite{tanaka05}, where the suppression of star formation and galaxy evolution appear to be taking place. Other processes, such as ram pressure stripping are expected to be important only at higher densities (e.g. Abadi, Moore \\& Bower 1999).  \\cite{kauffmann04}, \\cite{depropris04} and \\cite{croton05} show that most of the galaxies forming stars at the present epoch have relatively low luminosities and consist mostly of dwarf spirals and irregulars, while giants are mostly quiescent and tend to reside in the denser environments \\citep{yang05a}. This suggests that a study of the behavior of dwarf (and intermediate luminosity) galaxies in the group environment, together with a comparison of star-forming and quiescent galaxies may provide useful insight into the mechanisms that transform field galaxies into typical cluster members.  One of the most versatile (if simplistic) tools for studies of galaxy population variations  is the galaxy luminosity function (LF). The characteristic magnitude $M^*$ provides a measure of the variation of giant galaxy luminosity or mass, while the faint end slope $\\alpha$ can be related to the properties of the dwarf population. LFs in different wavebands, or selected according to appropriate color cuts, may be used to understand the stellar populations and star formation histories of the galaxy population under study.  In this paper we consider the variation of the galaxy LF in a sample of nearby groups spanning a wide range of masses.  This topic was previously investigated by \\cite{christlein00}, using Las Campanas Redshift Survey data (though this survey is now known to suffer from significant surface brightness selection effects), and by \\cite{eke04a}, who looked at three, dynamically determined, mass bins. Here, we consider the entire population, split this into several group luminosity bins and consider quiescent (red) and star-forming (blue) galaxies separately, in order to explore the possible roles of mergers and star formation suppression. We mainly concentrate on the behavior of lower luminosity galaxies, i.e. on the faint end of the LF, as a function of group properties, as this is where we expect to see the strongest signature of evolutionary effects. Dwarfs are believed to be fragile systems, whose properties are easily affected by dense environments (e.g. \\citealt{moore96}).  We start by describing the selection of groups and the derivation of the LF parameters in Section 2, and in Section 3 we discuss the observed trends and present possible interpretations  of our findings. We adopt a cosmology with $\\Omega_M=0.3$ and $\\Omega_{ \\Lambda}= 0.7$ (though this has minimal impact on our results) and normalize distances to H$_0=100$ km s$^{-1}$ Mpc$^{-1}$ (i.e. all quoted magnitudes should be read as $M - 5 \\log h$, where $h = H_0/100$ km s$^{-1}$ Mpc$^{-1}$). ",
        "conclusions": "Figure 5 summarizes the variation in LF parameters with group luminosity for each of the samples we consider. As in \\cite{eke04b} we find two noticeable trends in the data as a whole, viz. that $M^*$ brightens and $\\alpha$ steepens with increasing group mass. With the finer resolution we adopt, we are able to see that, in fact, the LF parameter values appear to reach approximately constant levels, comparable to those seen in  rich clusters \\citep{depropris03}, for groups brighter than $M^G_{b_J} \\simeq -22.5$. This is similar to the behaviour reported by \\cite{dominguez02} and \\cite{martinez02} from their analysis of groups in the 2dFGRS 100K galaxy release (with a different grouping algorithm) and is presumably related to the observation that galaxy properties appear to change at some `threshold' density in the neighborhood of clusters \\citep{kodama01,lewis02,tanaka05}. Of course, once we reach these asymptotic values, the group luminosity which we use as a discriminator will directly mirror the richness of the group.  Consider each of these trends in turn. As noted above, the characteristic magnitude $M^*$ levels off (or may even go fainter again) for groups with magnitudes brighter than $M^G_{b_J} \\simeq -22$ ($\\simeq 10^{11} h^{-2} L_{\\odot}$). The value seen for these more luminous groups is similar to that seen in rich clusters, even though they start from only moderate mass groups; assuming $h=0.7$ and a corresponding average universal ${\\cal M}/L$ ratio of 300 in solar units (e.g. \\citealt{bahcall00}), they have group (halo) masses ${\\cal M}_h \\simeq 6 \\times 10^{13}$ to $3 \\times 10^{14} {\\cal M}_{\\odot}$. We obtain a similar value for the mass at the `turnover' point, ${\\cal M}_h \\simeq 10^{13.6}$, if we use Eke et al.'s (2004b) luminosity dependent mass-to-light ratios.  Among the less luminous groups, there is a clear faintening of $M^*$ as one goes from intermediate to low luminosity. The dimming is seen in each left hand panel of Figure 3, though it is less clear for the blue galaxies. Assuming a simple linear relation we find slopes in $M^*$ vs. $M^G_{b_J}$, calculated as bisector fits of $0.77 \\pm 0.09$ for the whole sample, $0.63 \\pm 0.12$ for the red galaxies and $0.28 \\pm 0.13$ for the blue galaxies, with goodness of fit coefficients $R^2$ of 0.983, 0.960 and 0.717, respectively.  One could consider that this is merely a selection effect due to the enforced absence of very luminous galaxies: clearly we can not have an $M_{b_J} = -21$ galaxy in an $M^G_{b_J} = -20$ group. However, we would argue that it is still an interesting and informative exercise to determine how the available baryonic mass in a low mass group {\\it does} organise itself (i.e. the `conditional' LF; see e.g. van den Bosch, Yang \\& Mo 2003). From Figure 1, it is clear that the mass in the small groups is still shared out in a Schechter function-like distribution, but with a steadily declining bright end characteristic luminosity as we go to smaller groups\\footnote{\\cite{degrijs06} present an interesting discussion of the related question of how the mass function of the stars in small star clusters is `filled-up' until the available mass is used}.  We have also tested the `null' hypothesis of selecting out groups with low total luminosities from sets of galaxies (mock groups) produced by statistically sampling a fixed LF (both the 2dFGRS cluster LF of \\citealt{depropris03} and the field LFs of \\citealt{croton05}). As expected, $M^*$ becomes fainter for the mock groups with low total luminosity (due to the obvious lack of bright galaxies). However, these mock groups then turn out to be overpopulated with fainter objects compared to the real groups, leading to a steepening, not flattening, of the faint end slope as we go to lower mass groups. This gives us further confidence that the changes in slope we observe are real physical effects. It is more difficult to disentagle whether the change in $M^*$ is due to the effect of sampling the LF or is a real physical effect. However, though for a given small multiplicity group it is unlikely that a bright galaxy will be selected from a given purely statistical representation of the LF, once we sum over many groups, as here, the original LF will be restored (i.e. by chance, a few groups {\\it will} have a bright member).  Note that \\cite{hansen05} also find a shallower slope and fainter $M^*$ for their SDSS `clusters' with a low number of bright, red members. These objects are selected differently from our groups (they do not use spectroscopic redshifts, as we do), but their Figure 10 shows the same `saturation' of LF parameters for more massive systems as we find in our study.  Indeed, the observed behaviour is in excellent agreement with that predicted by the halo occupation model of \\cite{yang05d}. At low group masses (which, as here, they use group luminosity to represent), they find that the luminosity of the brightest galaxies should increase as ${\\cal M}_h^{2/3}$, with a much slower increase (slope $\\simeq 1/4$) for more massive groups. Moreover, they predict that the turnover in slope should occur at group halo masses a few times $10^{13} M_{\\odot}$, again in good agreement with our interpretation above (see also \\citealt{vale04,cooray04}). \\cite{yang05d} argue that this change in behaviour occurs at the transition from efficient to inefficient cooling in different mass halos, and is the same feature required in semi-analytical modelling of galaxy formation in order to match the bright end of the overall galaxy LF (see e.g. \\citealt{benson03}). \\cite{eke05} find that the group mass-to-light ratios (in the $K$ band) as a function of group mass level off at this same point, while in the semi-analytic models of \\cite{delucia06} there is a trend for early type galaxies in halos above this mass to be older, redder and metal richer (as against bluer, younger and metal poorer in the lower mass halos).  As noted earlier, the changes appear more clearly and are larger in amplitude for the red galaxies: perhaps unsurprisingly, small groups are even more deficient in bright red galaxies than bright blue ones. Even so, the generally similar trends for quiescent and star forming galaxies imply that the brightening of $M^*$ in larger groups is primarily a result of increasing galaxy masses, rather than just a variation in star formation and therefore in ${\\cal M}/L$. A reasonable interpretation would then be that what we are seeing is a direct result of hierarchical growth of halos and the merging of galaxies within them (e.g. \\citealt{dubinski98}; see also \\citealt{vandokkum05,bell06}).  Now consider the LF slope parameter $\\alpha$. Unlike $M^*$, this is more descriptive of the moderate to low luminosity (dwarf) galaxy population (though we should note that we do not sample very deep into the dwarf regime, so will be insensitive to any turn-up, as seen in some rich clusters (e.g. \\citealt{driver94,depropris95}; Smith, Driver \\& Phillipps 1997) and X-ray bright groups \\citep{miles04}, below $M_B \\simeq -17$). The slope should then be less influenced by the enforced absence of giant galaxies in the smaller groups. Nevertheless we again see a trend with group luminosity, with steeper slopes for the larger groups (as in \\citealt{zabludoff00}), and again a levelling off to the values observed for clusters once we reach $M^G_{b_J} \\simeq -22$. Notice that even though $\\alpha$ and $M^*$ will be linked by the fitting process, we generally sample to at least 2 magnitudes below $M^*$, so $\\alpha$ should genuinely measure the faint end slope. Nevertheless, this parameter linkage may account for the unusually small value of $\\alpha$ in the lowest group luminosity bin, where we have the smallest available fitting range.  Recall that the low luminosity groups are not equivalent to the general field but are a more specific environment. It is therefore not necessary (and indeed it is not the case) that the faint end of the LF in such groups should match, or asymptotically tend to, that of the field.  Field, and possibly void, LFs tend to have $\\alpha \\sim -1.2$ in the 2dF and SDSS studies, rather than our shallower slope. However, the two environments are not comparable, because local densities are computed over much larger volumes than the typical field size. Attempting to accommodate all the LFs with a fixed $\\alpha=-1.1$ results in poor fits for the low luminosity groups, but even then we still see systematic changes in $M^*$ with group luminosity.  Again, this behaviour matches well with that prediced by \\cite{yang05d} from the halo occupancy statistics (see also \\citealt{cooray05}). They find flatter slopes ($\\alpha \\geq -1$) for halos below about $10^{13.5} M_{\\odot}$ and steeper ones above that, consistent with both the direction and the crossover point we see in our variation of $\\alpha$ with group luminosity.  The blue galaxies show a variation of $\\alpha$ from about $-1.1$ to $-1.5$ (or even steeper), as previously observed for many clusters compared to the field (e.g. \\citealt{smith97}). The red galaxies show an even greater variation, the lowest luminosity groups showing very clearly declining LFs ($\\alpha > -1$) at the faint end. This argues for a general shortage of dwarf early type galaxies in these low density environments, consistent with their status as the most highly clustered of all galaxy types \\citep{ferguson91}.  Now consider the physical interpretation of the variation in faint end slope. If we are assuming that the larger groups form by the hierarchical build up of smaller ones, then we might look to the same sort of merger processes that we invoked earlier in respect to the brightening of $M^*$ (c.f. \\citealt{gonzalez05}). As we would expect mergers to take place most easily between the largest galaxies (e.g. \\citealt{ makino97}), then the bright end of the LF should build up at the expense of intermediate luminosities (in any case, only such major mergers would substantially change the brightness of a large galaxy). This will then deplete the mid-section of the LF (c.f. \\citealt{miles04}) and therefore create a steeper faint end, by default. However, this can not be the whole picture, as the lowest luminosity groups have, in particular, rather few faint red galaxies to start with; the change in slope of the red galaxy LF is substantially greater than that of the total sample. We therefore need to create extra faint red galaxies, the obvious source being intermediate to low luminosity late type galaxies fading and reddening due to the cutting-off of their star formation \\citep{lin83,davies88,christlein03}. This then reinforces the idea of reduced star formation rates in increasingly dense environments \\citep{lewis02, depropris03,gomez03,depropris04,balogh04}.  The steepening of the LF for blue galaxies at around $M^G_{b_J} = -22.5$ is perhaps surprisingly abrupt, but it is reminiscent of earlier claims for very steep LFs in clusters of galaxies \\citep{driver94,depropris95}, especially in their outskirts \\citep{phillipps98,boyce01}. Conceivably this effect may be due to star formation continuing in fainter dwarfs when it has ceased in intermediate mass ones (mimicking the global `downsizing' picture of cosmic star formation; Cowie et al. 1996). Another intriguing possibility would be that the blue galaxies represent a new population of tidal dwarfs (\\citealt{mirabel92}; Hunsberger, Charlton and Zaritsky 1996), which tend to have steep luminosity functions (Hunsberger, Charlton \\& Zaritsky 1998), though they are not usually thought to be as bright as the objects in our LFs. For the largest groups there may be a modest flattening of the slope again, perhaps reflecting the smaller number of dwarfs seen in some very dense environments \\citep{boyce01}.  When attempting to interpret the observed trends, we should also consider the appropriateness of our measure of the group or cluster environment. A group's total mass is evidently a critical measure, as halo mass is the key variable in structure formation models (\\citealt{press74} {\\it et seq.}; \\citealt{yang05b}). The total luminosity has been demonstrated to be a good measure of total stellar mass and total halo mass \\citep{padilla04,yang05c}. One could argue that the total $R$ band luminosity would be preferable to the $B$ band luminosity, as it is less subject to contributions from short term star formation, and that the $K$ band would be better still; $K$ band LFs should also more closely reflect the mass distribution of the individual group galaxies (e.g. Cole et al. 2001). In fact, $B$ and $R$ group luminosities turn out to be well correlated so experiments dividing the sample by group $R$ magnitudes presents the same picture as that already discussed, while $K$ band data is not yet available for a large enough sample to overcome statistical uncertainties (though see \\citealt{eke05} for a current analysis).  We could also divide the groups purely by multiplicity, though this will bring in additional (or at least different) selection effects, and theoretically is expected to have wide, poissonian, variations for a given halo mass (e.g. \\citealt{kravtsov04,yang05d}), or by a dynamical measure such as velocity dispersion \\citep{christlein00,eke04b} or parameters such as X-ray flux \\citep{miles04}. We defer a discussion of this and of a wider range of group properties to a subsequent paper.  To sum up, the results of our study suggest that changes in the galaxy population take place mostly among the lower luminosity groups and are consistent with the changes being brought about by merger/interaction-induced accretion and star formation suppression. We have argued that the observed changes in LF parameters are not merely a selection effect due to the total available luminosity.  The processes appear to be largely complete once the group luminosity reaches a threshold value $M^G_{b_J} \\simeq -22.5$ (with projected densities of a few giant galaxies per Mpc$^2$).  This is similar to earlier observations in groups and clusters \\citep{kodama01,dominguez02,martinez02,kodama05,sabatini05,tanaka05}. The `threshold' luminosity may arise during the growth of groups, when the internal velocity dispersion becomes so large as to inhibit further interactions despite the increased galaxy density. If the more massive groups then coalesce to produce clusters (as in the hierarchical picture), the remarkable homogeneity in the properties of cluster galaxies observed by \\cite{christlein03} and \\cite{depropris03} would result naturally."
    },
    "0608/gr-qc0608060_arXiv.txt": {
        "abstract": " ",
        "introduction": "Kinetic theory in general relativity is treated by a number of authors, and has been reviewed in much more detail by Ehlers$^{(1)}$, using the differential calculus on Riemannian manifolds. Recently the growing interest to the general relativistic kinetic theory in plasma appeared$^{(2)-(4)}$. However, the kinetic theory of electrical conductivity in conductors in the presence of gravitational field is not strongly discussed in the literature and its presentation may be of some interest (The formulation of the kinetic theory of electrical conductivity in the astrophysical matter in the flat spacetime can be found, for example, in$^{(5)}$). In this paper we consider the general relativistic kinetic theory needed to determine on microscopic level the influence of a stationary gravitational field on electromagnetic transport phenomena for normal conductors. This approach is interesting, even from purely theoretical point of view, because it involves the interplay of gravitation, electromagnetism, relativity and statistical physics.  The paper is organized in the following way. First, in $\\S2$ we consider Boltzmann's kinetic equation in a manifestly covariant form and the distribution function for electron gas in the local equilibrium. Then in the approximation of relaxation time we derive the general-relativistic nonequilibrium distribution function for electron gas in an applied electromagnetic field which is the key element in the construction of microscopic theory of conductivity: the general relativistic Ohm's law is defined by the corresponding average in velocity space. In $\\S3$ it is shown that the theoretical project$^{(6)}$ on measuring gyroscopic effect in two series connected coil conductor carrying current and rotating in opposite directions has to give the null result. It appears author of the paper$^{(6)}$ ignored the surface charges which compensate rotationally induced space charges.  Finally, we apply the derived Ohm's law to a few devices that can be used as detectors of gravitomagnetism and rotation. ",
        "conclusions": "We conclude:  \\noindent 1. The solution of the general-relativistic Boltzmann's kinetic equations for the nonequilibrium distribution function for electron gas leads to the Ohm's law for conduction current being valid in the fields of gravity and inertia. The latter has pure general-relativistic contributions arising from the absolute acceleration and relativistic rate of rotation.  \\noindent 2. `Gyroscopic' effect in the sense of the paper$^{(6)}$ does not exist since rotationally induced space charges in the solenoid with electric current are totally compensated by the surface charges generated at the inner $r=r_1$ and the outer $r=r_2$ boundaries of the solenoid.  \\noindent 3. Experiment is proposed to measure radial voltage through the rotating solenoid carrying azimuthal electric current.  \\noindent 4. Finally, the nonstationary situation, when alternate azimuthal current produces radial current across the conductor due to the effect of the Earth's Lense-Thirring field has been considered."
    },
    "0608/astro-ph0608357_arXiv.txt": {
        "abstract": "We present a study of the natal core harboring the class 0 protostar \\gf\\ (bolometric luminosity $\\simeq 0.3$ \\Lsun, bolometric temperature $\\lesssim 20$ K; Wiesmeyer et al.) in the filamentary dark cloud GF\\,9 (distance $=$ 200 pc) using the Nobeyama 45\\,m, CSO 10.4\\,m telescopes, and the OVRO mm-array. \\gf\\ stands unique in the sense that it shows \\wat\\ maser emission (Furuya et al.), a clear signpost of protostar formation, whereas it does not have a high-velocity large-scale molecular outflow evidenced by our deep search for \\co\\ wing emission. These facts indicate that \\gf\\ core is early enough after star formation so that it still retains some information of initial conditions for collapse. Our \\um\\ dust continuum emission image revealed the presence of a protostellar envelope with an extent of $\\simeq$ 5400 AU in the center of a molecular core $\\simeq$ 0.08 pc in size. The mass of the envelope is $\\simeq 0.6$ \\Msun\\ from the \\um\\ flux density, while LTE mass of the core is $\\simeq 3$ \\Msun\\ from \\NtwoH, \\HtCOp, CCS, and \\amm\\ line observations. Combining visibility data from the OVRO mm-array and the 45\\,m telescope, we found that the core has a radial density profile of $\\rho(r)\\propto r^{-2}$ for $0.003\\lesssim r/{\\rm pc}\\lesssim 0.08$ region. Molecular line data analysis revealed that the velocity width of the core gas increases inward, while the outermost region maintains a velocity dispersion of a few times of the ambient sound speed. The broadened velocity width can be interpreted as infall. Thus, the collapse in \\gf\\ is likely to be described by an extension of the Larson-Penston solution for the period after formation of a central star. We derived the current mass accretion rate of $\\simeq 3\\times 10^{-5}$ \\Msun\\ year$^{-1}$ from infall velocity of $\\simeq$ 0.3 \\kms\\ at $r\\simeq$ 7000 AU. Furthermore, we found evidence that a protobinary is being formed at the core center. All results suggest that \\gf\\ core has been undergoing gravitational collapse for $\\lesssim$ 5000 years since the formation of central protostar(s), and that the unstable state initiated the collapse $\\simeq 2\\times 10^5$ years (the free-fall time) ago. ",
        "introduction": "Stars form inside dense ($\\simeq 10^5$ \\cmc) and compact ($\\simeq 0.1$ pc) molecular cores (e.g., Myers \\& Benson 1983; Mizuno et al. 1994). One of the roles of such cores is to provide material for newly forming stars through dynamical accretion caused by gravitational collapse of the cores. For isolated low-mass star formation, two extreme paradigms for collapse of a core have been proposed. One is a slow, quasi-static model, while the other extreme is the dynamic, turbulent one. The former is represented by the Shu's similarity solution (Shu 1977), and the latter by the Larson-Penston's solution (Larson 1969; Penston 1969) which was extended for the period after a protostar formation by Hunter (1977) and Whitworth \\& Summers (1985). Hereafter, we refer to the latter extended solution as the Larson-Penston's solution. The two solutions are discriminated by radial density and velocity profiles of the core at $t =$ 0. The Shu's solution (1977) sets the density profile of $\\rho(r)=\\frac{c_{\\rm s}^2}{2\\pi G}~r^{-2}$ and is static ($v(r) =0$) at $t=0$ where $c_{\\rm s}$ and $G$ are the isothermal sound speed and the gravitational constant, respectively. At $t = 0$, the collapse of the core begins at the center, which produces an expansion or rarefaction wave that propagates outward with the local sound speed. This model is referred to as the ``inside-out'' collapse (Shu 1977; Shu, Adams \\& Lizano 1987). On the other hand, the Larson-Penston's solution approaches the profile of $\\rho(r)\\simeq 4.4~ \\frac{c_{\\rm s}^2}{2\\pi G}~r^{-2}$ and in infalling state ($v(r) \\approx -3.28~c_{\\rm s}$) when a protostar is about to be born ($t\\rightarrow -0$, Hunter 1977; Shu 1977; Foster \\& Chevalier 1993). Thus the Larson-Penston solution, referred to as the ``runaway'' collapse, predicts a 48 times higher mass infall rate than the Shu's. At $t >$ 0, the both solutions have a free-fall region of \\RadpInf\\ which expands with sound speed for the Shu solution and with supersonic velocity for the Larson-Penston that accompanies no rarefaction wave.\\par      Molecular line and dust continuum observations have identified a number of gravitationally bound starless cores (e.g., Mizuno et al. 1994), referred to as ``preprotostellar'' cores (Ward-Thompson et al. 1994). Thermal dust continuum emission imaging at mm and submm wavelengths revealed that radial density profiles of the preprotostellar cores have much flatter profiles than $\\rho (r)\\propto r^{-2}$ at radii smaller than a few thousand AU (e.g., Ward-Thompson et al. 1994; Andr\\'e et al. 1996). Molecular line studies measuring velocity field of the cores are also powerful tool to verify collapse models (e.g., Zhou 1992). For this purpose, it is the most important to search for a star-forming core which retains the initial conditions for the gravitational collapse, namely, a star-forming core at an extremely early evolutionary phase before dispersal by a large-scale molecular outflow. Furuya et al (1999, 2000) proposed that the intermediate-mass class 0 source S106\\,FIR is such an ideal core, because the source is characterized by (1) a cold SED (Richer et al. 1993), (2) \\wat\\ maser emission, i.e., strong evidence for the presence of a protostar, and (3) absence of a large-scale high-velocity CO outflow.\\par    We have successfully identified a low-mass counterpart, \\gf\\ by performing multi-epoch \\wat\\ maser survey (Furuya et al. 2001, 2003) and \\co\\ outflow mapping survey toward the selected low-mass protostars with the Nobeyama 45\\,m, SEST 15\\,m, and CSO 10.4\\,m telescopes (R. S. Furuya et al. in preparation). In our sample, \\gf\\ is the only protostar satisfying the above characteristics of (1)--(3). This protostar is located in the east-west filamentary dark cloud GF\\,9 where several YSOs are located every $\\simeq 0.75$ pc in the plane of the sky (Schneider \\& Elmegreen 1979; Wiesemeyer 1997; the filament is cross-identified as Barnard\\,150 and Lynds\\,1082. The distance to \\gf\\ is estimated to be 200 pc by Wiesemeyer (1997). Wiesemeyer et al. (1997, 1999) identified \\gf\\ as class 0 and obtained bolometric temperature (\\Tbol) of $\\leq$ 20 K and bolometric luminosity (\\Lbol) of $0.3 L_{\\odot}$. We believe that \\gf\\ offers unique opportunity to investigate the initial conditions for gravitational collapse of a star-forming core. This paper presents a detailed observational study of \\gf. In particular, we combined mm-interferometer and single-dish telescope visibility data to obtain high fidelity images of the core. ",
        "conclusions": "\\label{s:Discussion} \\subsection{Evolutionary Stage of the GF\\,9-2 Core: A Very Young Protostar Before Formation of a Large-Scale Molecular Outflow} \\label{ss:evl_flow}  With its low luminosity ($\\simeq 0.3$ \\Lsun) and cold SED (\\Tbol\\ $\\leq 20$ K; Wiesemeyer et al. 1997), we believe that the protostar embedded in \\gf\\ core is at an extremely early phase of low-mass star formation, evidenced by the absence of a large-scale high-velocity molecular outflow, the possible detection of a compact low-velocity outflow, and the detection of \\wat\\ masers. We estimate the protostar's age from dynamical age (\\tdyn) of the possible compact low-velocity outflow ($\\S\\ref{ss:res_outflows}$). The terminal velocity ($V_{\\rm t}$) of the spectra would give the line-of-sight component of the outflow velocity of $|V_{\\rm t}-V_{\\rm sys}|=V_{\\rm flow}\\cos i$ where $i$ is the inclination angle of the outflow axis to the line-of-sight. Given the projected lobe length of $l_{\\rm lobe}\\sin i = 1,1000$ and 6600 AU and $V_{\\rm flow}\\cos i \\simeq$ 3.3 and 3.1 \\kms\\ ($\\S$\\ref{ss:res_outflows}), we have $t_{\\rm dyn}=l_{\\rm lobe}/V_{\\rm flow}\\simeq 1.6\\times 10^4$ and $1.0\\times 10^4$ years with $i=$ 45\\degr\\ for the blue- and redshifted \\co\\ (3--2) lobes, respectively. These dynamical ages are comparable to those for the extremely young class 0 sources, e.g., (0.6--3)$\\times 10^4$ years for VLA\\,1623 and (0.8--2)$\\times 10^4$ years for IRAM\\,04191 (Table 1 in Andr\\'e, Motte \\& Backman 1999, and references therein).\\par      The youth of \\gf\\ is demonstrated in Figure \\ref{fig:gfevl} where we present a plot of CO outflow momentum rate (\\Fco) normalized by the bolometric luminosity (\\FcocLbol) vs. \\Tbol\\ for nearby ($d\\leq 350$ pc) low-mass (\\Lbol\\ $\\leq 50$ \\Lsun) YSOs. We employ \\FcocLbol\\ because molecular outflows in low-mass protostars are likely to be momentum-driven (e.g., Cabrit \\& Bertout 1992; Richer et al. 2000, and references therein). The temperature \\Tbol\\ is believed to trace the evolutionary stage of YSOs (Myers \\& Ladd 1993). We calculate 3$\\sigma$ upper limit of \\FcocLbol\\ $= \\frac{M_{\\rm lobe}V_{\\rm flow}^2}{l_{\\rm lobe}}\\cdot \\frac{c}{L_{\\rm bol}}$ $\\simeq$ 9 and 20 for the \\co\\ (3--2) blue- and red lobes, respectively, with $i=$ 45\\degr. Figure \\ref{fig:gfevl} illustrates that ``class 0'' is the most powerful outflow phase and ``class I'' is the decline phase (e.g., Moriarty-Schieven et al. 1994). The most remarkable result is that \\gf\\ shows the lowest outflow activity, among the coldest low-mass class 0 sources, although there exists enough circumstellar material ($\\S$\\ref{ss:Mlte}) to be dragged by a possible protostellar jet to develop a large-scale molecular outflow. Note that the prototypical class 0 sources of VLA\\,1623 (AWB93) and IRAM\\,04191 (Andr\\'e et al. 1999) have powerful CO outflows, despite that their bolometric luminosities are comparable to that of \\gf: 1 \\Lsun\\ for VLA\\,1623 and 0.15 \\Lsun\\ for IRAM\\,04191. We thus believe that the low luminosity of \\gf\\ is not responsible for the lowest outflow activity.  Consequently, it is very likely that the protostar of our interest is the youngest class 0 protostar ever identified and is just at the beginning of the outflow launch.\\par       \\subsection{The Initial Conditions for Gravitational Collapse of a Low-Mass Star Forming Core} \\label{ss:ICofSF}   \\subsubsection{Stability of the \\gf\\ Core} \\label{ss:stab}  If we interpret that the non-thermal velocity width obtained in \\gf\\ core ($\\S$\\ref{ss:widthmap}) is attributed to turbulent motions, we expect that the turbulence can help support the core against its self-gravity. To test this hypothesis, we compare the LTE mass of the core with virial mass \\Mvir\\ including the turbulence. Assuming a spherical system, the virial mass is given by $M_{\\rm vir}\\simeq R_{\\rm eff}\\langle\\Delta v_{\\rm int}\\rangle^2/G$ where \\meandVint\\ includes the non-thermal width, i.e., the turbulence. Given \\meandVint\\ and \\Reff\\ values (Tables \\ref{tbl:ResultsLine} and \\ref{tbl:Mcore}), we found that \\Mvir\\ ranges from 0.39 to 2.9 \\Msun (case I in Table \\ref{tbl:Mcore}). Since \\Mvir\\ $\\simeq$ \\MLTE\\ within the uncertainties, the core is most likely to be in virial equilibrium. However, it should be noted that this hypothesis does not reconcile with the line broadening toward the core center (Figure \\ref{fig:dVint}). Moreover, as far as we know, no cores which are supported by turbulence exhibit a \\RadpSimSoltn\\ density profiles (see Figure \\ref{fig:radprofile}). Therefore, we conclude that the non-thermal line width is not caused by turbulence.\\par     Alternatively, if we consider that infall causes the non-thermal line width, that is, if \\gf\\ core is gravitationally unstable, the drawbacks in the previous hypothesis would be solved. The presence of infalling motion is supported by the following facts:~ (1) the presence of the blueshifted gas and the less prominent redshifted counterpart in the central $r\\lesssim$ 1000 AU ($\\S$\\ref{ss:pv}), (2) the spectroscopic signature of infall, i.e., the ``blue-skewed'' profile observed in the optically thick \\isolated\\ line (Figure \\ref{fig:spectra}), and (3) the detection of the \\wat\\ masers, a clear signpost of star formation, naturally leading to the presence of accretion flow onto the star-disk system. We, however, point out a drawback in the above (2); the optically thin \\HtCOp\\ emission does not show a single Gaussian profile expected from the infall interpretation, while the CCS does. In this interpretation, \\Mvir\\ should be estimated only from the thermal motion as \\Mvir\\ $=\\frac{R_{\\rm eff}}{G}\\frac{kT_{\\rm k}}{\\mu m_{\\rm H}}$. The estimated \\Mvir\\ from the 4 lines takes a range between 0.078 and 0.25 \\Msun\\ (case II in Table \\ref{tbl:Mcore}) which is one order of magnitude smaller than \\MLTE. The small \\Mvir\\ is consistent with our interpretation that the core is under gravitational collapse. As discussed in $\\S\\ref{ss:widthmap}$, the core would be rotating with $\\Lambda\\simeq 10^{-14}$ s$^{-1}$. The collapse cannot be supposed by the centrifugal force due to the rotation because the rotational energy is less than a few percent of total energy required to support the core against its self-gravity. For such an unstable core, the gas is expected to be more perturbed toward the center, and the velocity width should increase inward as a result. In fact, Figure \\ref{fig:dVint} clearly indicates that \\dVint\\ in \\gf\\ increases up to $\\simeq 0.6$ \\kms\\ toward the center.\\par        \\subsubsection{Initial Conditions for the Gravitational Collapse of the Core} \\label{ss:IC}  To bring out the initial conditions for the gravitational collapse in \\gf\\ core, we examine the density and velocity structures by comparing with some key characteristics of the Shu (1977) and Larson-Penston (Larson 1969; Penston 1969) solutions. The radial H$_2$ column density profile of the core is given by $(7.7\\pm 4.1)\\times 10^{23}\\left(\\frac{{\\rm X(N_2H^+)}}{9.5\\times 10^{-10}}\\right )^{-1} \\left(\\frac{r}{200 \\rm AU}\\right )^{-(1.0\\pm 0.1)}$ \\cmq\\ for 0.001 $< \\frac{r}{\\rm pc} <$ 0.05 from the combined data and $(8.9\\pm 4.5)\\times 10^{23}\\left(\\frac{{\\rm X(N_2H^+)}}{9.5\\times 10^{-10}}\\right )^{-1} \\left(\\frac{r}{200 \\rm AU}\\right )^{-(1.1\\pm 0.2)}$ \\cmq\\ for 0.005 $< \\frac{r}{\\rm pc} <$ 0.08 from the 45\\,m data (Figure \\ref{fig:radprofile}). It should be noted that the dense \\NtwoH\\ gas seem to obey the power-law profile with the index of $-1$ even at the innermost region of $r\\lesssim$ 1000 AU where a possible compact outflow might begin to destroy the core. The power-law indices of $-1.0$ and $-1.1$, converted into $-2.0$ and $-2.1$ for the volume density, well agree with those in the two solutions after the formation of the protostar (Shu 1977; Whitworth \\& Summers 1985). The best-fit H$_2$ column densities at $r =$ 200 AU fall between the two extreme predictions at $T =$ 9.5 K: $(2.8\\pm 1.5)\\times 10^{23}$ \\cmq\\ for the Shu and $(13\\pm 7)\\times 10^{23}$ \\cmq\\ for the Larson-Penston solutions. Here we used a relation of $N_0({\\rm H_2}) =$ 1.0 and $4.4\\times\\frac{c_{\\rm s}^2}{2Gr_0}\\cdot\\frac{{\\rm X(N_2H^+)}}{\\mu m_{\\rm H}}$ for the Shu's and Larson-Penston's cases, respectively (Hunter 1977; Shu 1977; Foster \\& Chevalier 1993). Because of the abundance uncertainty, however, we cannot conclude which solution describes well the collapse in \\gf\\ core.      Contrary to the above discussion based on the density profile, the velocity width maps (Figure \\ref{fig:dVint}) help to discriminate the two theoretical scenarios for the collapse of \\gf\\ core. Notice that the influence of the compact outflow on the velocity width cannot be discerned toward the 3\\,mm source. Figure \\ref{fig:dVint} shows that the velocity width decreases outward, but remains $\\simeq$ 0.3 \\kms, which is twice the ambient \\Cs, even in the outermost region of $r\\simeq 50\\arcsec$ (corresponding to 0.05 pc) where the density profile shows \\RadpSimSoltn\\ (Figure \\ref{fig:radprofile}). Recall that \\RadpSimSoltn\\ profile characterizes outer part of a collapsing core. These facts clearly indicate that the extended Larson-Penston solution for $t>$ 0 is more likely to be the case because it requires the infalling state when a protostar is about to be born ($v(r) \\sim -3.28 c_{\\rm s}$ at $t \\rightarrow -0$; Hunter 1977), whereas the Shu solution requires a static state ($v(r) =$ 0 at $t =$ 0). It is, therefore, highly probable that the gravitational collapse of \\gf\\ core is well described by the runaway collapse scenario.\\par     Subsequently, we estimate the age of the \\gf\\ protostar (\\tpstar) based on the lack of the density profile of \\RadpInf\\ down to the innermost region of $r\\simeq 3\\arcsec$ (see Figure \\ref{fig:radprofile}). Recall that \\RadpInf\\ profile is expected for a central free-fall region which appears in both of the solutions at $t >$ 0. The non-detection of a free-fall density profile indicates that the central free-fall region, propagating outward with supersonic velocity, has not yet developed up to a radius ($r_{\\rm ff}$) of $\\simeq$ 600 AU, suggestive of the youth of the \\gf\\ protostar. However, we do not exclude the possibility that the absence of a detection of the \\RadpInf\\ profile at the center could be caused by the formation of binary (described in $\\S\\ref{ss:pbinary}$) whose projected separation is $\\simeq$ 1200 AU. Using Eqs. (9) and (10) in Ogino, Tomisaka \\& Nakamura (1999), we estimate \\tpstar\\ to be less than \\tpstar\\ $\\simeq 5000\\left(\\frac{T}{\\rm 9.5~K}\\right)^{-\\frac{1}{2}}$ years since the formation of the central protostar for the Larson-Penston solution. Although the age estimate is a factor of 2--3 smaller than the outflow dynamical age of \\tdyn\\ $\\simeq (1.0-1.6)\\times 10^4$ years ($\\S$\\ref{ss:evl_flow}), we argue that the central protostar has just begun to exhibit its jet activity traced by the \\wat\\ masers ($\\S$\\ref{ss:evl_flow}).\\par    Consequently, we propose that \\gf\\ core has started its gravitational collapse with initially unstable state. Here, the term {\\it initial} means $t= -t_{\\rm ff}$ where \\tff\\ is free-fall time of the core (Larson 1969; Penston 1969). For \\gf, we calculate \\tff\\ $\\simeq 2\\times 10^{5}$ years. We therefore believe that the age of the core, \\tcore\\ $=$ \\tff\\ $+$ \\tpstar, would be $\\simeq 2\\times 10^{5}$ years, which has a consistency with the time scale of chemical evolution ($\\S$\\ref{ss:Xccsamm}).\\par     \\subsubsection{Infall Velocity and Mass Accretion Rate in the Core} \\label{ss:Mdot}  Assuming that \\dVint\\ represents a velocity difference of blue- and redshifted infalling gas along the line-of-sight, let us discuss infall velocity (\\Vinf) of the core gas. We found that \\Vinf\\ varies from 0.2 to 0.4 \\kms\\ using a relation of \\Vinf\\ $\\simeq \\frac{1}{2}|v_{\\rm inf, red}-v_{\\rm inf, blue}|/\\cos i\\simeq$ $\\frac{\\dVint}{2\\cos i}$ for the \\dVint\\ range of (0.28 -- 0.63) \\kms\\ over the core (see Figure \\ref{fig:dVcenter}) and $i=$ 45\\degr. We now calculate the mass infall rate from $\\frac{dM}{dt} = 4\\pi R_{\\rm eff}^2 \\mu m_{\\rm H}\\cdot n_{\\rm H_2}(R_{\\rm eff})\\cdot v_{\\rm inf}$ $\\simeq 2\\pi R_{\\rm eff}\\cdot \\frac{\\mu m_{\\rm H}\\cdot N_{\\rm N_2H^+}(R_{\\rm eff})}{{\\rm X(N_2H^+})}\\cdot v_{\\rm inf}$ for a sphere with the radius of \\Reff. Here $n_{\\rm H_2}(R_{\\rm eff})$ is a volume density of a thin shell whose radius is \\Reff\\ and is approximated as $\\frac{N_{\\rm N_2H^+}(R_{\\rm eff})}{2R_{\\rm eff}\\cdot{\\rm X(N_2H^+)}}$. Given \\Vinf\\ $=$ 0.29 \\kms\\ at $r \\simeq$ \\Reff\\ from the \\NtwoH\\ \\dVint\\ map (Figure \\ref{fig:dVint}), we have $\\frac{dM}{dt}\\simeq 3\\times 10^{-5}$ \\Msun\\ year$^{-1}$ with uncertainty of a factor 1.6 for \\Reff\\ in Table \\ref{tbl:Mcore}, the column density profile ($\\S\\ref{ss:resdensity}$), and X$_{\\rm m}$ in Table \\ref{tbl:Xabs}. The inferred mass accretion rate falls within the range between the predicted values of $1.5\\times 10^{-6}$ and $ 7.4\\times 10^{-5}$ \\Msun\\ year$^{-1}$ for the Shu and Larson-Penston solutions, respectively, at 10 K. These accretion rates are calculated by $m_0 c_{\\rm iso}^3/G$ where $m_0$ is the non-dimensional factor of 0.975 for the Shu and 46.92 for the Larson-Penston solutions (Hunter 1977). Our small mass accretion rate, compared to that of the Larson-Penston's solution, is attributed to both the smaller column density and infall velocity.      \\subsection{Is a Protobinary System Forming ?} \\label{ss:pbinary}  We found that the position of 3\\,mm continuum source is displaced by $ \\simeq 6\\arcsec$ to the east of the \\NtwoH\\ column density peak, which would be explained by the presence of a protobinary system. In this interpretation, the 3\\,mm source represents a protostar and the \\NtwoH\\ peak a preprotostellar condensation. This is because the \\NtwoH\\ condensation is deeply embedded in the extended \\um\\ emission (Figure \\ref{fig:Nmaps}b), but is not accompanied with a compact 3\\,mm source. To investigate the nature of the binary members, we made a \\dVint\\ map from the combined \\HtCOp\\ data shown in Figure \\ref{fig:dVcenter} as the line is expected to be optically thin. In $\\S\\ref{ss:rescont_dust}$, we asserted that the 3\\,mm continuum emission would represent circumstellar material around a protostar. The line broadening toward the 3\\,mm source seen in Figures \\ref{fig:pv} and \\ref{fig:dVcenter} suggests the local enhancement of mass accretion, leading to the existence of a protostar. On the other hand, the \\NtwoH\\ peak (white contour) seems to be a gas condensation or a preprotostellar core rather than an apparent peak caused by local abundance enhancement because X(\\NtwoH)$_{8\\arcsec}$ inferred over the 3$\\sigma$ level contour of \\um\\ emission is comparable to those obtained with low spatial-resolution studies ($\\S\\ref{ss:Xnh}$). It should be noted that the \\HtCOp\\ emission has a secondary peak near the \\NtwoH\\ condensation as seen in the OVRO maps (Figure \\ref{fig:hikaku}). Furthermore, the combined \\NtwoH\\ map (Figure \\ref{fig:hikaku}) does not show a shell-like structure which is expected to be seen if \\NtwoH\\ disappears from the gas phase at the core center (e.g., Kuiper, Velusamy \\& Langer 1996; Belloche \\& Andr\\'e 2004). We therefore conclude that the positional displacement of \\NtwoH\\ peak is not caused by chemical evolution, but the \\NtwoH\\ peak represents a preprotostellar condensation, making it younger than the eastern object. The detection of the two objects reinforces the assertion that we witness a forming binary system with a separation of $\\simeq$ 1200 AU.    The masses of the circumstellar material ($M_{\\rm cm}$) for each binary member would be an order of 0.1 \\Msun. For the eastern object, the 3\\,mm continuum flux density ($\\S$\\ref{ss:rescont_dust}) integrated inside its 3$\\sigma$ level contour (see Figure \\ref{fig:dVcenter}) implies that the gas plus dust mass ranges from $\\simeq$ 0.1 (case I in $\\S\\ref{ss:rescont_dust}$) to 0.2 \\Msun\\ (case II). For the western object, we estimate (0.1 -- 0.3) \\Msun\\ from the \\NtwoH\\ intensity integrated inside the white contour, whose extent is similar to that of the 3\\,mm source (Figure \\ref{fig:dVcenter}). The inferred $M_{\\rm cm}$ seem reasonable for an amount of circumstellar material with 100 AU scale around a newly formed protostar at the very early accretion phase because of $M_{\\rm cm} \\simeq M_{\\rm protostar} = \\frac{dM}{dt}\\times t_{\\rm protostar} \\simeq 0.15$ \\Msun\\ ($\\S$\\ref{ss:IC} and \\ref{ss:Mdot}). In this context, the EW elongated structure at the core center seen in the 45\\,m maps (Figure \\ref{fig:hikaku}) may be interpreted as a 1000 AU scale circumbinary disk-like structure because the size would be smaller than the extent of the protostellar envelope traced by the \\um\\ continuum emission (6600$\\times$4500 AU; $\\S$\\ref{ss:rescont_dust}). It is impossible to uniquely define spatial extent of the disk-like structure. However, if we arbitrarily define the disk-like structure by a region enclosed by the 18 or $21\\sigma$ level contours of the combined \\HtCOp\\ map (Figure \\ref{fig:dVcenter}), mass of the circumbinary disk-like structure (\\Mcbd) calculated from the \\NtwoH\\ data would range (0.2 -- 0.5) \\Msun\\ with an approximate effective radius of $\\simeq$ 2000 AU. Considering the uncertainties in spatial extents and in $\\kappa_0$, $\\beta$ ($\\S$\\ref{ss:rescont_dust}) and X(\\NtwoH) ($\\S$\\ref{ss:Xnh}), the mass of disk-like structure is likely to be plausible to be hosted in the (0.4 -- 0.8) \\Msun\\ envelope ($\\S$\\ref{ss:rescont_dust}) and to host the two binary members with the 0.1 \\Msun\\ order circumstellar material. The mass distribution of $M_{\\rm core} > M_{\\rm env} \\gtrsim M_{\\rm cbd} \\gtrsim M_{\\rm cm}$ does not contradict with our hypothesis that the central protostar(s) is at extremely early evolutionary phase.\\par    If a circumbinary disk exists, the radius of disk-like structure would be physically related to the rotation of the natal core, described in $\\S\\ref{ss:pv}$. Although the axis for the large-scale rotation of the core (P.A.$\\simeq$ 135\\degr) is not perpendicular to the major axis of the disk-like structure (P.A.$\\simeq$ 90\\degr), a rough estimate of radius of the disk-like structure would be given by the centrifugal radius, \\Rcent, calculated from the specific angular moment of the core gas. With a relation of $R_{\\rm cent}=(4R_{\\rm eff}^4\\Lambda^2)/(GM_{\\rm cbd})$, we obtained \\Rcent\\ $\\simeq 900$ AU for $\\Lambda$ of $3\\times 10^{-14}$ s$^{-1}$ ($\\S\\ref{ss:pv}$) and \\Reff\\ in Table \\ref{tbl:Mcore}. The centrifugal radius is roughly twice the radius of disk-like structure, probably supporting our interpretation. Lastly, let us estimate the Toomre's $Q$ parameter to examine whether or not gravitational fragmentation of the disk-like structure led to the formation of the binary. The $Q$ parameter is given by $\\frac{\\Lambda c_{\\rm s}}{\\pi G \\Sigma}$ for Keplerian rotation. Here $\\Sigma$ is the surface density of the disk-like structure given by $M_{\\rm cbd}/{\\cal A}_{\\rm cbd}$; an area of ${\\cal A}_{\\rm cbd}$ $\\simeq 4\\times 10^{33}$ cm$^2$ is calculated for the region described above. We obtained $Q\\simeq 0.01 \\left (\\frac{\\Lambda}{3\\times 10^{-14} \\rm ~s^{-1}}\\right ) \\left (\\frac{T}{\\rm 10~ K}\\right )^{\\frac{1}{2}} \\left (\\frac{M_{\\rm cbd}}{0.4 M_{\\odot}}\\right )^{-1} \\left (\\frac{{\\cal A}_{\\rm cbd}}{4\\times 10^{33}~ \\rm cm^2}\\right ) $. Considering the uncertainties in the input parameters, it is certain that the $Q$ value is less than unity. We, therefore, suggest that the binary system is likely to be forming in the core center due to the gravitational fragmentation of the circumbinary disk-like structure, whose radius appears to be determined by the rotation of the core."
    },
    "0608/astro-ph0608482_arXiv.txt": {
        "abstract": "Since the advent of SCUBA on the JCMT, a series of complementary surveys has resolved the bulk of the far-infrared extragalactic background into discrete sources. This has revealed a population of heavily dust-obscured sources at high redshift ($z>1$) undergoing an intense period of massive star-forming activity with inferred star formation rates of several hundred to several thousand solar masses per year. Taken together, these existing surveys cover a total area of 460 sq. arcmin to a range of depths, but combining the results has hitherto been complicated by the fact that different survey groups have used different methods of data reduction and source extraction. In this paper we re-reduce and analyse all of the blank field surveys to date in an almost identical manner to that employed in the ``SCUBA 8\\,mJy Survey''. Comparative source catalogues are given which include a number of new significant source detections as well as failing to confirm some of those objects previously published. These new source catalogues have been combined to produce the most accurate number counts to date from 2 to 12.5\\,mJy. We find $\\rm N(>4\\,mJy) =620^{+110}_{-190}\\, deg^{-2}$, $ \\rm N(>6\\,mJy)=310^{+\\phantom{0}60}_{-100}\\, deg^{-2}$, $ \\rm N(>8\\,mJy)=150^{+\\phantom{0}40}_{-\\phantom{0}60}\\, deg^{-2}$ and $\\rm N(>10\\,mJy)=\\phantom{0}40^{+\\phantom{0}20}_{-\\phantom{0}20}\\, deg^{-2}$ after correcting for the effects of incompleteness, flux-density boosting and contamination from spurious / confused detections. Furthermore the cumulative number counts appear to steepen beyond $\\rm S_{850} > 8\\, mJy$, which could indicate an intrinsic turn-over in the underlying luminosity function placing an upper limit on the luminosity (and hence mass) of a high redshift galaxy. We have also investigated the clustering properties of the bright $\\rm S_{850} >5 \\,mJy$ SCUBA population by means of 2-point angular correlation functions. We find a $\\rm \\simeq 3.5 \\sigma$ excess of pairs within the first 100 arcsec over that expected from a Poisson distribution. Fits of a standard power-law of the form $\\rm w(\\theta)=A \\theta^{-\\delta}$ to the angular correlation functions for $\\rm S_{850} >5 \\,mJy$ are limited in accuracy by the small number of source detections but appear to be broadly consistent with that measured for EROs. Nearest-neighbour analyses further support strong clustering on arcmin scales, rejecting the null hypothesis that the distribution of the submm sources is random at the 95\\% confidence level for $\\rm S_{850} >5\\,mJy$, and at the 99\\% confidence level for $\\rm S_{850} >7\\,mJy$. ",
        "introduction": "Over the past seven years, a series of complementary deep $850\\,{\\rm \\mu m}$ surveys (eg. Smail et al. 1997, Hughes et al. 1998, Blain et al. 1999, Barger et al. 1998, Barger et al. 1999, Eales et al. 2000, Scott et al. 2002, Cowie et al. 2002, Borys et al. 2003, Webb et al. 2003a) carried out using SCUBA (Holland et al. 1999) on the JCMT has resolved the bulk of the far-infrared (FIR) extragalactic background into discrete sources. These surveys vary in size and depth from ultra-deep surveys exploiting gravitational lensing from intervening clusters to study the very faintest submm sources (Smail et al. 1997, Cowie et al. 2002), small and deep blank field surveys such as the HDF (6 sq. arcmin to a uniform $1\\sigma_{\\rm rms}\\simeq 0.5$ mJy/beam; Hughes et al. 1998, Serjeant et al. 2003), through to moderate area and comparatively shallower blank field surveys such as the ``SCUBA 8-mJy Survey'' (a total of 250 sq. arcmin to a uniform $1\\sigma_{\\rm rms}\\simeq 2.5$\\,mJy/beam; Scott et al. 2002). These surveys have revealed a population of heavily dust-enshrouded high-redshift ($z>1$) morphologically-irregular galaxies. Although there is still considerable uncertainty in the fraction of SCUBA sources hosting either low luminosity or Compton-thick AGNs (Smail et al. 2002, Ivison et al. 2002), deep X-ray observations with the Chandra and XMM-Newton telescopes have suggested that even when an AGN is present, it rarely dominates the far-infrared / submillimetre emission from the galaxy (Frayer et al. 1998, Alexander et al. 2003). Hence, if the thermal dust emission is dominated by reprocessed starlight,  the inferred star formation rates (SFRs) in the very brightest submillimetre galaxies may be as high as several thousand solar masses per year. This is sufficient to form the most massive elliptical galaxies observed in the present-day Universe on timescales of $\\sim 1$\\,Gyr, implying that the submillimetre sources may be the progenitors of todays massive ellipticals, and that the star-formation rate density in the early Universe was at least a factor of 2 larger than that derived from optical / UV observations (Steidel et al. 1999).  A key test of whether the bright submillimetre sources really are the progenitors of present-day massive spheroids is to measure their clustering properties. If they are indeed proto-massive ellipticals then they should be strongly clustered. This is an inevitable result of gravitational collapse  from Gaussian initial density fluctuations: the rare high-mass peaks are strongly biased with respect to the mass (eg. Benson et al. 2001). Each of the submillimetre survey consortia have performed their own reduction, source extraction and simulations on their individual datasets, in order to study the nature of the SCUBA population in general. However, each survey alone images only $\\sim 100 - 200$ sq. arcmin of sky spread over several fields, resulting in discrepancies in the cumulative number counts at $\\rm S_{850} > 8\\,mJy$ by over a factor of 5 due to cosmic variance, and also potentially the effects of clustering. Furthermore, no individual survey has identified enough sources to make a significant measurement of the clustering properties, although tentative evidence was obtained for strong clustering on scales of 1-2 arcmin from the ``SCUBA 8\\,mJy Survey'', the largest of these surveys undertaken to date (Scott et al. 2002).  In this paper we combine the data from all of the blank field surveys completed up to the completion date of the ``SCUBA 8\\,mJy Survey''. In order to do this we downloaded the raw data for the ``Canada UK Deep Submillimetre Survey (CUDSS)'', the ``Hawaii Flanking Fields Survey'' and the ``Hubble Deep Field (HDF) North Pencil Beam Survey'' from the Canadian Astronomy Data Centre (CADC) archive and re-reduced it in an identical manner to that employed in the ``SCUBA 8\\,mJy Survey''. This is decribed in Section 2. The source extraction algorithm developed by Scott et al. (2002) was used to identify significant submillimetre sources in each of the fields, and this method is described briefly in Section 3. In Section 4 we present and discuss the results of simulations carried out both in conjunction with the real data, and through the production of fully-simulated survey areas, for all of the survey fields. Comparative reduction methods and source lists down to a signal-to-noise ratio of 3.00 are given in Section 5. These lists include some previously unidentified significant sources but also cast doubt on the reality of other published detections. Section 6 combines these new catalogues into a master source list, from which the most accurate cumulative number counts to date in the flux density range $2-12.5$\\,mJy have been calculated. A number of models are tested against these data points. In Section 7, the new master catalogue has been used in an attempt to measure the clustering properties of the bright ($>5$\\,mJy) submillimetre sources, through 2-point angular correlation functions and nearest-neighbour analyses. Section 8 outlines the resulting motivation and aims of the current blank field survey work being undertaken by the JCMT + SCUBA, and the conclusions from this work are presented in Section 9. ",
        "conclusions": "Deep blank-field surveys conducted with SCUBA on the JCMT have successfully resolved $\\simeq 30-50\\%$ of the far-infrared extragalactic background into discrete sources down to the confusion limit of $\\rm S_{850}\\simeq 2-3\\,mJy$, with deeper surveys making use of gravitational lensing from intervening massive clusters probing the very faintest submillimetre sources. This has revealed a population of heavily dust-enshrouded galaxies at high redshift ($z>1$) undergoing a burst of massive star-forming activity. The work presented here is an investigation into the nature of the most luminous $\\rm 850\\,\\mu m$ sources ($\\rm S_{850}>5\\,mJy$), with particular consideration of their link with the formation and evolution of the most massive elliptical galaxies visible in the present-day Universe.  Combining the data from the ``SCUBA 8\\,mJy Survey'' with other existing blank field surveys, re-reduced and analysed in an identical manner, approximately doubles the area of sky observed with SCUBA analysed in a homogeneous manner. The combined datasets allow the determination of the most accurate source counts to date for $\\rm S_{850}>2\\,mJy$. The differential source counts for $\\rm S_{850}<8\\,mJy$ are well satisfied by a power-law of the form $\\rm dN(>S)/dS=N_{0}/(a+ S^{\\alpha})$ where the best-fit values of the free parameters were $\\rm N_{0} = 2.67 \\times 10^{4}$, $\\rm a=0.49$ and $\\rm \\alpha=3.14$. If one assumes this is true over the full range of flux densities from the very faint to the very bright, this predicts a total $\\rm 850 \\, \\mu m$ background of $\\rm 3.8 \\times 10^{4}\\,mJy\\,deg^{-2}$, mid-way between the $\\rm 850 \\, \\mu m$ extragalactic values of $\\rm 3.1 \\times 10^{4}\\,mJy\\,deg^{-2}$ and $\\rm 4.4 \\times 10^{4}\\,mJy\\,deg^{-2}$. However, there are some indications that the source counts steepen for $\\rm S_{850}>8\\,mJy$. This may simply be a result of small number statistics, but is possibly indicative of a high-mass cutoff of the SCUBA population. A steepening of the source counts would also make the bright SCUBA population more prone to the effects of gravitational lensing. More physical models of the source counts, beginning with a local $\\rm 850\\, \\mu m$ luminosity function, are able to reasonably successfully model the number counts using only strong pure luminosity evolution with redshift, and assuming a single temperature optically thin greybody to represent the thermal spectral energy distribution of the dust with characteristic dust temperatures in the range $\\rm 30<T_{d}<50\\,K$, and emissivity index in the range $\\rm 1.0 < \\beta < 2.0$.  Tentative indications that the bright SCUBA population clusters on scales of $\\rm 1-2$ arcmin were first identified in the ``SCUBA 8\\,mJy Survey'' (Scott et al. 2002), but these results proved inconclusive due to the small number of significant detections. However, the full source catalogue derived from combining all of the survey fields contains a sufficient number of significant detections to place more meaningful constraints on the clustering properties of the bright submillimetre population. Measurements of angular correlation functions and nearest neighbour statistics for $\\rm S_{850}>5$, 6, and 7\\, mJy, imply strong clustering on scales of $\\sim 40-50$\\,arcsec at a significance level of $3.5 - 4\\sigma$. Attempts have been made to fit the angular correlation function data points with a traditional power law of the form $\\rm w(\\theta)=A\\theta^{-\\delta}$, although the small number of sources detected even by combining all of the existing blank field surveys together mean that the values of A and $\\rm \\delta$ are not well constrained. However, the data points are in fairly good agreement with the clustering strength measured for extremely red objects by Daddi et al. (2000), suggesting a strong link between the SCUBA and ERO populations.  The results from this work have provided a major part of the motivation for a new wider-field extragalactic submillimetre survey, which has now been underway for three years at the JCMT. The aim of this survey is to map 0.5 square degrees of sky, $\\sim7$ times the total area of the ``SCUBA 8\\,mJy Survey'' to a comparable depth of $\\rm \\sigma_{rms}\\simeq 2.5\\,mJy/beam$ at $\\rm 850\\, \\mu m$. The ``SCUBA Half Degree Extragalactic Survey (SHADES)'' is split over two regions of sky; the Lockman Hole East (a continuation of the 8\\,mJy Survey), and the Subaru-XMM Deep Field. Based on the $\\rm 850 \\, \\mu m$ source counts, a survey of this nature is expected to yield $200-300$ sources detected at the $\\rm >3.5\\sigma$ level. This survey also has associated deep radio, mid-infrared, near-infrared, optical and X-ray imaging. Early example results from SHADES are discussed by Dunlop (2005), while full description of the SHADES science goals, design and data analysis techniques is given in Mortier et al. (2005). In van Kampen et al. (2005) a variety of theoretical models describing galaxy formation and evolution are presented with particular emphasis on their predictions regarding the redshift distribution and clustering properties of the SCUBA population. The first results on the sub-millimetre number counts from SHADES will be presented by Coppin et al. (2006), while the results of a search for radio and mid-infrared identifications in deep VLA and Spitzer data will be presented by Ivison et al. (in preparation)."
    },
    "0608/astro-ph0608161_arXiv.txt": {
        "abstract": "Spiral shocks are potentially a major source of turbulence in the interstellar medium.  To address this problem quantitatively, we use numerical simulations to investigate gas flow across spiral arms in vertically stratified, self-gravitating, magnetized models of galactic disks.  Our models are isothermal, quasi-axisymmetric, and local in the quasi-radial direction while global in the vertical direction.  We find that a stellar spiral potential perturbation promptly induces a spiral shock in the gas flow. For vertically stratified gas disks, the shock front in the radial-vertical plane is in general curved, and never achieves a steady state. This behavior is in sharp contrast to spiral shocks in two-dimensional (thin) disks, which are generally stationary.  The non-steady motions in our models include large-amplitude quasi-radial flapping of the shock front.  This flapping feeds random gas motions on the scale of the vertical disk thickness, which then cascades to smaller scales. The induced gas velocity dispersion in quasi-steady state exceeds the sonic value for a range of shock strengths, suggesting that spiral shocks are indeed an important generator of turbulence in disk galaxies. ",
        "introduction": "Turbulence in the interstellar medium (ISM) is observed to be pervasive and highly supersonic.  Shocks created by random gas motions produce a rich variety of structures in the diffuse, atomic ISM as well as in gravitationally-bound giant molecular clouds (GMCs). Recent work has shown that turbulence is crucial to control of star formation within GMCs, and it is also believed to affect GMC formation processes. Turbulence thus regulates star formation on both local and global scales (e.g., \\citealt{bal06}).  In the absence of driving, shock dissipation and nonlinear cascades cause turbulence to decay on a time scale comparable to flow crossing times even for a medium with equipartition-strength magnetic fields (e.g., \\citealt{sto98,mac99,pad99}), amounting to a few tens of Myr for the diffuse ISM.  This implies that the ISM must be continuously stirred by one or (likely) more driving sources.  A number of mechanisms have been proposed to drive turbulence, including \\ion{H}{2} region expansion, supernova explosions, and fluid instabilities involving magnetic fields and gravity (see \\citealt{mac04,elm04} for recent reviews).  Although less well recognized, galactic spiral shocks are also an appealing means to generate ISM turbulence. \\citet{woo76} showed that cloud deformation by passage of a shock, and the associated Kelvin-Helmholtz instabilities, together trigger random motions in the cloud.  \\citet{wad04} suggested that an in-plane ``wiggle instability'' of spiral shocks can also drive turbulent motions, although this appears to be suppressed when three-dimensional effects are included \\citep[hereafter Paper I]{kim06}.  Very recently, \\citet{bon06} and \\citet{dob06} demonstrated that passage of a distribution of clouds through a spiral shock gives rise to internal turbulent motions that follow Larson's (1981) empirical scaling law fairly well.  Most of the studies cited above investigate two-dimensional dynamics driven by spiral shocks, neglecting the vertical degree of freedom.  When the vertical dimension is taken into account and well-resolved, Paper I showed that non-steady motions of the spiral shock develop.\\footnote{The early low-resolution simulations of \\citet{tub80} did not show this effect.} Non-steady behavior of spiral shocks and generation of vertical motions has in fact been seen in previous numerical models by other authors (e.g., \\citealt{mar98,gom02,gom04,bol06}).  In this Letter, we clarify the physical causes of non-stationarity for spiral shocks in stratified disks.  We also quantify the level of the induced gas motions to show that vertical spiral shocks should indeed be an important source of turbulence in spiral galaxies. ",
        "conclusions": "\\begin{figure} \\epsscale{1.0} \\plotone{fig3.eps} \\caption{({\\it a}) Time evolution of the density-weighted velocity dispersions $\\sigma_i$ in model G. ({\\it b}) Total velocity dispersion $\\sigma_{\\rm tot}$ {\\it vs.} the effective Mach number $\\mathcal{M}_{\\rm eff}$ of the spiral shock in models B to G. The errorbars represent the standard deviations in the temporal fluctuations of $\\sigma_i$. \\label{fig-sig}} \\end{figure}  Galactic rotation can supply an effectively inexhaustible amount of kinetic energy to power turbulence in the ISM \\citep{fle81}.\\footnote{To maintain a fixed level of turbulence, the energy dissipation per unit mass per unit time must be balanced by the net stresses, $\\Omega \\langle\\rho v_x \\delta v_y - B_x B_y/(4\\pi)\\rangle/\\bar\\rho$ (e.g. \\citealt{pio05}). If this power is supplied at the expense of overall ISM accretion toward the galactic center, the accretion time is $\\sim (2\\Omega)^{-1} \\bar\\rho (\\Omega R)^2 \\langle \\rho v_x \\delta v_y \\rangle^{-1}$,  which exceeds the Hubble time if $\\langle \\rho v_x \\delta v_y - B_x B_y/(4\\pi) \\rangle \\sim \\bar \\rho c_s^2$.} In this Letter, we have shown that spiral shocks are in general non-stationary in the radial-vertical plane, and can efficiently transform some of the available bulk rotational energy into random gas motions.  Vertical force imbalance drives radial flapping of the shock front at $|z|/H \\simgt 0.2$, which then feeds turbulence at smaller scales.  The random gas motions induced by the spiral shock persist despite strong shock dissipation, and yield time-averaged in-plane velocity dispersions $\\sim 7-10\\kms$ for a range of shock strengths, similar to the observed line widths of cold and warm atomic gas in the Milky Way (e.g., \\citealt{hei03}).  Vertical velocity dispersions are lower, but still amount to $\\sim 1/2$ of the thermal velocity dispersion.  Although our presentation has focused on shocks with realistic parameters, in fact turbulence generation by spiral shocks appears to require neither self-gravity nor magnetic fields. Our model simulations, including those unlisted in Table~\\ref{tbl-model}, show that anything (e.g., an imposed variation of $\\Phi_{\\rm sp}$ with $z$) that makes the primary shock non-vertical will end up producing a non-steady flow -- simply because vertical oscillation periods need not (and in general do not) agree with the horizontal crossing time between arms.  Since strong spiral arms usually imply a high rate of star formation, spiral shocks and stellar energy sources may work together in generating turbulence in many galaxies.  The absence of observed correlation between spiral arm phase and turbulent amplitude (e.g., \\citealt{dic90}), however,  suggests that these two processes are not the only important sources of turbulence.  In addition, radio observations of extended \\ion{H}{1} disks in face-on galaxies show that the total vertical velocity dispersions are as large in the non-star-forming outer parts as in the star-forming inner regions \\citep{van99}, again suggesting that additional sources of turbulence are able to compensate when needed.  One compensating source of turbulence may be provided by the magnetorotational instability (MRI) \\citep{sel99,kim03,dzi04}.  The MRI should naturally be strongest exactly where spiral shocks are absent: in the outer galaxy, where gravity is weak, the disk flares, and the density drops; and in interarm regions, where the angular velocity decreases outward (in arms, shear is reversed so that MRI cannot occur).  While in this Letter we adopt isothermal conditions for the gas and take the mean magnetic field parallel to the spiral arm, the real ISM has a multi-phase structure and is threaded also by weak vertical magnetic fields.  \\citet{pio04,pio05} demonstrated that MRI is able to generate random gas motions up to $\\sim 8\\kms$ in two-phase gas under favorable conditions.  It will be interesting to see how turbulent driving in spiral shocks develops for a multiphase medium, and whether MRI and spiral shocks can indeed establish a geographically-compensating balance of power."
    },
    "0608/astro-ph0608175_arXiv.txt": {
        "abstract": "{We performed an extended analysis of the parameter space for the interaction of the Magellanic System with the Milky Way (MW). The varied parameters cover the phase space parameters, the masses, the structure, and the orientation of both Magellanic Clouds, as well as the flattening of the dark matter halo of the MW. The analysis was done by a specially adopted optimization code searching for the best match between numerical models and the detailed H\\,I map of the Magellanic System by Br\\\"uns et al.~(\\cite{Bruens05}). The applied search algorithm is a genetic algorithm combined with a code based on the fast, but approximative restricted N--body method. By this, we were able to analyze more than $\\mathrm{10^6}$ models, which makes this study one of the most extended ones for the Magellanic System. Here we focus on the flattening $q$ of the axially symmetric MW dark matter halo potential, that is studied within the range $0.74 \\leq q \\leq 1.20$. We show that creation of a trailing tail (Magellanic Stream) and a leading stream (Leading Arm) is quite a common feature of the Magellanic System--MW interaction, and such structures were modeled across the entire range of halo flattening values. However, important differences exist between the models, concerning density distribution and kinematics of H\\,I, and also the dynamical evolution of the Magellanic System. Detailed analysis of the overall agreement between modeled and observed distribution of neutral hydrogen shows that the models assuming an oblate ($q < 1.0$) dark matter halo of the Galaxy allow for better satisfaction of H\\,I observations than models with other halo configurations. ",
        "introduction": "The idea of \\emph{dark matter} (DM) was introduced by Zwicky~(\\cite{Zwicky33}). His dynamical measurements of the mass--to--light ratio of the Coma cluster gave larger values than those known from luminous parts of nearby spirals. That discrepancy was explained by the presence of DM. Ostriker et al.~(\\cite{Ostriker74}) proposed that DM is concentrated in a form of extended galactic halos. Analysis of rotation curves of spiral galaxies (Bosma~\\cite{Bosma81}; Rubin\\,\\&\\,Burstein~\\cite{Rubin85}) denotes that their profiles cannot be explained without presence of non--radiating DM. Hot X--ray emitting halos have been used to estimate total galactic masses (McLaughlin~\\cite{McLaughlin99}). Corresponding mass--to--light ratios exceed the maximum values for stellar populations, and DM explains the missing matter naturally. The presence of DM halos is expected by the standard CDM cosmology model of hierarchical galaxy formation.  The classical CDM halo profile (NFW) is simplified to be spherical. However, most CDM models expect even significant deviations from the spherical symmetry of DM distribution in halos.  The model of formation of DM halos in the universe dominated by CDM by Frenk et al.~(\\cite{Frenk88}) produced triaxial halos with a preference for prolate configurations. Numerical simulations of DM halo formation by Dubinski\\,\\&\\,Carlberg~(\\cite{Dubinski91}) are consistent with halos that are triaxial and flat, with (c/a) = 0.50 and (b/a) = 0.71. There are roughly equal numbers of dark halos with oblate and prolate forms.  Observationally, the measurement of the shape of a DM halo is a difficult task. A large number of varying techniques found notably different values, and it is even not clear if the halo is prolate or oblate. Olling\\,\\&\\,Merrifield~(\\cite{Olling00}) use two approaches to investigate the DM halo shape of the Milky Way (MW), a rotation curve analysis and the radial dependence of the thickness of the H\\,I layer. Both methods lead consistently to flattened oblate halos.  Recently, the nearly planar distribution of the observed MW satellites, which is almost orthogonal to the Galactic plane, raised the question if they are in agreement with cosmological CDM models (Kroupa et al.~\\cite{Kroupa05}) or if other origins have to be invoked.  Zentner et al.~(\\cite{Zentner05}) claim that the disk--like distribution of the MW satellites can be explained, provided the halo of the MW is sufficiently prolate in agreement with their CDM simulations. On the other hand, it is not clear if there exists a unique prediction of the axis ratios from CDM simulations, as the scatter in axis ratios demonstrates (Dubinski\\,\\&\\,Carlberg~\\cite{Dubinski91}). Based on $\\Lambda$CDM simulations, Kazantzidis et al.~(\\cite{Kazantzidis04}) emphasize that gas cooling strongly affects halo shapes with the tendency to produce rounder halos.  Another promising method to determine the Galactic halo shape are stellar streams because they are coherent structures covering large areas in space. Thus, their shape and kinematics should be strongly influenced by the overall properties of the underlying potential.  A good candidate for such an analysis is the stellar stream associated with the Sagittarius dwarf galaxy. By comparison with simulations, Ibata et al.~(\\cite{Ibata01}) found that the DM halo is almost spherical in the galactocentric distance range from 16 to 60 kpc. Helmi~(\\cite{Helmi04}) studied the Sagittarius dwarf stream and argues for a non--spherical shape of the Galactic halo. However, Helmi~(\\cite{Helmi04}) also warned that the Sagittarius stream might be dynamically too young to allow for constraints on the halo shape. Johnston et al.~(\\cite{Johnston05}) argue that the analysis of the orbital planes of the Sagittarius dwarf galaxy prefers oblate over prolate models. In flattened oblate halos, the non--polar satellite orbits tend to become co--planar and contribute to the disk formation. In the case of Monoceros tidal stream discussed by Pe\\~narrubia et al.~(\\cite{Pena05}), the best--fitting models predict a significantly oblate halo.  In this paper we use the Magellanic Stream to derive constraints on the halo shape of the MW. The basis are the new detailed H\\,I observations of the Magellanic System (including the Large Magellanic Cloud (LMC) and the Small Magellanic Cloud (SMC)) by Br\\\"uns et al.~(\\cite{Bruens05}).  As remnants of the LMC--SMC--MW interaction, extended structures connected to the System are observed. Among them, the \\textit{Magellanic Stream} -- an H\\,I tail originating in between the Clouds and spreading over $\\mathrm{\\approx 100^\\circ}$ of the plane of sky -- has been a subject of investigation for previous studies (see, e.g.,\\ Fujimoto\\,\\&\\,Sofue~\\cite{Sofue76}; Lin\\,\\&\\,Lynden--Bell~\\cite{Lin77}; Murai\\,\\&\\,Fujimoto~\\cite{Murai80}; Heller\\,\\&\\,Rohlfs~\\cite{Heller94}; Gardiner\\,\\&\\,Noguchi~\\cite{Gardiner96}; Bekki\\,\\&\\,Chiba~\\cite{Bekki05}; Mastropietro et al.~\\cite{Mastropietro05}). Due to the extended parameter space related to the interaction of three galaxies and also due to the high computational costs of fully self--consistent simulations, simplifying assumptions were unavoidable. Many simulations neglected the self--gravity of the individual stellar systems by applying a restricted N--body method similar to the method introduced by Toomre\\,\\&\\,Toomre~(\\cite{Toomre72}). None of these simulations considered the self--gravity of all three galaxies. Often only one galaxy is simulated, including its self--gravity by means of a live disk and halo, whereas the other two galaxies are taken into account by rigid potentials of high internal symmetry. That is, none of the simulations so far adopted a live dark matter halo of the MW, but they applied \\mbox{(semi--)analytical} descriptions for the \\textit{dynamical friction} between the Magellanic Clouds and the MW. Also, a possible flattening of the MW halo has not been considered.  Having the numerical difficulties in mind, it is not surprising that a thorough investigation of the complete parameter space was impossible.  Modeling observed interacting galaxies means dealing with an extended high--dimensional space of initial conditions and parameters of the interaction. Wahde~(\\cite{Wahde98}) and Theis~(\\cite{Theis99}) introduced a \\emph{genetic algorithm} (GA) as a robust search method to constrain models of observed interacting galaxies. The GA optimization scheme selects models according to their ability to match observations. Inspired by their results, we employed a simple fast numerical model of the Magellanic System combined with an implementation of a GA to perform the first very extended search of the parameter space for the interaction between LMC, SMC, and MW. Here we present our results about the MW DM halo flattening values compatible with most detailed currently available H\\,I Magellanic survey (Br\\\"uns et al.~\\cite{Bruens05}, see Figs.~\\ref{2d_HI_map}, and~\\ref{pic_obs_radvel}). ",
        "conclusions": "We performed an extended analysis of the parameter space for the interaction of the Magellanic System with the Milky Way. The varied parameters cover the phase space parameters, the masses, the structure, and the orientation of both Magellanic Clouds as well as the flattening of the dark matter halo of the MW. The analysis was done by a specially adopted optimization code searching for a best match between numerical models and the detailed HI map of the Magellanic System by Br\\\"uns et al.~(\\cite{Bruens05}). The applied search algorithm is a genetic algorithm combined with a code based on the fast, but approximative restricted N--body method. By this, we were able to analyze more than $\\mathrm{10^6}$ models, which makes this study one of the most extended ones for the Magellanic System.  In this work we focused especially on the flattening of the MW dark matter halo potential $q$. The range $0.74 \\leq q \\leq 1.20$ was studied. It is equivalent to the interval of the density flattening $0.31 \\leq q_\\rho \\leq 1.37$ (see Eq.~\\ref{log_dens_flattening_2}).  We showed that the creation of a trailing tail (Magellanic Stream) and a leading stream (Leading Arm) is quite a common feature of the LMC--SMC--MW interaction, and such structures were modeled across the entire range of halo flattening values. However, important differences exist between the models, concerning density distribution and kinematics of H\\,I, and also dynamical evolution of the Magellanic System over the last 4\\,Gyr. In contrast to Murai\\,\\&\\,Fujimoto~(\\cite{Murai80}), Gardiner et al.~(\\cite{Gardiner94}), or Lin et al.~(\\cite{Lin95}), the Clouds do not have to be gravitationally bound to each other for the entire evolutionary period to produce the matter distribution that is in agreement with currently available H\\,I data on the Magellanic System.  Overall agreement between the modeled and observed distribution of neutral hydrogen in the System is quantified by the fitness of the models. The fitness value is returned by a $FF$, that performs a very detailed evaluation of every model (Appendix~\\ref{appendixB}). Analysis of fitness as a function of the halo flattening parameter $q$ indicates that the models assuming oblate DM halo of MW (model A) allow for better satisfaction of H\\,I observations than models with other halo configurations. Analysis of Fig.~\\ref{pic_fitness-q_max} does not indicate a drop in the value of fitness as the flattening parameter $q$ decreases. It suggests that models of a quality comparable to the fitness of model A may exist for even more oblate configurations of the MW DM halo. Unfortunately, the logarithmic potential defined by Eq.~\\ref{log_gravity} cannot be used for $q<1/\\sqrt{2}$ due to negative density on the axis of symmetry. Thus, the result represented by Fig.~\\ref{pic_fitness-q_max} established a strong motivation for further extension of the parameter study using a more general and realistic model of the Galactic halo (e.g., a tri--axial distribution of matter or the flattening parameter depending on position). Nevertheless, the above mentioned problem does not alter our result preferring flattened halo configurations and discriminating nearly--spherical shapes within the galactocentric radius of $\\approx 200$\\,kpc containing typical orbits of the Clouds and so analyzed in this study.  We did not involve surveys of stellar populations in the Magellanic System in the process of fitness calculations. This is due to the nature of test--particle models that do not allow for the distinction between stellar and gaseous content of studied systems. However, we still have to face one of the most interesting observational facts connected to the Magellanic Clouds -- the absence of stars in the Magellanic Stream (Van den Bergh~\\cite{Vandenbergh00}) -- because both the LMC and SMC contain stellar populations, and so every structure emanating from the Clouds should be contaminated by stars. It is an additional constraint on the models. It cannot be involved in fitness calculation because of the limits of our numerical code, but has to be taken into account.  Stellar populations of the SMC are very young and the mass fraction in the form of stars is extremely low. Our models show that the evolution of the Magellanic Stream has been lasting $\\mathrm{2\\,Gyr}$ at least (model B). Thus, the fraction of matter in the Magellanic Stream, that is of SMC origin, was torn off before significant star formation bursts occurred in the SMC, and stars should not be expected in the Stream. Nevertheless, we found both LMC and SMC matter in the Magellanic Stream for every model of the System. Similarly to the case of the SMC, if the LMC star formation activity was increased after the matter transport into the Magellanic Stream was triggered, stars would naturally be missing in the Stream. Such a scenario is doubtful however. Observational studies argue for a massive star formation burst started in LMC at $\\mathrm{\\approx -3\\,Gyr}$ (Van den Bergh~\\cite{Vandenbergh00}), which is rather close to the age of the Magellanic Stream, as indicated by our models (see Sect.~\\ref{stream_struct}). Our results concerning the LMC and SMC orbits introduced an acceptable solution to the problem of missing stars. We showed that the evolution of the Clouds in aspherical MW DM halos (models A and C) does not lead to extremely close encounters disturbing inner parts of the LMC disk ($r_\\mathrm{disk} < 5$\\,kpc). Since the distribution of gaseous matter in galaxies is typically more extended than the stellar content, the Magellanic Stream matter coming from outer regions of the Clouds does not necessarily have to contain a stellar fraction.  Previous discussion of stellar content of the Magellanic System supports discrimination of the configurations with nearly spherical halos (model B) that was discovered by the GA search. On the other hand, many papers on the dynamical evolution of the Magellanic Clouds dealing with a spherical MW halo (Murai\\,\\&\\,Fujimoto~\\cite{Murai80}, Gardiner et al.~\\cite{Gardiner94}, Bekki\\,\\&\\,Chiba~\\cite{Bekki05}) argue that the observed massive LMC star formation bursts $\\mathrm{3\\,Gyr}$ ago was caused by close LMC--SMC encounters. Our model B shows close approaches of the Clouds $\\Delta r \\approx 10$\\,kpc at around the mentioned time. For aspherical halos, such encounters do not induce the formation of particle streams. However, close LMC--MW and SMC--MW encounters appeared to be efficient enough to trigger massive matter redistribution in the System leading to formation of the observed structures. Then, they could also be responsible for the triggering of star bursts."
    },
    "0608/astro-ph0608496_arXiv.txt": {
        "abstract": "In gravitational lensing, steeper mass profiles generically produce longer time delays but smaller magnifications, without necessarily changing the image positions or magnification ratios between different images.  This is well known.  We find in this paper, however, that even if steepness is fixed, time delays can still have significant model dependence, which we attribute to shape modeling degeneracies. This conclusion follows from numerical experiments with models of 35 galaxy lenses.  We suggest that varying and twisting ellipticities, features that are explored by pixelated lens models but not so far by parametric models, have an important effect on time delays. ",
        "introduction": "In the gravitational lensing of quasars by galaxies, time delays between images are highly prized because they are proportional to the Hubble time \\citep{schech04,jakob05,kochanek06,morgan06,vuissoz06,scmw06}. But a given set of image positions and brightness ratios---in fact any images of sources at single redshift---can be produced by very different lensing-mass distributions.  In particular, making a lens profile steeper lengthens the time delays and reduces the overall magnification, but otherwise has little or no effect on the images.  A more precise version of the previous statement is that replacing $(1-\\kappa)$ everywhere on a lens by $\\lambda(1-\\kappa)$---where $\\kappa$ is the projected density in units of the critical density and $\\lambda$ is a constant---multiplies all time delays by $\\lambda$ and multiplies all magnifications by $\\lambda^{-2}$, but changes nothing else.  In fact the transformation only needs to be applied within a circle larger than all the images.  The simplest interpretation is a stretching of the arrival-time surface by a factor of $\\lambda$ along the time axis. Multiplying $(1-\\kappa)$ by a constant naturally makes the mass profile steeper or shallower.  That is not exactly the same as changing the radial index, but quite similar to it over the scales of interest.  This degeneracy has a long history and several names, having been independently discovered at least four times.  \\cite{fgs85} derived it as a consequence of the lens equation, and the same authors in \\cite{gfs88} named it the `magnification transformation'. \\cite{paczynski86} discovered it in the context of microlensing. \\cite{ss95} found it in cluster lensing and called it a `global invariance transformation'. \\cite{wp94} came upon it as a parameter degeneracy in galaxy-lens models.  Nowadays the common name is `mass-sheet degeneracy'; ADS first shows the phrase used by \\cite{bn95}, but it seems the name was already in spoken usage by then. Unfortunately, the name `mass-sheet degeneracy' can give the incorrect impression that simply adding/removing a mass sheet is a degeneracy.  It seems preferable to use the more descriptive term {\\em steepness degeneracy,} thus avoiding the possible confusion.  In this paper we will use `steepness degeneracy' in both strict and rough senses: the strict meaning being rescaling $(1-\\kappa)$ within a circle enclosing all the images, and the rough meaning being changing the radial index.  Whatever the name, the steepness degeneracy has been much discussed in recent years \\citep{bradac04,schech04,treu04,oguri02,wuck02}. On the other hand, there has been little research on whether any other degeneracies are important for the time-delay problem.  Several known lensing degeneracies are summarized in \\cite{saha00}, along with a derivation of the arrival-time interpretation above, but apart from steepness and the obvious monopole degeneracy, none of them are applicable in the context of lensed quasars.  It is easy to imagine further degeneracies: we can simply make the stretching factor a function of position.  In other words we replace the arrival-time surface $\\tau(\\btheta)$ by \\begin{equation} \\tau'(\\btheta) = \\lambda(\\btheta) \\tau(\\btheta) . \\label{trans} \\end{equation} We must require $\\nabla\\lambda=0$ at the image positions to preserve said image positions, $\\nabla(\\lambda\\tau)\\neq0$ except at the images so as not to introduce new images, and $\\nabla^2(\\lambda\\tau)\\geq0$ everywhere to keep the density non-negative.  But otherwise the transformation (\\ref{trans}) are arbitrary. We may call such transformations {\\em shape degeneracies,} because they change the shape of the arrival-time surface and the mass profile in some complicated way.  General shape degeneracies change magnification ratios between different images and time-delay ratios between different pairs of images, though particular shape degeneracies may preserve some or all of these. In contrast, the steepness degeneracy preserves all time-delay ratios and magnification ratios.  Hence the effect of steepness degeneracies will be reduced if such data are present.  If sources at multiple redshifts are present, then steepness degeneracy is broken, while shape degeneracies can be greatly reduced.  The only explicit example of a shape degeneracy in the literature is a special but intriguing model constructed by \\cite{zq03}, to which we will return later.  The main aim of this paper, however, is to assess whether shape degeneracies are important in galaxy lenses independently of particular examples.  We can do so using pixelated modeling, which is the best available way to explore the full range of shape degeneracies because shape degeneracies are generically present in free-form lens models. (Parametric modeling, on the other hand, allows only for a restricted set or sets of shape degeneracies.)  The trick is to somehow `turn off' the steepness degeneracy, and then see how degenerate time delays remain. ",
        "conclusions": "\\label{disc}  The steepness degeneracy in lensing is now well understood.  The above numerical experiments attempt to estimate the effect of other degeneracies.  This is done by searching through mass models at fixed image-positions, time-delay ratios (where applicable), and mean annular density $\\kann$.  The additional degeneracies, quantified approximately by $\\Delta\\varphi/\\varphi$ at fixed $\\kann$, turn out for some lenses to be as important as steepness.  What then are the additional important degeneracies beyond steepness? Do common parametric forms for lenses already allow for the other degeneracies, and if not, what new parameters are needed? Detailed answers to these questions require more research, but we can deduce partial answers by thinking about the arrival-time surface. In the Introduction we classified degeneracies into steepness and shape, with the stipulation that the latter category can be further subdivided depending on how many image observables we care to consider. In this Section we go a little further and attempt a more quantitative, but still intuitive classification.  Recall that the steepness degeneracy amounts to a homogeneous stretching or shrinking of the time scale in the arrival-time surface. Imagine now that we stretch the time scale on the E side and shrink it on the W side, preserving the image positions. No change is required in the circularly averaged $\\kann$. The resulting models are not steepness-degenerate, but the time delay between E and W images will change, producing a shape-degeneracy transformation. This particular kind is allowed only in asymmetric lenses, but there it may well be as important as the steepness degeneracy.  Next, let us imagine stretching the time scale on the E and W quadrants while shrinking it on the N and S quadrants.  Such a transformation, allowed in inversion symmetric lenses, is likely to most affect core quads, and inclined quads and doubles to a lesser extent, but not axial systems.  Further, we can imagine a transformation that shrinks the time scale at small radii and stretches it at large radii.  We can thus imagine a hierarchy of lensing degeneracies, from an $m=0$ mode (the steepness degeneracy) through $m=1,2,$ etc.\\ representing various shape degeneracies. This is reminiscent of basis functions in cylindrical coordinates, but we emphasize that shape degeneracies are not additive modes in the arrival-time surface, still less so in the mass profile --- they are multiplicative modes in the arrival-time surface, and in the mass profile their form will be more complicated.  The steepness degeneracy is special in that it rescales the arrival time surface homogeneously, leaving time-delay ratios and magnification ratios unaffected, while there is no guarantee that shape degeneracies will preserve time-delay and magnification ratios. The image elongation information by itself, as measured in weak lensing does not break shape degeneracies, but having having many weakly lensed images would help to constrain the shape of the arrival time surface. Sources at multiple redshifts will break steepness, and help reduce shape degeneracy.  We can try and guess the sort of mass-profile feature that will produce an $m=2$ mode.  By analogy with the steepness degeneracy, suppose an elliptical mass profile is steeper along the long axis than the short axis; this corresponds to ellipticity decreasing with increasing radius, and it seems plausible that it will increase time delays along the long-axis direction and decrease delays along the short-axis direction.  In general we suggest that ellipticity varying or twisting with radius as the signature of $m=2$ and higher modes. Re-examining our early models of the inclined quad B1115+080 \\citep{sw97} the role of such features in fitting time delays is already apparent; at the time we commented briefly on it but had no interpretation.  The above suggests interpreting the degeneracy given by \\cite{zq03} as a mixture of steepness and shape degeneracies.  Their Fig.~2 illustrates the transformation of an arrival-time surface, which appears to be an $m=2$ stretching/shrinking followed by an $m=0$ stretching with the effects canceling at the image positions but not globally.  (Note that the left- and right-hand sides of their arrival-time plot actually correspond to a $90^\\circ$ change of position angle, not $180^\\circ$.)  In the Zhao-Qin example, the ellipticity in the potential comes entirely from external shear and the main lens is circular.  But in our Figs.~\\ref{map1}--\\ref{map4}, varying and twisting ellipticity is a common feature, especially in inclined systems.  The axial systems in these figures tend not to show twisting ellipticity.  Recall also from our numerical results that axial systems like B1422+231 tend to have the lowest $\\Delta\\varphi/\\varphi$, that is to say, steepness dominates.  Individual models of axial systems may still contain twisting ellipticity; however, clockwise and anti-clockwise twists are equivalent if the image morphology is axial, hence such twists will tend to cancel in the ensemble average. For inclined image morphologies, clockwise and anti-clockwise twists in the density are not equivalent, and will tend to survive in an ensemble average.  We may ask whether the pixelated method tends to exaggerate twisting ellipticity.  The blind tests in \\cite{ws00} are reassuring in this regard; no spurious twisting appears in the ensemble-average models.  We remark that in galaxy dynamics, twisting ellipticity arises naturally in at least two ways: differential rotation leading to spiral features, and projection of triaxial features. Because these and other shape features can be important in real lenses, the errors in derived $H_0$ must incorporate all of the degeneracies\\citep{scmw06}.  The arguments in this Discussion are hand-waving, but they indicate that the issue of varying/twisting ellipticity needs closer attention.  One project that is now called-for is to map the degeneracies in pixelated models in detail, using principal components analysis or similar on model ensembles, to see if a hierarchy of degeneracies indeed emerges. Another project is to incorporate ellipticities that can vary or twist with radius into parametric models."
    },
    "0608/astro-ph0608669_arXiv.txt": {
        "abstract": "{We present a detailed study on the kinematics of metal-rich stars with and without planets, and their relation with the Hyades, Sirius and Hercules dynamical streams  in the solar neighbourhood. Accurate kinematics have been derived for all the stars belonging to the CORALIE planet search survey. We used precise radial velocity measurements and CCF parameters from the CORALIE database, and parallaxes, photometry and proper motions from the HIPPARCOS and Tycho-2 catalogues. The location of stars with planets in the thin or thick discs has been analysed using both kinematic and chemical constraints. We compare the kinematic behaviour of known planet-host stars to the remaining targets belonging to the volume-limited sample, in particular to its metal-rich population. The high average metallicity of the Hyades stream is confirmed. The planet-host targets show a kinematic behaviour similar to that of the metal-rich comparison subsample, rather than to that of the comparison sample as a whole, thus supporting a primordial origin for the metal excess observed in stars with known planetary companions. According to the scenarios proposed as an explanation for the dynamical streams, systems with giant planets could have formed more easily in metal-rich inner Galactic regions and then been brought into the solar neighbourhood by dynamical streams. ",
        "introduction": "Precise spectroscopic studies have revealed that stars with planets seem to be particularly metal-rich when compared with ``single'' field dwarfs \\citep[e.g.][]{Gonzalez1997,Gonzalezetal2001,Santosetal2001, Santosetal2004b,Santosetal2005}. This represents the first strong link between the planetary companion and its parent star. In the last few years we have published a series of papers on chemical abundances in stars with known giant planets. Abundances of volatile and refractory elements in a large set of planet-host stars and in a volume-limited comparison sample were presented in \\citet{Ecuvillonetal2004a, Ecuvillonetal2004b, Ecuvillonetal2006a, Bodagheeetal2003, Beiraoetal2005, Gillietal2006}. The volatile/refractory abundance ratios have recently been investigated in \\citet{Ecuvillonetal2006b}. Be and Li light elements were also analysed \\citep{Santosetal2002b,Santosetal2004c, Israelianetal2004}. We are now interested in investigating the kinematic behaviour of planet-host stars.  The kinematics of stars in the solar neighbourhood can give fundamental information for our understanding of the structure and evolution of the Milky Way. ESA's astrometric satellite {\\it Hipparcos} (ESA 1997) has provided us with accurate positions and trigonometric parallaxes, as well as absolute proper motions, for a large and homogeneous sample of tens of thousands of stars near the Sun. This has offered the opportunity to investigate the velocity distribution in the solar neighbourhood, not only for early-type stars, but also for the old population of the Galactic disc. Several studies have been performed on this topic since \\citep[e.g.,][]{Dehnen&Binney1998,Dehnen1998, Alcobe&Cubarsi2005, Chereul&Creze&Bienayme1999, Hoogerwerf&Aguilar1999, Asiainetal1999, Mignard2000}.  The kinematic distribution of stars in the Solar neighbourhood is far from smooth. The inhomogeneities can be spatially confined groups of young stars (e.g.\\ young clusters), or spatially extended groups with the same kinematics. \\citet{Eggen1994} defined a ``supercluster'' (SC) as  a group of gravitationally unbound stars that share the same kinematics and may occupy extended regions in the Galaxy, and a ``moving group'' (MG) as the part of the supercluster that enters the solar neighbourhood and can be observed all over the sky.  According to the standard interpretation, stars in a moving group were formed simultaneously in a small phase-space volume, so that we can still observe a stream of young stars with similar velocities. On the other hand, old stellar populations should be completely mixed and present a smooth distribution function, due to phase mixing and scattering processes. Actually, there are basically two factors acting against the persistence of a moving group in the general stellar background: Galactic differential rotation, which tends to spread them out very quickly in the direction of the Galactic rotation \\citep{Woolley1960}, and disc heating \\citep[e.g.][]{Lacey1991}, which produces an increase in the velocity dispersion of disc stars with age.\\citet{Wielen1977} reported that a typical star in the solar neighborhood would forget its initial peculiar velocity within $2\\times10^8$ years.  However, several works have obtained results that raise problems with this interpretation. On the one hand, it has been reported that early-type stars belonging to the same supercluster span a wide range of ages, in contradiction with the assumption of a common origin \\citep{Chereul&Creze&Bienayme1998,Chereul&Creze&Bienayme1999}. On the other hand, it is striking to verify that some of the classical MGs are several 10$^8$ yr old, and that significant clumpiness in velocity space has been reported for late-type stars \\citep{Dehnen1998}. Other mechanisms of different nature have been proposed as responsible for the substructures observed in velocity space. Streams could be the remnants of merger events between the Milky Way and a satellite galaxy, as recently proposed by \\citet{Navarroetal2004} for the Arcturus group, although the chance that two of them have left such important signature in the disc near the Sun is statistically very low\\citep{Famaeyetal2005a}. Moreover, \\citet{Raboudetal1998} reported that the local anomaly found in the ($u,v$) plane is caused by intermediate- to high-metallicity stars, while a merger would involve rather metal-poor stars \\citep{Dehnen2000}.  Another possibility is that purely dynamical mechanisms, such as non-axysimmetric components of the gravitational potential (for example, the rotating galactic bar or spiral waves), are the cause of substructures in the velocity distribution \\citep[e.g.][]{Dehnen1998, Raboudetal1998, Sellwood&Binney2002}. Stars in an MG could have formed in the inner part of the disc and have become trapped in resonant orbits. \\citet{Dehnen1999} associates the bimodality of the local velocity distribution of late-type disc stars, lately identified as the Hercules stream, with high-metallicity stars thrown out from the inner disc, and \\citet{Dehnen2000} concluded that outer Lindblad resonant scattering off the Galactic bar is presently the only viable explanation for it. Other simulations have tested the effect of other non-axysimmetric components, in particular transient spiral waves, in the Galactic disc \\citep[e.g.][]{Sellwood&Binney2002, Desimoneetal2004}, with the result that they may cause radial migration near their corotation radius.  It is not unlikely that known exoplanetary systems could also have suffered such radial displacements. The typical contributors to the kinematic structures in the solar neighbourhood are metal-rich old stars \\citep{Raboudetal1998}. The well-known metal-rich nature of planet-host stars \\citep[e.g.][]{Gonzalez1997,Gonzalezetal2001,Santosetal2001, Santosetal2004b,Santosetal2005} might thus be related with dynamical streams. It is probable that planet-host stars may have formed in protoplanetary clouds with a high metal content, such as those located in inner regions of the Milky Way, and then be brought closer to the Sun by dynamical streams \\citep{Famaeyetal2005b}. \\citet{Famaeyetal2005b} have found an indicative clump of stars hosting planets from \\citet{Santosetal2003} in the Hyades streams. Several works have investigated the kinematics of stars with extrasolar planets \\citep[e.g.][]{Gonzalez1999, Reid2002, Barbieri&Gratton2002, Santosetal2003}, without finding any significative peculiarity.  Two main scenarios have been proposed to explain the metal excess found in stars with exoplanets. The ``self-enrichment'' hypothesis suggests that the infall of a large amount of protoplanetary rocky material on to the parent star may have enhanced its primordial metallicity \\citep{Gonzalez1997,Laughlin2000}. In contrast, \\citet{Santosetal2000,Santosetal2001} have proposed a ``primordial'' origin for the iron overabundance of planet-host stars, so that planetary system formation would be triggered by a high metal content of the primordial cloud. Several works have obtained results supporting the ``primordial'' scenario \\citep{Pinsonneaultetal2001,Santosetal2004b,Reid2002, Sadakaneetal2002,Ecuvillonetal2004a,Ecuvillonetal2004b, Ecuvillonetal2006a, Ecuvillonetal2006b, Gillietal2006}. However, some clear evidence of pollution events have been reported \\citep{Gonzalez1998,Laws&Gonzalez2001,Smithetal2001,Israelianetal2001,Israelianetal2003}, so that the ``self-enrichment'' scenario cannot be completely discarded \\citep{Vauclair2004}.  This article studies the kinematics of planet-host stars and their relation with the dynamical streams in the solar neighbourhood and investigates whether it corresponds to the behaviour of the metal-rich population or not. The idea is that if the metal-rich nature of planet-host stars is intrinsic, then they should be kinematically identical to the metal-rich comparison sample, while if the metallicity excess is due to planet-related enrichment, then they should be kinematically identical to the total comparison sample.  We present accurate kinematics for all the stars belonging to the CORALIE planet search survey \\citep[see][]{Udryetal2000}. Precise radial velocity measurements and CCF parameters were extracted from this unique database, whereas HIPPARCOS \\citep{Hipparcos1997} parallaxes and photometry, and Tycho-2 \\citep{Hogetal2000} proper motions were adopted. The location of stars with planets in the thin or thick discs is analysed using both kinematic and chemical constraints. We study the kinematic behaviour of known planet-host targets with respect to the remaining stars belonging to the kinematically unbiased volume-limited sample and we investigate possible links between the presence of planets, the stellar metallicity and the dynamical streams Hyades, Sirius and Hercules.  \\begin{figure} \\centering \\includegraphics[width=8cm]{figure1.eps} \\caption{Probability distributions of location in the thin disc ({\\it upper panel}) and  the thick disc ({\\it central panel}), and the thick-disc-to-thin-disc relative probability  ({\\it lower panel}). Diamonds and crosses represent stars with and without planets, respectively. In the lower panel, the solid lines indicate the limits for reliable membership to each population.} \\label{ProbDist} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=8cm]{figure2.eps} \\caption{Thick-disc-to-thin-disc relative probability  as a function of [Fe/H]. Diamonds and crosses represent stars with and without planets, respectively. The solid lines indicate the limits for the reliable membership to each population.} \\label{MetProb} \\end{figure} ",
        "conclusions": "This is the first study carried out on a volume-limited sample of stars with and without exoplanets aimed at investigating the comparative kinematics of planet-host stars and their possible relation with the dynamical streams of the solar neighbourhood. With this goal, we have compared the group of stars harbouring planets with the rest of the CORALIE sample, distinguishing  subgroups of metallicity and age. This was made possible by the precise radial velocity measurements and the CCF parameters included in the CORALIE database, and using the HIPPARCOS and Tycho-2 catalogues. For targets with detailed abundance determinations, a complete study on whether they belong to the thick and thin disc populations, and on possible signatures in the abundance trends, has been carried out. The analyses related to the dynamical streams have been carried out using two independent methods that have produced equivalent results.  The main outcomes of this work are the following: \\begin{itemize} \\item The CORALIE planet-host target \\object{HD\\,4308} has been found to be a very probable member of the Galactic thick disc, together with another 21 comparison stars; the dependence of the abundance trends on the thin/thick disc populations has been studied for a subset of the CORALIE sample with homogeneous abundance determinations from previous studies. \\item No significant differences or peculiar enrichment in $\\alpha$-elements are observed in the abundance patterns of the subset of stars with homogeneous abundance determinations from previous studies. However, the small number of thick disc members prevents us from reaching conclusive results. \\item The tentative existence of thick disc stars at supersolar metallicities is confirmed by the appearance of one thick member over 21 with high metallicity (\\object{HD\\,152391} with [Fe/H] = 0.03). \\item The high average metallicity of the Hyades stream reported by previous works is confirmed. Our results are in agreement with the scenario of non-axysimmetric perturbations, such as transient spiral waves, proposed by several results of simulations and observations. According to this model, the Hyades stream would be the outward-moving stream of high-metallicity stars born at inner Galactic radii and then pushed into the solar neighborhood. We have also found a new clue confirming the solar and subsolar average metallicity suggested for the Sirius stream. \\item The group of planet-host stars shows a kinematic behaviour much more similar to the metal-rich comparison subsample, than to the whole comparison sample, in the sense that the overdensity in the Hyades stream observed for stars with planets is much higher than that observed for the whole comparison group. This issue has been reached by the two methods independently and strongly supports a primordial origin for the metal excess observed in stars with giant planets. \\item If we interpret our results within the scenario proposed to explain the presence and origin of dynamical streams in the solar neighbourhood, stars with giant planets could have formed more easily at inner Galactic radii in a more metal-rich intergalactic medium and then suffered radial displacements due to a non-axysimmetric component of the galactic potential, pushing them into the solar neighbourhood. \\end{itemize}  This paper presents the results of the first complete analysis on this topic. However, the uncertainties affecting our results are large, due to the intrinsic errors of the kinematics analysis, but also, and mainly, to the inevitably small number of targets with planets in the CORALIE sample. Consequently, substantial improvements will be possible in the near future, when the steadily increasing number of known stars with planetary companions will have risen enough to permit more definite conclusions. \\acknowledgement{This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France. We thank C. Carretero for many fruitful discussions and useful suggestions. Support from Funda\\c{c}\\~ao para a Ci\\^encia e a Tecnologia (Portugal) to N.C.S. in the form of a scholarship (reference SFRH/BPD/8116/2002) and a grant (reference POCI/CTE-AST/56453/2004) is gratefully acknowledged.}"
    },
    "0608/astro-ph0608343_arXiv.txt": {
        "abstract": "In its first part, this paper summarizes recent work on the mass and shape of the Galactic dark halo. The second part presents a review of the large-scale structure of the Milky Way, and of the evidence that the inner Galaxy is dominated by baryonic matter. This is briefly compared with the predictions of $\\Lambda$CDM and MOND. Finally, a summary is given of bulge formation from clumpy, gas-rich disks, a process which may give rise to old, disk-like, $\\alpha$-rich bulges similar to the Galactic bulge. ",
        "introduction": "Our position inside the Milky Way has both advantages and drawbacks. On the one hand, a number of experiments are possible that give complementary information on Galactic structure to that attainable for external galaxies. Notable examples are microlensing surveys, three-dimensional velocity data, local dynamical and stellar population measurements, determination of line-of-sight density distributions, etc. On the other hand, our view from inside the disk is often obscured, distance measurements are often needed but difficult, and deciphering the large-scale structure of the Galaxy is a much more involved process than from the image of an external galaxy. Both aspects are evident in the following sections: in \\S2, we review recent work on measuring the radial mass distribution and shape of the Galactic dark matter halo. Then in \\S3, we summarize constraints on the distribution of baryonic matter in the Milky Way, based on NIR, gas kinematic, and microlensing data. Finally in \\S4, we discuss a model of bulge formation from clumpy gas-rich disks which has some promise for understanding Galactic bulge observations. ",
        "conclusions": ""
    },
    "0608/astro-ph0608033_arXiv.txt": {
        "abstract": "Thermal X-ray emission which is simultaneous with the prompt gamma-rays has been detected for the first time from a supernova connected with a gamma-ray burst (GRB), namely GRB060218/SN2006aj. It has been interpreted as arising from the  breakout of a mildly relativistic, radiation-dominated shock from a dense stellar wind surrounding the progenitor star. There is also evidence for the presence of a mildly relativistic ejecta in GRB980425/SN1998bw, based on its X-ray and radio afterglow. Here we study the process of repeated bulk Compton scatterings of shock breakout thermal photons by the mildly relativistic ejecta. During the shock breakout process, a fraction of the thermal photons would be repeatedly scattered between the pre-shock material and the shocked material as well as the mildly relativistic ejecta  and, as a result,  the thermal photons get boosted to increasingly higher energies. This bulk motion Comptonization mechanism will produce nonthermal gamma-ray and X-ray flashes, which could account for the prompt gamma-ray burst emission in low-luminosity supernova-connected GRBs, such as GRB060218.  A Monte Carlo code has been developed to simulate this repeated scattering process, which confirms  that a significant fraction of the thermal photons get ``accelerated\" to form a nonthermal component, with a dominant luminosity. This interpretation for the prompt nonthermal emission of GRB060218 may imply that either the usual internal shock emission from highly relativistic jets in these low-luminosity GRBs is weak, or alternatively, that there are no highly relativistic jets in this peculiar class of bursts. ",
        "introduction": "Supernova shock breakout has been predicted for a few decades (e.g. Colgate 1974; Klein \\& Chevalier 1978; Ensman \\& Burrows 1992; Blinnikov et al. 1998, Matzner \\& McKee 1999). In core collapse supernovae, a shock wave is generated which propagates through the progenitor star and ejects the envelope. As the shock propagates through the envelope, it is mediated by radiation: The post shock energy density is dominated by radiation, and the shock transition is mediated by Compton scattering. As the shock approaches the edge of the star, the optical depth of the plasma lying ahead of the shock decreases. At the radius where the optical depth drops to $\\sim c/v_s$, where $v_s$ is the shock velocity, Compton scattering can no longer sustain the shock, and the shock undergoes a transition to a viscous shock (or a collisionless shock, Waxman \\& Loeb 2001). This transition is accompanied by a bright ultraviolet/X-ray burst, as photons escape far ahead of the shock due to the low optical depth. The term \"shock breakout\" is commonly used to refer to the emergence of the shock from the edge of the star. However, if the star is surrounded by an optically thick wind, the radiation mediated shock would continue to propagate into the wind, up to the point where the wind optical depth drops below $\\sim c/v_s$. In this case, shock breakout, i.e. the transition from a radiation mediated shock to a viscous (or collisionless) one accompanied by a bright ultraviolet/X-ray burst, occurs as the shock propagates through the wind, at a radius which may be significantly larger than the star's radius. We use here the term \"shock breakout\" to denote this transition in general, regardless of whether it occurs as the shock reaches the edge of the star or further out within an optically thick wind.  Shock breakout has previously never been directly detected from any supernovae due to its transient nature and its very early (minutes to hours) occurrence, in the absence of a suitably prompt trigger alert. Recently, thanks to its sensitive gamma-ray trigger and rapid slewing capability,  {\\it Swift} has detected early thermal X-rays emission from a supernova associated with a GRB, namely GRB060218/SN2006aj (e.g. Campana et al. 2006; Pian et al. 2006; Mazzali et al. 2006; Modjaz et al. 2006; Sollerman et al. 2006; Mirabel et al. 2006;  Cobb et al. 2006; Soderberg et al. 2006; Liang et al. 2006a), which has been interpreted as arising from the breakout of a radiation-dominated shock (Campana et al. 2006; Waxman et al. 2007).  The large energy carried by thermal X-rays, $\\sim10^{49}$~erg, and the long duration ($\\sim 1000$~s) of the thermal X-ray emission implies that shock breakout must occur at a radius of $\\sim10^{13}$~cm. The lack of Hydrogen lines implies, on the other hand, a compact progenitor, a WR star of radius $\\sim10^{11}$~cm. This suggests that the progenitor was surrounded by an optically thick wind, and that shock breakout occurred at a radius $\\sim10^{13}$~cm, where the optical depth of the wind becomes small enough. The inferred mass loss rate, $\\sim3\\times10^{-4}M_\\odot/{\\rm yr}$  (Campana et al. 2006) is consistent with that typically expected for WR stars (e.g. Felli \\& Panagia 1982; Nugis \\& Lamers 2000, 2002)\\footnote{The arguments based on radio observations against the presence of an optically thick wind (Fan et al. 2006), are not conclusive, since radio observations do not allow one to determine the explosion parameters. In particular, the estimates for the kinetic energy range between $\\sim10^{48}$~erg and $\\sim10^{50}$~erg, and the ambient medium density estimates range between $\\sim10^0{\\rm cm}^{-3}$ and $\\sim10^2{\\rm cm}^{-3}$ (Soderberg et al. 2006; Fan et al 2006). }. Another possibility is that the shock breaks out from an optically thick shell pre-ejected from the progenitor (the required shell mass is only $10^{-7} M_\\sun$, see Campana et al. 2006).  The simple result, that a wind photospheric radius $R_{ph}\\sim10^{13}$~cm is required to account for the thermal emission as due to shock break out from a wind, is supported by the calculations of Li (2007), who finds that the radius at which the wind optical depth equals 20 is required to be $R_*\\sim100R_\\odot$. Since the wind models investigated in Li (2007) have $1.5\\lesssim R_{ph}/R_*<3$ (compare his figures 12 and 3), his conclusion implies that $R_{ph}\\sim10^{13}$~cm is required, consistent with our conclusion. Note, however, that the claim that $R_*\\sim100R_\\odot$ is required is not robust, since the ratio $R_{ph}/R_*$ depends on the adopted wind velocity profile, and is likely to be $>3$ for WR stars. { Note also that when the mass loss is well above $10^{-4} M_\\odot {\\rm yr^{-1} }$, the wind becomes very optically thick and the apparent stellar radius grows, so that the star no longer looks like a WR star (Nino Panagia, personal communication)}. It should also be pointed out that the analysis of the required wind properties in Li (2007) is far too restricted. A model is adopted by Li (2007) where the shock is produced by a spherical explosion with fixed energy ($2\\times10^{51}$~erg) and fixed ejecta mass ($2M_\\odot$), and its acceleration is due to propagation through the model wind profile. Beyond the fact that dependence on the explosion parameters is not explored, it is not at all clear that this is the correct description of shock dynamics.  In fact, it is known that such shock acceleration models (Matzner \\& McKee 1999; Tan et al. 2001) can not account for the energy deposited in the mildly relativistic ejecta, $v/c=0.8$, observed in GRB 980425/SN1998bw (Kulkarni et al. 1998, Waxman \\& Loeb 1999, Chevalier \\& Li 1999, Waxman 2004). It could be, e.g., that this mildly-relativistic shock is driven instead by a choked jet propagating through the progenitor.  From the thermal energy density of the shock, Campana et al. (2006) have inferred that the shell driving the radiation-dominated shock in GRB060218/SN2006aj must be mildly relativistic, with a velocity $\\Gamma\\beta\\sim1-2$, where $\\beta$ is the velocity in units of the speed of light $c$ and $\\Gamma=(1-\\beta^2)^{-1/2}$ is the Lorentz factor. This shock could be driven by the outermost parts of the envelope that get accelerated to a mildly relativistic velocity when the supernova shock accelerates in the density gradient of the envelope of the supernova progenitor (Colgate 1974; Matzner \\& McKee 1999; Tan et al. 2001), probably a Wolf-Rayet (WR)  star. The more compact configuration of a WR star, compared with a blue or red supergiant, and the large explosion energy characterizing a hypernovae may help to accelerate a considerable amount of matter to a mildly relativistic, and probably anisotropic velocity distribution (see however the comment above on SN1998bw).  Due to the connection with a GRB, there are also other possibilities for the origin of this radiation-dominated shock. In the collapsar model for GRBs, the shock could be driven by the cocoon of a relativistic jet (M\\'esz\\'aros \\& Rees 2001; Ramirez-Ruiz et al. 2002; Zhang et al. 2003) or the slower ($\\Gamma\\beta\\sim1-2$), high-latitude outer parts of a relativistic jet. The emergence of the shock in this case would be highly non-spherical. As pointed out in Camapana et al. (2006), observations require a non-spherical shock breakout. This does not necessarily imply a shock driven by a cocoon/jet, since the wind itself may be highly non-spherical\\footnote{We note in this context that the claim in Ghisellini et al. (2006b), that the optical emission and the thermal X-ray emission are not due to the same thermal component is not necessarily valid, since Ghisellini et al. (2006b) assume spherical symmetry and neglect in addition light travel time effects.}. Observations of high optical polarization in SNe at early times of 3-5 days (e.g. Gorosabel et al. 2006) are also consistent with non-spherical shocks.  The  relatively flat light curve of the X-ray afterglow of the nearest GRB, GRB980425/SN1998bw, up to $\\sim 100$ days after the burst, has been argued to be caused by the coasting phase of a mildly relativistic shell with energy of a few times $10^{49}\\rm erg$ (Waxman 2004). Similarly, the X-ray afterglow of another nearby GRB, GRB031203/SN2003lw also had an early flat light curve (Watson et al. 2004). Thus, it is possible that a mildly relativistic ejecta is ubiquitous in supernovae-associated GRBs. In this paper, we suggest that the nonthermal X-ray  and gamma-ray emission in such objects  can arise from repeated bulk-Compton scatterings (i.e. bulk motion Comptonization) of thermal {shock breakout X-rays photons.}    Blandford \\& Payne (1981) first noted the importance of bulk motion acceleration of photons in a radiation-dominated shock. They found that photons are preferentially upscattered by the bulk motion rather than by the thermal motions of the electrons, and a power-law spectrum extending to high energies forms, when the electron thermal velocity is less than the shock velocity $v_s$. Repeated scatterings using the energy of the bulk motions of two approaching relativistic shells in the context of GRB internal shocks was studied by Gruzinov \\& \\Meszaros (2000). They found that the seed synchrotron photons can be boosted to much higher energies, which is confirmed by their Monte Carlo simulations. This process is equivalent to the Fermi acceleration mechanism of particles (Blandford \\&  Eichler 1987)  or  photon scattering off Alvf\\'{e}n waves (Thompson 1994), but here the mechanism, instead, uses the relative bulk motion and accelerates photons. { The bulk motion  Comptonization mechanism has also been invoked to interpret the  high-energy spectra of  accreting black holes (e.g. Shrader \\& Titarchuk 1998; Titarchuk \\& Shrader 2005).} ",
        "conclusions": "We have discussed a model for the early gamma--ray and X-ray emission of low luminosity GRBs which are associated with supernovae, based on bulk Comptonization of thermal photons by a semi-relativistic ejecta. One expects the following basic characteristics for the gamma-ray bursts produced through the above mechanism:    1) Smooth light curve profiles. The production of the gamma-rays is through the repeated bulk-Compton scattering of thermal photons emitted by the radiation dominated shock. The light curve should follow the time behavior of the shock breakout and should therefore be smooth. The light curves would generally have a simple profile without multi-peak structure. The three nearby low-luminosity bursts, GRB060218, GRB980425 and GRB031203, all have smooth light curves, consistent with this picture.  Given that GRB980425 and GRB031203 are also believed to have produced mildly relativistic ejecta, it is possible that the gamma-ray emission is due to supernova shock breakout also in these two bursts. { Based on the association between GRB980425 and SN1998bw, Bloom et al. (1998) speculated that there may exist a subclass of smooth light curve GRBs produced by SNe.} The duration of GRB980425 and GRB031203 is only tens of seconds, which can be interpreted as due to shock break out from the progenitor envelope, implying that the wind surrounding the progenitor is optically thin. The characteristic variability time of GRB980425 and GRB031203 is a few seconds, corresponding to $R(\\tau=1)\\sim10^{11}\\rm cm$, which is comparable to the stellar radius of a Wolf-Rayet star. The speculation that the stellar wind surrounding  GRB980425 is optically thin is consistent with the wind mass-loss rate inferred from the X-ray and radio afterglow of GRB980425/SN1998bw (Waxman 2004; Li \\& Chevalier 1999). The longer timescale of GRB 060218 can, on the other hand, be interpreted in terms of the shock breaking out from an optically thick wind.  \\begin{table*} \\caption{The spectrum of three nearby low-luminosity GRBs} \\begin{center} \\begin{tabular}{lllllllll} \\hline \\hline GRB/SN & z & $E_{\\gamma, {\\rm iso}}($\\rm erg$) $ &  $\\alpha$& $\\varepsilon_c (\\rm KeV)$ &  references$^*$\\\\ \\hline GRB980425/SN1998bw & 0.0085&$(8.5\\pm0.1)\\times10^{47}$ &$0.45\\pm0.22$ & $\\sim 200$ &1;3 \\\\ GRB031203/SN2003lw & 0.105 & $(4\\pm1)\\times10^{49}$ &$0.63\\pm0.06$ &$>190$ & 2;4\\\\ GRB060218/SN2006aj & 0.0331& $(6.2\\pm0.3)\\times10^{49}$ &0.45 & $\\sim 30$$^\\S$& 1;3;5; 6\\\\ \\hline \\end{tabular} \\end{center}  $^{*}$The redshift $z$, spectral index $\\alpha$ ($F_\\nu\\propto \\nu^{-\\alpha}$) and cutoff energy $\\varepsilon_c$ of the three nearby, sub-luminous SN-connected GRBs. References: (1) Ghisellini et al. (2006); (2) Sazonov et al. 2004; (3)Kaneko et al. 2006; (4) Sazonov et al. 2004; (5) Campana et al. 2006; (6) S. Campana, private communication. \\\\ $^\\S$ For the time $<300$ seconds after the trigger. \\end{table*}  2)The spectrum is expected to be composed of a simple power-law, with a high energy cutoff lower than few $\\times100$~keV, and a thermal X-ray component. The spectrum of the gamma-ray emission can be modelled as a cutoff power law rather than the Band function for the usual bursts. The spectrum is expected to evolve from hard to soft since the spectral index becomes smaller and the cutoff energy decreases as the optical depth decreases. The spectra of three nearby SN-GRBs, indicated in Table 1, are consistent with this cutoff power-law spectrum.    Cosmological GRBs with isotropic equivalent energes $E\\sim10^{51}-10^{54} {\\rm erg}$ are generally believed to be produced through internal shocks in relativistic jets with Lorentz factor $\\Gamma\\ga 100$ (for a recent review, see \\Meszaros 2006). In the collapsar scenario, relativistic jets can break free from the star  along the rotation axis of a collapsing stellar core, provided that the central engine feeding time of the jets is sufficient long. Light curve breaks in the afterglow emission attributed to jet effects have been seen in many bursts. On the other hand, three nearby GRBs, GRB980425/SN1998bw, GRB031203/SN2003lw and GRB060218/SN2006aj, have isotropic energies in the range $\\sim10^{48}-10^{49}{\\rm erg}$, much lower than typical cosmological bursts. Due to their proximity, the inferred intrinsic rate of these sub-energetic events is, however, much higher (Soderberg et al. 2006; Liang et al. 2006b), raising difficulties in interpreting these events as typical GRBs observed off-axis (Cobb et al. 2006). Both the properties of the prompt gamma-ray emission and of the afterglows are different from those of typical cosmological bursts. Up to now there is no straightforward evidence for the presence of highly relativistic jets in these sub-energetic bursts, but instead, there does exist evidence for mildly relativistic ejecta in these bursts. Thus, it could be that the relativistic jets in these events are much weaker, giving most of their energy to the outer slow material when they are burrowing their way though the star, which leads to a much wider anisotropic pattern of outflow. This could be the case even if the jets are chocked before emerging outside the stellar envelope, due e.g. to an insufficient  central engine feeding time. Whether there is a relativistic jet or not can be probed through sub-GeV to GeV observations during the shock breakout (Wang \\& \\Meszaros 2006)    Mildly relativistic ejecta may also exist in the usual high-luminosity long GRBs, since  shock acceleration is expected to accompany the supernova. This might explain the low-velocity component, other than the relativistic jet, found by Berger et al. (2003) and Kartik et al. (2003) in another nearby supernova-connected GRB, namely GRB030329/SN2003dh, from the radio and optical afterglow more than 1.5 days after the explosion. In this burst, the gamma-ray emission from the mildly relativistic ejecta is lower than that of the relativistic jet, since the latter has a much higher isotropic energy. As suggested by Berger et al (2003), the total energy in the low velocity component and the relativistic jet may be a roughly constant quantity, the only variable being the ratio of these two components. In this vein, we would predict that at high redshifts there could be many GRB/SN with stronger lower velocity and weaker relativistic jet components, which would be harder to detect, even though their rates of occurrence might exceed that of the more easily detectable bursts with weak low-velocity and stronger relativistic jets.    {"
    },
    "0608/astro-ph0608205_arXiv.txt": {
        "abstract": "A sequence of three {\\it Chandra} X-ray Observatory High Resolution Camera images taken over a span of five years reveals  arc-second-scale displacement of RX~J0822--4300, the stellar remnant (presumably a neutron star) near the center of the Puppis A supernova remnant. We measure its proper motion to be $0.165 \\pm 0.025$\\ arcsec yr$^{-1}$\\ toward the west-southwest. At a distance of 2 kpc, this corresponds to a transverse space velocity of $\\sim$1600 km s$^{-1}$. The space velocity is consistent with the explosion center inferred from proper motions of the oxygen-rich optical filaments, and confirms the idea that Puppis A resulted from an asymmetric explosion accompanied by a kick that imparted roughly 3$\\times$10$^{49}$ ergs of kinetic energy (some 3 percent of the kinetic energy for a typical supernova) to the stellar remnant. We discuss constraints on core-collapse supernova models that have been proposed to explain neutron star kick velocities. ",
        "introduction": "There has long been broad consensus that core-collapse supernovae---the explosions of massive progenitors that produce Types II, Ib, and Ic events at least---leave behind a compact stellar remnant: either a neutron star or a black hole. Early on, this model was marred by the paucity of observed compact objects associated with supernova remnants (SNRs).  The discovery in recent years of numerous compact X-ray sources associated with SNRs, especially with the oxygen-rich SNRs that are presumably the young remnants of core-collapse SNe, has removed this blemish \\citep[e.g.][for a recent review]{manchester01}. Compact stellar remnants have been identified near the centers of all three of the known oxygen-rich SNRs in the Galaxy---Cas A \\citep{tananbaum99, chakrabarty01}, Puppis A \\citep{petre96}, and most recently G292.0+1.8 \\citep{hughes01, camilo02}---plus 0540--69.3 in the LMC \\citep{seward84}.  There is also growing evidence, both observational and theoretical, that the explosions of core-collapse SNe are asymmetric. In young, oxygen-rich SNRs, whose fast filaments provide the best opportunities for observing ejecta from the cores of massive SNe, anisotropies appear to be typical:  e.g., Cas A  \\citep{lawrence95, reed95, fesen01, hwang00, hwang04}, and the LMC remnants N132D  \\citep{morse95}, and 0540--69.3 \\citep{kirshner89}, in addition to Puppis A \\citep{winkler85}.  Furthermore, two-dimensional models for core-collapse SNe show instabilities just prior to collapse that produce asymmetries in the ejecta distribution \\citep[e.g.,][]{burrows95}.  Simple conservation of momentum requires that if the explosion of a progenitor star expels the ejecta preferentially in one direction, the compact core must recoil in the opposite direction.  By measuring the displacements of young pulsars from the apparent centers of their associated SNR shells and using the pulsar spin-down periods as age estimates, \\citet{caraveo93} and \\citet{frail94} inferred that pulsars are typically born with transverse velocities of $500 \\kms$, and that velocities $\\gtrsim 2000 \\kms$\\ may occur. Direct radio and optical measurements have found transverse velocities ranging from $60\\kms $\\  \\citep[for Vela -- ][]{dodson03} to $240 \\kms$\\ \\citep[for PSR B1951+32 in CTB~80 --][]{migliazzo02}.  Here we report measurement of the recoil of RX~J0822--4300, the presumed neutron star near the center of Puppis A, obtained through images taken at three epochs spanning 5 years from the High Resolution Camera (HRC) on the {\\it Chandra} X-ray Observatory. Past optical studies of the most prominent ejecta-dominated filaments show that they are concentrated in the northeast quadrant of the remnant, and that their motions are northward and eastward, consistent with undecelerated expansion from a common center \\citep{winkler85, winkler88}.   The present measurement of the stellar remnant's rapid motion to the west-southwest completes a picture of asymmetric ejecta and neutron star recoil resulting from a core-collapse supernova.  As the original version of this  paper was nearing completion, we became aware of a paper by \\citet[][henceforth HB06]{hui06a} that presents a similar analysis based on two of the three {\\it Chandra} observations used here.  While the HB06 result is qualitatively similar to our own, our analysis leads to a more precise measurement and a significantly higher velocity for RX~J0822--4300.  We compare their methods with our own in \\S 3.5. ",
        "conclusions": "As a summary, we may highlight the following points: (1) A precise position for RX~J0822--4300, the presumed neutron star, has been measured from three independent \\chandra\\ HRC observations spread over a baseline of over five years. Comparison  among these yields a proper motion of $0\\farcs 165 \\pm 0\\farcs 025\\; {\\rm yr}^{-1}$\\ at  position angle $248\\degr \\pm \\, 14\\degr$. The uncertainty is dominated by the precision of position fits for two fiducial stellar X-ray sources used in the analysis.  (2) This motion implies a transverse velocity of $1570\\pm 240\\,(d/2\\; {\\rm kpc}) \\kms$\\ toward the west-southwest.  This direction is roughly opposite to the motion of a handful of oxygen-rich optical filaments---presumably near-undiluted ejecta from the outer core of the progenitor---that are scattered throughout the northeast quadrant of Puppis A\\@.  Furthermore, backwards extrapolation of the motion of RX~J0822-4300 indicates an origin consistent with the measured expansion center for the O-rich filaments.    This completes---in at least one SNR---what is becoming almost a  canonical picture for core-collapse supernovae:  an asymmetric supernova explosion accompanied by recoil of the compact stellar remnant, here almost surely a neutron star.  (3) The kinetic energy associated with the transverse motion of RX~J0822-4300 is $\\sim 3\\times 10^{49}\\; (d/2\\; {\\rm kpc})^2\\; {\\rm ergs}$, only about 3\\% \\ of the total of $\\sim 10^{51} \\; {\\rm ergs}$\\ expected in a typical supernova.   Some 2--3 dozen O-rich knots like those now glowing optically are sufficient to balance the momentum of the neutron star.  (4) The physics of the explosion mechanism necessary to produce such a fast neutron star remains elusive, but the high kick velocity and lack of apparent pulsations from RX~J0822-4300 do constrain possible models.   Both these observations argue against  neutrino/magnetic-field driven or electromagentically driven mechanisms for the kick. The most likely candidate appears to be some mechanism through which hydrodynamic instabilities in the explosion lead to recoil of the compact remnant. However, the most specific such model, that of \\citet{burrows07}, is strained to explain both the high kick velocity and the apparent absence of iron-rich ejecta from the inner core of the Puppis A progenitor.   A complete kinematic study  of Puppis A, including both the oxygen knots and the stellar remnant, will be interesting and should be carried out.   CCD data should enable measurement of the motions for significantly more oxygen knots than the handful used by \\citet{winkler88}, who based their study on only 11 individual knots whose motions could be measured on photographic plates.   A third-epoch \\chandra\\ observation with the HRC-I and/or HRC-S would further cement the kinematic picture."
    },
    "0608/hep-th0608068_arXiv.txt": {
        "abstract": "A class of Einstein-dilaton-axion models is found for which almost all  flat expanding homogeneous and isotropic universes undergo recurrent periods of acceleration.  We also extend recent results on eternally accelerating open universes. ",
        "introduction": "\\setcounter{equation}{0}  The discovery that the expansion of our Universe is accelerating, presumably due to ``dark energy'', together with the compatibility of recent astronomical observations with the hypothesis of an earlier inflationary epoch, has given a new impetus to studies of cosmologies driven by scalar fields. Models with a single scalar field have a long history. In more recent times, models with two or more scalars have attracted much attention; in particular models with both dilaton and axion \\cite{Linde:1991km}. Of most  relevance here is the work of  Billyard et al. \\cite{Billyard:2000cz} because the models we study  include one  studied by them, and our methods are similar.  We suppose that a dilaton field $\\sigma$ and axion field $\\chi$, which together define a map from a $D$-dimensional spacetime to a hyperbolic space with an $Sl(2;\\bR)$-invariant metric, are coupled to gravity. The $Sl(2;\\bR)$-invariance must be broken by any non-constant scalar field potential, but if we require only invariance under dilatations then an exponential potential for the dilaton is allowed. These considerations lead to a low-energy Lagrangian density of the form \\be\\label{laginitial} {\\cal L} = \\sqrt{-\\det g}\\left[ R - {1\\over2}\\left(\\partial\\sigma\\right)^2 - {1\\over2} e^{\\mu\\sigma}\\left(\\partial\\chi\\right)^2 - \\Lambda e^{-\\lambda\\sigma} \\right] \\ee where $\\mu$ and $\\lambda$ are constants, as is the `cosmological constant' $\\Lambda$, which we here assume to be positive. We may choose $\\lambda\\ge0$ without loss of generality but then the sign of $\\mu$ is significant, so we allow for either sign. Note that $|\\mu|$ is proportional to the radius of the hyperbolic target space. When $\\mu=0$, the target space is flat and the axion field decouples; the resulting cosmologies were studied in \\cite{Bergshoeff:2003vb,Townsend:2004zp,Collinucci:2004iw} so we assume here that $\\mu\\ne0$. Models of this type can be motivated in various ways. For example the Freedman-Schwarz $D=4$ supergravity theory \\cite{Freedman:1978ra} has a consistent truncation to (\\ref{laginitial}) with $2\\lambda=-\\mu =2$.  We shall study flat FLRW cosmological solutions of the above model as a function of the parameters $\\mu$ and $\\lambda$. The $D=4$ model with $\\lambda=-\\mu=2$ was previously studied by Billyard et al. \\cite{Billyard:2000cz} with a string-theory motivation that led them to consider a `string-frame' metric rather than the Einstein-frame metric implicit in (\\ref{laginitial}). They found an interesting quasi-cyclic evolution of flat cosmological models that they interpreted as successive periods of expansion and contraction. As we point out here, the Einstein-frame universe is either eternally expanding or eternally contracting, but each undergoes successive periods of {\\it acceleration}. The `string-frame'  cosmologies of \\cite{Billyard:2000cz} also undergo successive periods of acceleration because  this is an immediate consequence of  re-expansion following a period of  contraction, but acceleration in non-Einstein frames does not  imply acceleration in the Einstein frame; in particular, a flat Einstein-frame universe cannot smoothly pass from  contraction to expansion (see e.g. \\cite{Khoury:2001bz}) so this source of cosmic acceleration is not available in the Einstein frame. We shall show that, in Einstein frame, almost all flat expanding FLRW cosmologies undergo quasi-cyclic alternation between periods of  acceleration and deceleration provided that both $\\lambda$ and $-\\mu$ are  sufficiently large. The late-time behaviour is a power-law expansion which may be accelerating or decelerating, depending on the value of $\\mu/\\lambda$.   Of  course, there is no physical significance to the choice of frame within the models we consider  but this choice  becomes significant once one considers the additional fields needed for a realistic theory.  In this larger context, the choice of Einstein frame is equivalent to the assumption that  the energy-momentum-stress tensor of the matter fields responsible for structure couple to the Einstein-frame metric, in the sense that the motion of a free particle associated with such a field is geodesic in this frame, an assumption that  is consistent with observations.  The model defined by (\\ref{laginitial}) can be consistently truncated by taking $\\chi=\\chi_0$, for any constant $\\chi_0$.  Solutions of the resulting Einstein-dilaton  model are therefore also solutions of the Einstein-dilaton-axion model, and this  provides a useful check on our results. The Einstein-dilaton model has been much studied in the past, but some features in relation to eternal cosmic acceleration of open FLRW universes in models arising from hyperbolic compactification have come to light only recently. For the values of $\\lambda$ that arise in such compactifications,  Chen et al. \\cite{Chen:2003dc} have shown that  late-time `eternal'  acceleration is possible for open universes (see also \\cite{Neupane:2003cs}).  Andersson and Heinzle have further shown that `eternal' acceleration is possible in the stronger sense  of  ``at all times''  (and not merely in the future of some initial onset of  acceleration), and that there is a unique open universe cosmology with this property \\cite{Andersson:2006du}. Here we extend this result, in the  form of a conjecture, to other values of $\\lambda$, including those that arise from flux compactifications. ",
        "conclusions": "\\setcounter{equation}{0}    We have presented a study of the homogeneous and isotropic cosmological models that arise in a two-parameter class of models of gravity coupled to dilaton and axion fields in arbitrary spacetime dimension $D$; the two parameters $(\\lambda,\\mu)$ determine the dilaton coupling constant (and hence the `strength' of an `exponential' potential) and the radius of curvature of the 2-dimensional hyperbolic target space (which we assumed to be finite).  It is well-known that all  possible cosmological solutions of  the zero-axion model can be understood in terms of trajectories of a two-dimensional autonomous dynamical system, but we have explored some new features that arise from consideration of the global nature of the phase space of this system.  It is less well-known that a similar reduction to a two-dimensional autonomous dynamical system is possible for the generic dilaton-axion model provided that the restriction is made to {\\it flat}  cosmologies; a particular $D=4$ dilaton-axion model was analysed by this method in \\cite{Billyard:2000cz} and we have extended this analysis to the generic dilaton-axion model. In addition, we have explored here the consequences for  {\\it Einstein frame} cosmology, which are quite different from those in the `string-inspired' frame of \\cite{Billyard:2000cz}. In particular, all flat Einstein-frame cosmologies are either eternally expanding or eternally contracting, but may undergo one or more periods of cosmic acceleration.  Our main result is that {\\it there is a large range of  the parameters $(\\lambda,\\mu)$ for which almost all flat expanding universes undergo an epoch of oscillation between acceleration and deceleration}, with either continuous late-time acceleration or continuous late-time deceleration depending on the value of the ratio $\\lambda/\\mu$ (and eternal oscillation at a critical value of this ratio).  The early-time behaviour depends strongly on whether $\\lambda$ is greater or less than the `hypercritical' value $\\lambda_h$. For $\\lambda>\\lambda_h$, the global phase space is topologically equivalent to the one given in \\cite{Billyard:2000cz}, in which almost all trajectories are asymptotic in one direction to a heteroclinic cycle and in the other direction to a focus within this cycle. The focus can be stable or unstable according to the chosen direction of time. In a non-Einstein frame, such as that considered in  \\cite{Billyard:2000cz}, for which the universe undergoes quasi-cyclic periods  of expansion and contraction, one may suppose that the focus is unstable and the the heteroclinic cycle is approached at late times. In the Einstein frame, for which the universe is either eternally expanding or eternally contracting, the heteroclinic cycle is approached at late times only by a contracting universe. If we assume expansion, then the heteroclinic cycle dominates the early-time behaviour, which exhibits oscillation between acceleration and deceleration. This oscillation continues until the focus is approached sufficiently closely that its behaviour, of continuous acceleration or continuous deceleration, takes over. This is the origin of the epoch of oscillation between cosmic acceleration and deceleration when $\\lambda>\\lambda_h$. For $\\lambda<\\lambda_h$ there is no heteroclinic cycle and the early-time behaviour of expanding universes is quite different, but the medium-time behaviour may again exhibit oscillation between cosmic acceleration and deceleration for values of the parameters that put the focus sufficiently close to the borderline between acceleration and deceleration.  The fact that there are models for which almost all {\\it flat} universes undergo cycles of cosmic acceleration and deceleration is of little use for inflationary purposes because one aim of inflation is to explain why the Universe is flat, and for that purpose one needs to consider universes that are not flat. In this case, however, the cosmological evolution cannot be reduced to the study of a 2-dimensional phase-plane. It would be of interest to use other methods to learn more about generic non-flat cosmologies (with non-constant axion)  but at present we must assume that there exist some non-flat cosmologies that become approximately flat at late-times, in which case their late-time behaviour  will determine the  `initial' conditions for a flat cosmology. Thus, the early-time behaviour of the flat cosmologies studied in this paper is probably not physically relevant. The late-time behaviour may not be relevant either, depending on whether an almost flat universe becomes more or less flat as time progresses. It is plausible that flatness is a stable feature of universes that exhibit continuous late-time acceleration, and an unstable feature of those that exhibit continuous late-time deceleration. However, the medium-time recurrent acceleration that we have found to be typical of many dilaton-axion models could be relevant to the observed cosmic acceleratiion of our Universe. At the very least, it shows that current observations of cosmic acceleration provide no evidence that this acceleration will continue forever.  With many scalar fields, the future is uncertain!"
    },
    "0608/astro-ph0608563_arXiv.txt": {
        "abstract": "We report 24 and/or 70 \\um measurements of $\\sim$160 A-type main-sequence stars using the Multiband Imaging Photometer for {\\it Spitzer} (MIPS). Their ages range from 5 to 850 Myr based on estimates from the literature (cluster or moving group associations) or from the H-R diagram and isochrones. The thermal infrared excess is identified by comparing the deviation ($\\sim$3\\% and $\\sim$15\\% at the 1-$\\sigma$ level at 24 and 70 \\mm, respectively) between the measurements and the synthetic Kurucz photospheric predictions.  Stars showing excess infrared emission due to strong emission lines or extended nebulosity seen at 24 \\um are excluded from our sample; therefore, the remaining infrared excesses are likely to arise from circumstellar debris disks.  At the 3-$\\sigma$ confidence level, the excess rate at 24 and 70 \\um is 32\\% and $\\ge$33\\% (with an uncertainty of 5\\%), considerably higher than has been found for old solar analogs and M dwarfs. Our measurements place constraints on the fractional dust luminosities and temperatures in the disks.  We find that older stars tend to have lower fractional dust luminosity than younger ones. While the fractional luminosity from the excess infrared emission follows a general 1/$t$ relationship, the values at a given stellar age vary by at least two orders of magnitude. We also find that (1) older stars possess a narrow range of temperature distribution peaking at colder temperatures, and (2) the disk emission at 70 \\um persists longer than that at 24 \\mm. Both results suggest that the debris-disk clearing process is more effective in the inner regions. ",
        "introduction": "From various lines of evidence (e.g., theoretical modeling, the cratering record on the Moon, returned Lunar samples, and isotopic data) it is believed that embryo terrestrial planets in the Solar System formed in $\\sim$10-30 Myr, evolved through a period of potentially immense collisions, and then underwent a reduced, but still significant collisional period that ended with another violent episode -- the Late Heavy Bombardment -- at about 700 Myr \\citep{kleine02,chambers04,gomes05,strom05}. Thereafter, the system was in a relatively settled state, setting the stage for life to emerge on Earth. We are unsure whether this sequence was typical or exceptional. We are not even sure of the details of the steps because much of the evidence has been obliterated over time. Hard evidence must often be supplemented with relatively poorly tested theory to assemble a complete picture. The obvious solution would be to observe these processes as they are occurring in other planetary systems. However, observing terrestrial planets around nearby stars passing through parallel stages of evolution is virtually impossible. The planets are too dim for detection against the glare of the central stars; they are not sufficiently massive to be detected through gravitational recoil; and relatively few young stars lie close to the Sun.    However, there is a promising alternative approach, to study planetary debris disks. Debris disks arise from collisions of asteroidal (or planetesimal) bodies that lead to cascades of collisions among the resulting debris. Eventually, significant amounts of this material are ground down to dust grains. Because the surface area per unit mass is large for dusty material, when these grains are heated by the central star they can produce a readily detectable level of excess emission in the mid- and far-infrared. The behavior of these infrared excesses can trace the different zones within a planetary system. For grains around an A-type star and large enough to be in pseudo-stable orbits, the mid-IR band ({\\it Spitzer} 24 \\mm, {\\it IRAS} and {\\it ISO} 25 \\um passbands) is sensitive to material largely between $\\le$5 and $\\sim$50 AU, while the far-IR band ({\\it IRAS} and {\\it ISO} 60 \\um and {\\it Spitzer} 70 \\um passbands) is sensitive to material between $\\sim$50 and 200 AU. Detailed studies of the behavior with age of debris disks in these two bands can indicate how the stages deduced for the early evolution of the Solar System are playing out in hundreds of other planetary systems.   Pioneering studies with {\\it IRAS} and {\\it ISO} suggest a systematic drop in infrared excess with stellar age (e.g., \\citealt{habing01,spangler01}), qualitatively similar to the drop that would be deduced from the settling down of the Solar System. However, because of limitations in the sensitivity and the accuracy of measurements with these two missions, they were unable to show unambiguously how debris disks evolve. {\\it Spitzer} brings significant advances in both sensitivity and photometric accuracy, increasing our ability to detect low levels of infrared excesses around hundreds of stars. A-type stars provide an ideal laboratory to study the early stages in planetary system evolution, particularly in the zones relevant for evolution of terrestrial planets (Earth-like temperatures occur at about 5 AU from such stars). Their main sequence lifetimes are long enough ($\\sim$800 Myr) to encompass the entire period of interest, they are of high enough luminosity to light up their debris well, they are sufficiently abundant and bright that many can be observed readily, and they are cool enough that they are unlikely to create infrared excesses in the form of ionized gas. \\citet{rieke05} have already conducted a study of the 24/25 \\um excesses in such stars, based on a combination of {\\it IRAS} and Multiband Imaging Photometer for {\\it Spitzer} (MIPS) measurements of stars in the field and MIPS measurements of stars in clusters. At that time, relatively few field stars had been measured well with {\\it Spitzer} at 70 \\mm. In this paper, we report MIPS photometry at 24 and 70 \\um of 160 main-sequence early type stars, enough to give a much better understanding of the evolution of debris systems. ",
        "conclusions": "About 160 A-type main-sequence stars with ages ranging from 5 to 850 Myr were measured at 24 and/or 70 \\um using the Multiband Imaging Photometer for {\\it Spitzer}. With {\\it Spitzer's} unprecedented sensitivity, we are able to identify infrared excesses of $\\sim$6\\% at 24 \\um and $\\sim$55\\% at 70 \\um above the photospheres at the 3-$\\sigma$ levels. At this 3-$\\sigma$ confidence level, the infrared excess rate is 32$\\pm$5\\% and $\\ge$33$\\pm$5\\% for 24 and 70 \\mm, respectively. The excess detection rate around these early-type stars is significantly higher than the one found in the old FGK \\citep{bryden06} and M-type stars \\citep{gautier06}. However, recent studies by \\citet{gorlova06} and \\citet{siegler06} show that the excess rate for young FGK stars is also higher than for old FGK stars; most importantly, the evolution in the trend of excess rate vs.~age looks similar between A stars and FGK stars (only the absolute fraction is different) in their studies.  It is difficult to determine whether there is an intrinsic difference in the infrared excess rates between A-type and FGK stars because equivalent levels of fractional luminosity become increasingly hard to detect as the star becomes cooler. The incidence of debris disks is actually the convolution of how many stars actually possess debris disks with given properties, and how detectable the disks are with {\\it Spitzer}, and this convolution is a strong function of spectral type. Thus debris-disk excess emission is not a steep function of stellar type. Since the detectability of a given level of fractional excess decreases with decreasing stellar temperature, the incidence of excesses is consistent with a similar debris disk driven fractional excesses between A and K stars. Nevertheless, it appears that age is the most important factor in determining the detectability of infrared excess among A- to K-type main-sequence stars, not stellar mass.    The amount of excess emission decreases with stellar age and follows a simple $t_o/t$ relationship in general with $t_o \\sim$150 Myr for excesses at 24 \\um but $t_o \\ge$400 Myr at 70 \\mm; that is, 24 \\um excesses decay more rapidly than those at 70 \\mm. In addition, at any given age there exists a large variety in the amount of 24 and 70 \\um excess emission. The observed [24]$-$[70] color temperatures for a total of 44 debris disks with 24 and/or 70 \\um excesses range from $\\sim$60 to $\\sim$250 K, but the majority of the disks have a typical temperature of $\\sim$ 90 K. Furthermore, older stars (age $>$ 300 Myr) tend to have colder observed [24]$-$[70] color temperatures than young ones. Assuming the dust we detect is blackbody-like, the decline of the [24]$-$[70] color implies the debris is located further away from the stars in the older systems.  The fractional luminosity for these 44 early-type debris disks ranges from $\\sim$10$^{-3}$ to $\\sim$10$^{-6}$. The trend between the fractional luminosity and stellar age follows a general 1/$t$ relationship; older stars tend to have lower fractional luminosity than younger ones. An upper envelope of 1/$t$ is seen in the fractional luminosity vs.~age diagram, suggesting that collisions followed by radiation pressure blowing out small grains are the dominant process to remove grains in these systems. The decreases of observed [24]$-$[70] color temperatures and the fractional luminosity suggest that the debris disk clearing is an inside-out process."
    },
    "0608/astro-ph0608049_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "Grain alignment is a subject that is very rich in terms of physical processes. The advances in understanding of grain alignment processes made the theory predictive and allowed to explain the observational data available. This enables one to reliably interpret observed polarization in terms of the underlying magnetic fields.     \\subsection*"
    },
    "0608/astro-ph0608613_arXiv.txt": {
        "abstract": "We have recently shown that both the Prugniel-Simien model and S\\'ersic's function (hereafter referred to as the Einasto model when applied to internal density profiles) describe simulated dark matter halos better than an NFW-like model with equal number of parameters. Here we provide analytical expressions for the logarithmic slopes of these models, and compare them with data from real galaxies. Depending on the Einasto parameters of the dark matter halo, one can expect an extrapolated, inner (0.01--1 kpc), logarithmic profile slope ranging from $\\sim -0.2$ to $\\sim -1.5$, with a typical value at 0.1 kpc around $-0.7$.  Application of this (better fitting) model therefore alleviates some of the past disagreement with observations on this issue. We additionally provide useful expressions for the concentration and assorted scale radii: $r_s, r_{-2}, r_{\\rm e}, R_{\\rm e}, r_{\\rm vir}$, and $r_{\\rm max}$ --- the radius where the circular velocity profile has its maximum value.  We also present the circular velocity profiles and the radial behavior of $\\rho(r)/\\sigma(r)^3$ for both the Einasto and Prugniel-Simien models, where $\\sigma(r)$ is the velocity dispersion associated with the density profile $\\rho(r)$.  We find this representation of the phase-space density profile to be well approximated by a power-law with slope slightly shallower than $-2$ near $r=r_{-2}$. ",
        "introduction": "In an interesting turn of events, S\\'ersic's (1963, 1968) 3-parameter function, developed to describe the projected (on the plane of the sky) radial stellar distributions of galaxies, has been shown to also match the internal (3D) density profiles of simulated dark matter halos (Navarro et al.\\ 2004).  Intriguingly, this function was shown to provide a better fit than the 3-parameter NFW-like model with arbitrary inner power-law slope (Diemand, Moore, \\& Stadel 2004; Merritt et al.\\ 2005, 2006).  The functional form of S\\'ersic's equation was independently developed by Einasto (1965, 1968, 1969) and used to describe the internal density profiles of galaxies (see also Einasto \\& Haud 1989 and Tenjes, Haud, \\& Einasto 1994).  We shall therefore refer to this function as the S\\'ersic model when applied to projected distributions, and as the Einasto model when applied to internal density profiles.  One of the main concerns in (Merritt et al.\\ 2005, 2006) was whether the deprojected form of S\\'ersic's function might provide a better description of the density profiles than the Einasto model. Specifically, in Merritt et al.\\ (2006, hereafter Paper I), the analytical approximation to the deprojection of S\\'ersic's function given in Prugniel \\& Simien (1997) was tested along with the Einasto model, the NFW-like model, and various other fitting functions. Overall, the Prugniel-Simien and Einasto models performed the best, providing a good description to both the galaxy- and cluster-sized halos built from hierarchical $N$-body simulations, and also the monolithic cold collapse halos.\\footnote{In passing we note that the 3-parameter anisotropic Dehnen-McLaughlin (2005) model matched the galaxy-sized halos best, but it did not perform so well in describing the cluster-sized halos, and it was unable to describe the halos formed from spherical cold collapses.}  Curiously, while the Prugniel-Simien model provided the best fit to the cluster-sized halos, Einasto's model provided a better fit to the galaxy-sized halos.  In this paper we explore some of the properties of these two models and some of the consequences they imply, and comparisons they enable, with real galaxies and galaxy clusters.  In an effort to help clarify and unify the various parameters of the different models, Section~2 provides relations between such quantities as effective radius $R_{\\rm e}$, virial radius $r_{\\rm vir}$, the radius where the logarithmic slope of the model equals $-2$, $r_{-2}$, and `concentration' as measured by observers and by modelers.  In Section~\\ref{SecCube} we present the phase-space density profiles, or more specifically, the density profiles divided by the cube of their associated velocity dispersion profiles, showing how, for sufficiently large shape parameters $n$, both the Einasto model and the Prugniel-Simien model approximate a power-law $\\sim r^{-2}$ near $r_{-2}$.  In Section~\\ref{SecSlope} we derive the logarithmic slopes of the Einasto and Prugniel-Simien model, and compare these with real data. This is of particular interest because the innermost slope of these models is considerably shallower than $-1$ and in fact equal to zero in the case of the Einasto model at $r=0$.  The inward extrapolation of these models, rather than the NFW-like model, therefore noticeably reduces the disagreement between modelers and observers on this issue.  Our findings are summarized in Section~\\ref{SecSum}. ",
        "conclusions": "\\label{SecSum}   We have provided expressions to relate the half-mass radii of the Einasto and Prugniel-Simien models to a) the radius, $r_{-2}$, where the logarithmic slope of the density profile equals $-2$, b) the virial radius, $r_{\\rm vir}$, and c) the radius where the associated circular velocity profile has its maximum value, $r_{\\rm max}$.  We have shown the dependence of the `concentration' terms $r_{\\rm vir}/r_{-2}$ and $r_{\\rm vir}/r_{\\rm e}$ (and $r_{\\rm vir}/R_{\\rm e}$) on the ratio $\\rho_{\\rm e}/\\langle \\rho_{\\rm univ} \\rangle$, where $\\rho_{\\rm e}$ and $r_{\\rm e}$ (and $R_{\\rm e}$) are the Einasto (and Prugniel-Simien) scale density and half-mass radius.  We also show how the slope of these models at $r_{\\rm vir}$ depends solely on the ratio $\\rho_{\\rm e}/\\langle \\rho_{\\rm univ} \\rangle$ (Figure~\\ref{Conc}).  Over the radial range $10^{-2} < r/r_{-2} < 4 \\times 10^1$, we have shown both the Einasto and Prugniel-Simien models possess the property that $\\rho(r)/{\\sigma(r)}^3$ can be roughly described by a power-law $r^{-\\alpha}$ with the value of $\\alpha$ slightly less than 2 for profile shapes $n$ equal to or greater than the best-fitting values reported both here and elsewhere.  Analytical expressions for the logarithmic slope of the Einasto and Prugniel-Simien models have been derived, and the slope expected from the inward extrapolation of these models, inside of $\\sim$0.01 $r_{\\rm vir}$, is compared with that from observations of real galaxies.  The innermost ($r=0$) slope of the Prugniel-Simien model (0.73--0.87), as currently defined, appears too steep to match all the galaxy data, but agrees with theoretical expectations for a slope of $-0.75$ (Taylor \\& Navarro 2001) and $-0.78$ (Austin et al.\\ 2005). Future work should explore the optimal value of the quantity $p$, the inner logarithmic profile slope in the Prugniel-Simien model.  Setting $p=0$, one recovers the Einasto model, which appears capable of matching the inner profile slopes observed in real galaxies (Figure~\\ref{FigLSB}).  Indeed, the typical value of $\\sim -0.7$ at 0.1 kpc in our CDM halos agrees well with the galaxy data from Swaters et al.\\ (2003a) and Simon et al.\\ (2005), but is steeper than the value $-0.2 \\pm 0.2$ reported by de Blok et al.\\ (2003) and others. We also note that, at present, the pseudo-isothermal model appears inconsistent with the galaxy data (Figure~\\ref{FigIso})."
    },
    "0608/astro-ph0608339_arXiv.txt": {
        "abstract": "We have carried-out multi-configuration Breit--Pauli {\\sc autostructure} calculations for the dielectronic recombination (DR) of Fe$^{8+}$ -- Fe$^{12+}$ ions. We obtain total DR rate coefficients for the initial ground-level which are an order of magnitude larger than those corresponding to radiative recombination (RR), at temperatures where Fe $3\\rmp^q$ ($q=2-6$) ions are abundant in photoionized plasmas. The resultant total (DR+RR) rate coefficients are then an order of magnitude larger than those currently in use by photoionized plasma modeling codes such as {\\sc cloudy, ion} and {\\sc xstar}. These rate coefficients, together with our previous results for $q=0$ and $1$, are critical for determining the ionization balance of the M-shell Fe ions which give rise to the prominent unresolved-transition-array X-ray absorption feature found in the spectrum of many active galactic nuclei. This feature is poorly described by {\\sc cloudy} and {\\sc ion}, necessitating an {\\it ad hoc} modification to the low-temperature DR rate coefficients. Such modifications are no longer necessary and a rigorous approach to such modeling can now take place using these data. ",
        "introduction": "The iron M-shell ions (predominantly Fe$^{7+}$ -- Fe$^{13+}$) give rise to strong X-ray absorption lines due to $n=2-3$ electronic inner-shell transitions. These are seen in the spectra of active galactic nuclei (AGNs) observed by {\\it Chandra} and {\\it XMM-Newton} \\citep{Sak01} as an unresolved-transition-array (UTA). The shape of the UTA feature can be used as a powerful diagnostic --- see \\cite{Beh01} for a detailed discussion, and whose atomic database for describing it has been updated recently by \\cite{Gu06}. However, \\cite{Net03} have pointed-out problems in modeling this shape and so \\cite{Net04}, using {\\sc ion}, and \\cite{Kra04}, using {\\sc cloudy}, have suggested increasing the magnitude of the recombination rate coefficients for Fe ions, particularly the $3\\rmp^q$ ($q=1-6$) ions Fe$^{13+}$ -- Fe$^{8+}$, by postulating {\\it ad hoc} low-temperature dielectronic recombination (DR) rate coefficients. The net effect is to shift the ionization balance towards the neutral end. This brings the modeling results towards accord with observation, but it is not rigorous since there is little to constrain the DR rate coefficients which they use.  The extant recombination data is effectively all radiative (RR) at photoionized plasma temperatures \\footnote{When we refer to photoionized or electron-collisional plasma temperatures we mean the temperatures at which Fe $3\\rmp^q$ ions are abundant in such plasmas --- see \\cite{Kall01} ({\\sc xstar}) and \\cite{Mazz98}, respectively.}. The recommended DR data for these Fe ions \\citep{Arna92a} is based upon the high-temperature electron-collisional plasma results of \\cite{Jac77}, which give essentially zero-contribution from DR at photoionized plasma temperatures. There is much evidence from the Fe L-shell, both experimental and theoretical \\citep[see, e.g.,][]{Sav06}, that a significant contribution can be expected from DR at low temperatures, due principally to `non-dipole' core-excitations which were not (needed to be) considered by \\cite{Jac77}.   Recently, measurements by \\cite{Sch06} \\footnote{\\cite{Sch06} also give an up-to-date list of references of observations of the AGN X-ray absorption feature.} on Fe$^{13+}$ ($q=1$) at the Heidelberg  heavy-ion test storage ring have found the contribution from DR to the total recombination rate coefficient to be an order of magnitude larger than that due to RR, at photoionized  plasma temperatures. We have carried-out a detailed theoretical analysis of these experimental results \\citep{Bad06_Fe13}. We found some disagreements, especially at energies which affect the DR rate coefficient at photoionized  plasma temperatures. Nevertheless, our DR rate coefficient for Fe$^{13+}$ is also an order of magnitude larger than our RR one, but it is up to a third smaller than the experimentally-based one over $10^4 - 10^5$~K. We now report-on the results of calculations for the computationally demanding Fe $3\\rmp^q$  ($q=2-6$) ions.   The remainder of the paper is organized as follows: in Section 2 we describe our methodology, in Section 3 we present our results and make comparisons with the results of earlier works, and then make some concluding remarks. ",
        "conclusions": "We have reported new DR rate coefficients for Fe $3\\rmp^q$ ($q=2-6$) ions which give rise to total recombination rate coefficients which are an order of magnitude larger, at photoionized plasma temperatures, than those currently recommended \\citep{Arna92a} and routinely used by modeling codes such as {\\sc cloudy, ion} and {\\sc xstar}. These new rate coefficients can be expected to change significantly the ionization balance of the Fe M-shell ions which give rise to the important UTA X-ray absorption feature seen in the spectra of AGNs observed by {\\it Chandra} and {\\it XMM-Newton}, and to enable rigorous modeling of it to be carried-out now using these data.  Similarly, large low-temperature DR contributions can be expected for $3\\rmp^q$ ions of other elements of astrophysical importance, e.g., Si, S, Ar, Ni. A move into the $3\\rmd$ sub-shell of lower-charge Fe ions, using a level-resolved approach, is also desirable."
    },
    "0608/hep-th0608224_arXiv.txt": {
        "abstract": "The enormous red-shifting of the modes during the inflationary epoch suggests that physics at the Planck scale may modify the standard, nearly, scale-invariant, primordial, density perturbation spectrum. Under the principle of path-integral duality, the space-time behaves as though it has a minimal length~$L_{_{\\rm P}}$ (which we shall assume to be of the order of the Planck length), a feature that is expected to arise when the quantum gravitational effects on the matter fields have been taken into account. Using the method of path integral duality, in this work, we evaluate the Planck scale corrections to the spectrum of density perturbations in the case of exponential inflation. We find that the amplitude of the corrections is of the order of $({\\cal H}/M_{_{\\rm P}})$, where ${\\cal H}$ and $M_{_{\\rm P}}$ denote the inflationary and the Planck energy scales, respectively. We also find that the corrections turn out to be completely independent of scale. We briefly discuss the implications of our result, and also comment on how it compares with an earlier result. ",
        "introduction": "The inflationary scenario~\\cite{texts,mukhanov92} provides an attractive mechanism to generate the primordial density perturbations that are eventually responsible for the anisotropies in the cosmic microwave background~(CMB) and the formation of the large scale structure~(LSS).  Typically, the CMB and LSS data have been used to constrain the parameters of the inflationary model such as, for example, the height and the slope of the canonical scalar field potential~\\cite{lidsey97,bassett05}.  However, over the last few years, it has been realized that, in most of the models of inflation~\\cite{lyth99}, the period of acceleration lasts so long that length scales that are of cosmological interest today would have emerged from sub-Planckian length scales at the beginning of inflation.  In other words, inflation provides a magnifying glass to probe physics at the very high energy scales (say, $\\sim 10^{17}\\, {\\rm GeV}$) that are otherwise inaccessible to observations or experiments.  This has led to a considerable effort in the literature towards understanding the effects of Planck scale physics on the inflationary perturbation spectrum~\\cite{paddyolp}--\\cite{sriram05}, and the resulting signatures on the CMB~\\cite{minimalcmb,nccmb,lqgcmb,generalcmb}.  In an exponential (i.e. de Sitter) or a power-law inflationary scenario, the spectrum of perturbations is essentially given by the Fourier transform of the Wightman function of a free, massless and minimally coupled scalar field (see, for instance, Ref.~\\cite{mukhanov92}).  Therefore, in order to study the effects of Planck scale physics on the primordial spectrum, we need to understand as to how quantum gravitational effects modify the two-point function of a scalar field in an inflationary background.  But, due to the lack of a viable quantum theory of gravitation, one is forced to consider phenomenological models constructed by hand---models which are supposed to contain one or more features of the actual effective theory obtained by integrating out the gravitational degrees of freedom.  The different high energy models that have been considered in this context modify either the dynamics or the initial conditions (and, in some cases, both) of the canonical scalar field~\\cite{modifieddr}--\\cite{sriram05}.  A class of models that has been used extensively to study the Planck scale corrections to the primordial spectrum involves the violation of local Lorentz invariance at the energy scales of inflation~\\cite{modifieddr,shanki05}.  However, theoretically, there exists no apriori reason to believe that Lorentz invariance may be broken at the inflationary energy scale.  Moreover, certain observations and experiments seem to indicate that local Lorentz invariance may be preserved to extremely high energies (see Ref.~\\cite{vlorentzi}; in this context, also see Ref.~\\cite{vlorentzic}).  In such a situation, to study the Planck scale effects, it becomes important to consider models that {\\it preserve}\\/ local Lorentz invariance {\\it even as they contain a fundamental scale}\\/~\\cite{shanki04}.  One such model that introduces a fundamental length scale while preserving Lorentz invariance is the approach due to the principle of path-integral duality~\\cite{paddypid}. Recall that, in standard quantum field theory, the path integral amplitude for a path connecting events~${\\tilde x}$ and~${\\tilde x'}$ in a given space-time is proportional to the proper length, say, ${\\cal R}({\\tilde x},{\\tilde x'})$ between the two events. The duality principle proposes that the path integral amplitude should be invariant under the transformation ${\\cal R} \\rightarrow \\l(L_{_{\\rm P}}^2/{\\cal R}\\r)$, where $L_{_{\\rm P}}$ is assumed to be of the order of the Planck length.  A consequence of this postulate is that, in the two-point function in the Minkowski vacuum, the proper distance, say, $(\\Delta {\\tilde x})^2$ between two space-time events is replaced by~$\\left[(\\Delta {\\tilde x})^2 + 4\\, L_{_{\\rm P}}^2\\right]$. This suggests that, under the duality principle, the background space-time behaves as though it possesses a minimal length~$L_{_{\\rm P}}$, a feature that is expected to arise when the quantum gravitational effects are taken into account.  In this work, we shall utilize the locally Lorentz invariant approach due to path integral duality to evaluate the Planck scale corrections to the spectrum of density perturbations in the exponential inflationary scenario.  It should be pointed out here that, recently, another approach that is based on T-duality in string theory has been used to evaluate the Planck scale corrections to the standard spectrum of gravitational waves in exponential inflation~\\cite{chouha05}. There seem to exist some correspondence between the approaches of T-duality in string theory and the principle of path integral duality\\footnote{For instance, using the T-duality in string theory, it has been shown that the propagator for the string center of mass results in the same propagator as obtained through the principle of path integral duality. For details, see Ref.~\\cite{Fontanini}.}. However, we would like to emphasize here the following crucial difference between the approach motivated by T-duality in string theory that has been adopted in Ref.~\\cite{chouha05} and the approach due to path integral duality we shall adopt.  In Ref.~\\cite{chouha05}, the authors use a modified dispersion relation to mimic the effects due to T-duality on the propagator.  Clearly, such an approach assumes violation of local Lorentz invariance.  In this work, we shall instead use the proper time representation of the modified propagator~\\cite{paddypid}---an approach that explicitly preserves local Lorentz invariance---to evaluate the Planck scale corrections to the spectrum of perturbations.  As we shall see, the amplitude of the corrections prove to be the order of $({\\cal H}/M_{_{\\rm P}})$, where ${\\cal H}$ and $M_{_{\\rm P}}$ denote the inflationary and the Planck energy scales.  Moreover, in contrast to the results obtained in Ref.~\\cite{chouha05}, we find that the corrections turn out to be completely independent of scale.  The remainder of this paper is organized as follows.  The modifications to the two-point functions that arise due to the path integral duality are easily expressed in terms of the Schwinger's proper time representation for the Greens functions~\\cite{paddypid}. Therefore, in the following section, using the proper time method, we shall very briefly outline as to how the approach of path integral duality modifies the two-point functions of a quantum scalar field propagating in a curved space-time.  We shall also point out the initial conditions that need to be imposed for evaluating the perturbation spectrum through the proper time method.  In Section~\\ref{sec:psptm}, we shall obtain the standard, scale invariant perturbation spectrum in exponential inflation through the proper time formalism.  In Section~\\ref{sec:corrections}, we shall evaluate the Planck scale corrections to the scale invariant spectrum using the approach of path integral duality.  Finally, in Section~\\ref{sec:discussion}, we conclude with a summary and discussion of the results we have obtained.  A few words on our convention and notation are in order at this stage of our discussion.  The metric signature we shall adopt is $(+, -,-, -)$, and we shall set $\\hbar=c=1$.  Also, for convenience, we shall denote the set of four coordinates $x^{\\mu}$ as ${\\tilde x}$. ",
        "conclusions": "\\label{sec:discussion}  In this work, using the locally Lorentz invariant approach of path integral duality, we have evaluated the Planck scale corrections to the primordial perturbation spectrum in the case of exponential inflation.  We find that the Planck scale corrections are of the order of $\\l({\\cal H}/M_{_{\\rm P}}\\r)$.  Moreover, the corrections are completely scale-free, which implies that they simply change the amplitude of the power spectrum.  As we pointed out in the last section, the amplitude in itself cannot be determined from the observations and, as a result, the corrections we have obtained are unobservable.  However, the amplitude of the Planck scale corrections (whether they would depend on, say, $({\\cal H}/M_{_{\\rm P}})$ or $({\\cal H}/M_{_{\\rm P}})^{2}$) have been of theoretical interest and, our calculation assumes importance in this context.  At this stage, it is interesting to compare our results with that of Ref.~\\cite{chouha05}.  As we had mentioned in the introduction, in this work, the authors evaluate the power spectrum of the gravitational waves in exponential inflation.  For computational convenience, they use a particular modified dispersion relation to mimic the high energy effects.  For the case of the minimal energy initial state---which is similar to the vacuum state we have chosen in Eq.~(\\ref{eq:icwf})---they obtain a spectrum which has a blue tilt with superimposed oscillations.  In contrast, using the locally Lorentz invariant proper time method, we obtain a spectrum that is strictly scale invariant with the Planck scale corrections of order $\\l({\\cal H}/M_{_{\\rm P}}\\r)$.  In cosmological perturbation theory, it is well-known that, the metric perturbations do not couple to the matter perturbations in a de Sitter background (see, for instance, Ref.~\\cite{mukhanov92}).  For this reason, the exact de Sitter case is considered pathological and, ideally, one would like to carry out the analysis in slow-roll inflation, or, at least for the case of power-law inflationary scenario.  However, the wave functions $\\chi_{E}$ for $E\\ne 0$ are not known in an analytic form in such cases and one needs to resort to a WKB approximation to evaluate the wave functions and the resulting corrections.  We hope to address this issue in a forthcoming publication."
    },
    "0608/astro-ph0608362_arXiv.txt": {
        "abstract": "We present the results of a coronagraphic imaging search for extra-solar planets around the young main-sequence stars, $\\epsilon$ Eri and Vega. Concentrating the stellar light into the core of the point spread function by the adaptive optic system and blocking the core by the occulting mask in the coronagraph, we have achieved the highest sensitivity for point sources in close vicinity of the both central stars. Nonetheless we had no confidential detection of a point source around the stars. The observations give the upper limits on the masses of the planets to 4 $\\sim$ 6 Jupiter mass and 5 $\\sim$ 10 Jupiter mass at a few arcsecond from $\\epsilon$ Eri and Vega, respectively. Diffuse structures are also not detected around both stars. ",
        "introduction": "Searches for extra-solar planets are very successful; over 180 planets have been discovered. However all discoveries so far were made using indirect methods, i.e. Doppler shift measurements and the transit method. Direct detection --- imaging of the emission from an extra-solar planet --- will open the door to investigating chemistry, meteorology, and biology under conditions completely different from those of the Earth and the other planets in the Solar System.  The small angular separation between a planet and a central star and the huge difference in the brightness of the two objects make direct imaging difficult. The direct detection of extra-solar planets around pre-main sequence stars is less challenging in terms of the brightness difference, since young planets are bright in radiation. Itoh et al. (2005) discovered a young brown dwarf companion to the classical T Tauri star DH Tau, and the companionship was established by proper motion measurements. They derived the effective temperature by comparing its near-infrared spectrum with synthetic spectra of young low-mass objects (Tsuji et al. 2004). The mass is estimated to be 30 -- 50 Jupiter mass ($M_{\\rm J}$) by comparison to evolutionary tracks (Baraffe et al. 2003; D'Antona et al. 1997) on the HR diagram. Chauvin et al. (2004) presented a direct image of a giant-planetary mass object around a young brown dwarf in the TW Hya association. Neuh\\\"auser et al. (2005) announced the discovery of a proto-planet around GQ Lup. They estimated its mass to be 1 -- 40 $M_{\\rm J}$ using evolutionary tracks of Wuchterl et al. (2005) and Baraffe et al. (2003). However, in general, mass estimates of a proto-planet have large uncertainty. First, as pointed out by Itoh et al. (2005), determining the effective temperature by comparison of the spectrum to the spectra of field dwarfs tends to underestimate the value. Moreover, the evolutionary tracks of low-mass objects have large uncertainties. For example, we derive the mass of DH Tau B to be only 5 $M_{\\rm J}$, i.e. in the planetary mass range, if using the evolutionary track of Wuchterl et al. (2005). These two factors make it still unclear whether the low-luminosity companions are proto-planets.  Another approach to direct detection of extra-solar planets is, of course, detection of an extra-solar planet around a main-sequence star. Such an object does not have large ambiguities in the effective temperature estimate and in the mass estimate on the HR diagram. However, such an object is no longer bright in radiation. The flux ratio of the reflection light of a planet to a central star is described as $\\frac{F_{\\rm p}}{F_{\\rm s}} = \\frac{A}{2} (\\frac{R_{\\rm p}}{2a})^{2}$, where $A$, $R_{\\rm p}$, and $a$ are albedo, the radius, and the semi-major axis of the planet, respectively. Reflection lights from all the extra-solar planets discovered by the Doppler shift measurements or the transit methods are difficult to be directly detected in reflected light, due to their faintness ($>$ 21 mag at the $H$-band) and/or due to their small separation ($<$ 0.3\\arcsec) from the star.  Alternative targets for direct detection are unknown extra-solar planets around nearby young main-sequence stars. Such planets are expected to be still bright in a contraction phase. Here we report the results of the coronagraphic observations of extra-solar planets around such young dwarfs, $\\epsilon$ Eri and Vega. These stars are surrounded by dust rings, suggestive not only of their youths but also of the presence of a planet between the star and the ring. ",
        "conclusions": "We have carried out near-infrared coronagraphic observations of $\\epsilon$ Eri and Vega. The observations are one of the deepest near-infrared searches for extra-solar planets in close vicinity of both central stars so far. \\begin{enumerate} \\item We did not detect any trustworthy point source around $\\epsilon$ Eri. The upper limit on the mass of a planet is $4\\sim6 M_{\\rm J}$ in the region between 3\\arcsec (10 AU) and 7\\arcsec (23 AU) from the star. Location of one point source candidate is not inconsistent with the planet suggested by the Doppler shift measurements. Another epoch observation will make it clear whether the candidate is a real object and whether it orbits $\\epsilon$ Eri. \\item We did not detect any point source around Vega. Negative result of our observations puts an upper limit on the mass of a planet to $5\\sim10 M_{\\rm J}$ in the region between 5\\arcsec (40 AU) and 8\\arcsec (60 AU) from the star. \\end{enumerate}"
    },
    "0608/astro-ph0608681_arXiv.txt": {
        "abstract": "The acceleration of the cosmic expansion may be due to a new component of physical energy density or a modification of physics itself.  Mapping the expansion of cosmic scales and the growth of large scale structure in tandem can provide insights to distinguish between the two origins.  Using Minimal Modified Gravity (MMG) -- a single parameter gravitational growth index formalism to parameterize modified gravity theories -- we examine the constraints that cosmological data can place on the nature of the new physics. For next generation measurements combining weak lensing, supernovae distances, and the cosmic microwave background we can extend the reach of physics to allow for fitting gravity simultaneously with the expansion equation of state, diluting the equation of state estimation by less than 25\\% relative to when general relativity is assumed, and determining the growth index to 8\\%. For weak lensing we examine the level of understanding needed of quasi- and nonlinear structure formation in modified gravity theories, and the trade off between stronger precision but greater susceptibility to bias as progressively more nonlinear information is used. ",
        "introduction": "\\label{sec:intro}  Whether the acceleration of the cosmic expansion is due to a new physical component or a modification of gravitation, the answer will involve groundbreaking new physics beyond the current standard models for high energy physics and cosmology.  To obtain the clearest, unbiased picture of the fundamental physics we need to allow for the possibility of gravitation beyond Einstein and see where the data lead. This article presents simultaneous fitting of the cosmic expansion and the theory of gravity.  Direct measurements of the expansion history can be interpreted generally as the equation of state of the universe; this may or may not correspond to a physical component.  Measurements of the growth history of mass fluctuations combine information on the expansion and the theory of gravity.  The two in complementarity thus allow separation of the possible origins for cosmic acceleration -- a physical dark component, e.g.\\ a new field in high energy physics, or dark (new) physics, e.g.\\ a modification of Einstein gravity.  Within a specific theory of modified gravity one can attempt to calculate the cosmological observables and determine the goodness of fit with data.  This model dependent approach must proceed on a case by case basis and moreover suffers from difficulties in computation of many quantities due to the complexity of the theories.  Nor are the modifications necessarily well motivated or completely free from pathologies.  The alternate approach taken here is phenomenological, using a model independent yet physically reasonable and broad parameterization of the gravity modification to gain insight into the effect of generalizing Einstein gravity.  This is closely analogous to the widely successful equation of state approach to modifications of the expansion history, giving model independent constraints and understanding.  In \\S\\ref{sec:index} we discuss the development of the parametrized approach and its range of validity.  \\S\\ref{sec:data} lays out the cosmological probes, fiducial models, and survey data characteristics used in our analysis. We examine in \\S\\ref{sec:constrain} the ability of next generation cosmological probes to reveal new gravity simultaneous with fitting cosmological expansion, and the cosmological bias incurred if we neglect to allow the possibility of beyond Einstein gravity. Leverage and systematics from the nonlinear regime are discussed in \\S\\ref{sec:mods}. ",
        "conclusions": "This article presents an approach for simultaneous fitting of the cosmic expansion and the theory of gravity. We have advocated a ``minimalist'' strategy of distinguishing modified gravity from dark energy, which consists in measuring a single parameter, the growth index $\\gamma$. In addition to reproducing the linear growth function for essentially all standard gravity, dark energy models (parametrized with $w_0$ and $w_a$), the growth index also fits the linear growth of a single known modified gravity theory, the DGP braneworld scenario. Therefore, it is reasonable to expect that the growth index can be used to measure deviations from standard gravity: given the measurements of the background expansion rate (parameterized, say, by $\\Omega_M$, $w_0$ and $w_a$), standard gravity predicts the value of $\\gamma$, and a statistically significant deviation from this value can in principle be interpreted as evidence for -- and characterization of -- beyond Einstein gravity.  This is a first step, hence our emphasis in calling it Minimal Modified Gravity.  One could examine more complex schemes but these may be more model dependent or employ more parameters and are therefore likely to give weaker results.  The MMG approach therefore has more statistical power, being more suitable to near-future data, while we think the loss in generality is minimal.  We also emphasize the advantage in retaining the maximal physics information by treating the expansion effects on growth through the effective equation of state, giving clear separation from deviations in the gravity theory.  How accurately such a program can reveal beyond Einstein gravity -- dark physics vs.\\ physical darkness -- is the main topic of this paper. We have shown that measurements from weak gravitational lensing, Type Ia supernovae, and the CMB combined can measure the growth index to about 8\\%, or to $\\pm 0.04$ around its $\\Lambda$CDM value (and galaxy cluster data could potentially reduce this even further). At the same time, the constraints on other cosmological parameters are not appreciably degraded, essentially because the surveys probe a range of scales and thus their complementarity breaks the degeneracies between parameters.  One particular concern in the program of distinguishing general relativity from modified gravity is the nonlinear density regime of structure formation. Even for the limited number of well-defined modified gravity theories, details of nonlinear clustering are currently unknown. While in principle the nonlinear structure formation is calculable from N-body simulations of modified gravity, creating these simulations in practice is extremely difficult except for some very simple cases. The structure and evolution of galaxy clusters, which are nonlinear objects, is fairly strongly dependent on the nonlinear physics, and is consequently problematic.  Weak lensing, on the other hand, probes a range of scales, and we studied how our results behave if we drop small-scale (that is, nonlinear) information. Using the nulling tomography approach and a reasonably well motivated toy model for bias due to uncertainty in nonlinear structure, we found that cutting out the small scale information ($k\\gtrsim 1\\hmpcinv$) can lead to significant decrease in the resulting bias in the growth index, at the expense of increasing the statistical error in it by a factor of a few.  Other cosmological observables exist with sensitivity to the growth of fluctuations, and hence can be used to constrain MMG, but we have not discussed them in any detail.  For example, the bispectrum of weak gravitational lensing is a potentially powerful probe \\cite{takada_jain}; however, it has proved to be a tough task to calibrate the bispectrum even in standard general relativity.  The same holds for Lyman-alpha forest observations. The Integrated Sachs-Wolfe effect is a potentially strong discriminant of modified gravity models (see in particular recent predictions of the ISW in DGP models \\cite{song_sawicki_hu}); however, the metric potentials are particularly sensitive to the structure of the modified gravity theory.  Neither we nor anyone else has yet succeeded in finding a generic parametrization of the deviations from general relativity for the ISW effect.  Our work outlines a first step in treating modified gravity models. At the time of this writing, there is hardly a single well-defined modified gravity theory that does not look like standard general relativity in terms of observables, but is not already ruled out or disfavored by data. Considerable effort is underway to construct such theories (e.g.\\ \\cite{luess2,chiba,carroll_trodden,weller,odintsov,navarro,gabadadze,song,zhang, amendola,defelice,msg,gauss1,gauss2,ghost}; for a review see \\cite{lue_overview,uzan,trodden_review}) and test them experimentally.  It is heartening that the strong complementarity in cosmological probes such as the combination of weak lensing, supernovae, and the CMB does provide important information on the question of dark physics vs.\\ physical darkness.  Such data from next generation experiments, here calculated specifically for SNAP and Planck, does furnish a real test not just of individual models but of the physics framework beyond Einstein.  We will be able to measure the effective equation of state describing the cosmic expansion and simultaneously reveal the theory of gravity.  By continuing forward with advances in measurements, theory, and computation we can lift the darkness on the new physics."
    },
    "0608/astro-ph0608648_arXiv.txt": {
        "abstract": "The Sc galaxy M\\,99 in the Virgo cluster has been strongly affected by tidal interactions and recent close encounters, responsible for an asymmetric spiral pattern and a high star formation rate. Our {\\it XMM-Newton} study shows that the inner disk is dominated by hot plasma at $kT \\approx 0.30$ keV, with a total X-ray luminosity $\\approx 10^{41}$ erg s$^{-1}$ in the $0.3$--$12$ keV band. At the outskirts of the galaxy, away from the main star-forming regions, there is an ultraluminous X-ray source (ULX) with an X-ray luminosity $\\approx 2 \\times 10^{40}$ erg s$^{-1}$ and a hard spectrum well fitted by a power law of photon index $\\Gamma \\approx 1.7$. This source is close to the location where a massive H\\,{\\footnotesize I} cloud appears to be falling onto the M\\,99 disk at a relative speed $> 100$ km s$^{-1}$. We suggest that there may be a direct physical link between fast cloud collisions and the formation of bright ULXs, which may be powered by accreting black holes with masses $\\sim 100 M_{\\odot}$. External collisions may trigger large-scale dynamical collapses of protoclusters, leading to the formation of very massive ($\\ga 200 M_{\\odot}$) stellar progenitors; we argue that such stars may later collapse into massive black holes if their metal abundance is sufficiently low. ",
        "introduction": "The Sc galaxy M\\,99 (NGC\\,4254), located a distance of about $17$ Mpc (Tully 1988), is the brightest spiral in the Virgo Cluster ($M_B = -20.8$), with a gas mass $\\approx 5 \\times 10^{9} M_{\\odot}$ and a kinematic mass $\\approx 10^{11} M_{\\odot}$ (Vollmer, Huchtmeier \\& van Driel 2005; Phookun, Vogel \\& Mundi 1993). It shows a peculiar spiral structure\\footnote{As a matter of historical curiosity, it was the second galaxy ever in which a spiral pattern was identified (Rosse 1850), with the ``Leviathan of Parsonstown'', a few months after a similar discovery for M\\,51.}, with one arm less tightly wound and much brighter than the other two. Another unusual feature of M\\,99 is the presence of a string of H\\,{\\footnotesize I} ``blobs'' to the south and west of the stellar disk, and a low-surface-density H\\,{\\footnotesize I} gas tail to the north-west; the total gas mass in those extra-disc structures is  $\\approx 2 \\times 10^{8} M_{\\odot}$ (Phookun et al.~1993). It is likely that there is a physical connection between the unusual spiral structure and the disturbed gas distribution. In one scenario (Vollmer et al.~2005), the spiral structure was affected by a close encounter with another Virgo Cluster galaxy $\\approx 280$ Myr ago; the surrounding H\\,{\\footnotesize I} clouds are due to ongoing, face-on ram-pressure stripping, and are mostly moving away from the galaxy. An alternative scenario (Phookun et al.~1993) suggests that the H\\,{\\footnotesize I} clouds and tail are tidal debris of an entity that was at least partly disrupted in a close encounter with M\\,99; that debris is now infalling and merging with the galactic disk. Simulations show (Bekki, Koribalski \\& Kilborn 2005) that tidal debris could also have been stripped from the outer disk of M\\,99 itself, by the Virgo Cluster potential as the galaxy crossed the central region of the cluster; part of the stripped gas would later fall back onto the same galaxy. In this scenario, heavy gas infall onto the M\\,99 disk would be responsible for the lopsided spiral structure (Phookun et al.~1993; Bournaud et al.~2005).  Intriguingly, a large H\\,{\\footnotesize I} cloud (VIRGOHI\\,21), with a gas mass $\\approx 2 \\times 10^8 M_{\\odot}$ but without any associated stars, located $\\approx 120$ kpc to the north-west of M\\,99, has recently been discovered (Davies et al.~2004). One interpretation (Minchin et al.~2005, 2006) is that it is an old, bound ``dark galaxy'' that has never formed stars because of its low hydrogen surface density. In this scenario, VIRGOHI\\,21 would be dark-matter dominated, with a kinematic mass $\\approx 10^{11} M_{\\odot}$; if so, it would likely be the same galaxy responsible for the suspected close encounter with M\\,99 and possible tidal gas stripping. Alternatively, VIRGOHI\\,21 could itself be simply another, larger piece of tidal debris (without a dark matter halo), stripped from the outer H\\,{\\footnotesize I} disk of M\\,99 by the Virgo Cluster potential or during a close encounter with another galaxy (Bekki et al.~2005)  Gas-rich galaxies with recent tidal interactions generally show particularly active star formation, and M\\,99 is no exception, with a rate $\\approx 10 M_{\\odot}$ yr$^{-1}$ inferred from its H$\\alpha$ luminosity (Kennicutt et al.~2003), or $\\approx 10$--$20 M_{\\odot}$ yr$^{-1}$ from an optical/near-IR photometric study (Gonzales \\& Graham 1996). As a comparison, this is twice as high as in M\\,82 (Kennicutt et al.~2003), and three times the current total rate in our Local Group (Hopkins, Irwin \\& Connolly 2001). Another phenomenon often associated with tidal interactions and high star formation is the presence of ultraluminous X-ray sources (ULXs) with apparent isotropic X-ray luminosities $\\sim 10^{39}$--$10^{40}$ erg s$^{-1}$ (for recent reviews, see King 2006; Fabbiano \\& White 2006; Colbert \\& Miller 2004). They are interpreted as accreting binary systems, an order of magnitude more luminous than the Eddington limit for typical stellar-mass black holes (BHs) in the Local Group. Two ULXs were detected by {\\it ROSAT}/HRI within the $D_{25}$ region of M\\,99, on 1997 June 30 (Colbert \\& Ptak 2002; Liu \\& Bregman 2005). In this paper, we use {\\it XMM-Newton} to have a better understanding of the X-ray properties of M\\,99, and in particular of its brightest ULX. We compare radio and X-ray information to investigate the possible connection between ULX formation, star formation and collisional events in this galaxy. Finally, we outline a possible general scenario for the nature of ULXs, to be tested by future studies. ",
        "conclusions": "We studied the X-ray properties of M\\,99, the most massive spiral galaxy in the Virgo Cluster. As expected from its high SFR ($\\approx 10 M_{\\odot}$ yr$^{-1}$), the X-ray emission is dominated by a soft thermal-plasma component. The total unabsorbed luminosity (not including a bright ULX) inside the $D_{25}$ ellipse is $\\approx (1.2 \\pm 0.2) \\times 10^{41}$ erg s$^{-1}$ in the $0.3$--$12$ keV band; about $15\\%$ of this is due to resolved or unresolved discrete sources. The emission appears almost uniformly diffused across the inner disk (at $\\la 5$ kpc from the nucleus), and there is no starburst core. The temperature of the hot gas, $kT \\la 0.30$ keV, is also rather low, typical of disk emission rather than of a starburst environment.  We briefly discussed the morphology of the nucleus in the optical and UV bands, showing the presence of a (redder) massive nuclear star cluster and a few smaller, much bluer clusters of young stars around it. The spatial resolution of {\\it XMM-Newton} does not reveal whether there are faint point-like X-ray sources in the nuclear region, and if so, whether they are located in the old nuclear cluster or associated to the young stars around it.  The main goal of this paper was a study of the properties and origin of a bright ULX at the outer edge of the stellar disk. Its unabsorbed luminosity of $\\approx 2 \\times 10^{40}$ erg s$^{-1}$ in the $0.3$--$12$ keV band, together with a pure power-law spectrum with photon index $\\Gamma \\approx 1.7$, are difficult to interpret in the framework of classical spectral states for stellar-mass BHs. Such high/hard states are not uncommon in ULXs but are not generally seen in Galactic BHs. It suggests that we are not seeing the accretion disk directly, and that most of the gravitational power is efficiently transferred to a comptonizing region.  An intriguing new discovery is that there is a massive gas cloud (H\\,{\\footnotesize I} mass $\\sim 10^7 M_{\\odot}$) seen in projection very close to the ULX position. From its radial velocity, we suggested that this cloud is falling onto the galactic disk (from behind) and perhaps impacting it near the ULX location. The cloud occupies $\\approx 3\\%$ of the projected area of the stellar disk; therefore, we cannot entirely rule out the possibility that it is a chance association. However, we explored the speculative idea that there is a direct connection between collisional events and ULX formation, and not simply an indirect effect due to an enhancement in star formation. We suggested a possible scenario of ULX formation that would be consistent with this interpretation. We argued that cloud collisions may trigger a large-scale dynamical collapse of molecular clumps, rather than Jeans-mass fragmentation; in turn, this may lead to the formation of very massive stellar progenitors in the cluster core. If the metal abundance is low ($Z \\sim 0.1 Z_{\\odot}$,  not primordial abundances), we speculate that the star  might retain enough of its gas to directly-collapse into a BH massive enough (up to $\\sim 100 M_{\\odot}$) to explain the luminosity of the brightest ULXs.  This scenario has to be tested with further individual and population studies of ULXs. At the individual level, the suggestion that ULXs have masses up to an order of magnitude higher than Galactic BHs (but not more) needs to be tested with more advanced X-ray spectral modelling, and a better understanding of the thermal disk and comptonized emission components in such systems. At the population level, we need a more systematic study of the spatial association between ULXs, low-metallicity environments and collisional events. Further work on this issue is currenly under way (D. Swartz et al., in preparation)."
    },
    "0608/gr-qc0608113_arXiv.txt": {
        "abstract": "A paradigm deeply rooted in modern numerical relativity calculations prescribes the removal of those regions of the computational domain where a physical singularity may develop. We here challenge this paradigm by performing three-dimensional simulations of the collapse of uniformly rotating stars to black holes without excision. We show that this choice, combined with suitable gauge conditions and the use of minute numerical dissipation, improves dramatically the long-term stability of the evolutions.  In turn, this allows for the calculation of the waveforms well beyond what previously possible, providing information on the black-hole ringing and setting a new mark on the present knowledge of the gravitational-wave emission from the stellar collapse to a rotating black hole. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608224_arXiv.txt": {
        "abstract": "The temporal structure of Gamma Ray Bursts can be interpreted assuming as a inner engine a neutron star which undergoes a progressive compactification via production of strangeness (hyperons and kaons) and quarks. We will propose a tentative identification of various emission periods of the burst with specific structural changes of the star. Each of these modifications of the composition of the compact star takes place as a deflagration and not as a detonation, so the energy released in the transition goes mainly into heat and not into a mechanical wave. This is important in order to avoid an excessive baryonic contamination of the region surrounding the compact star. In this way a ultrarelativistic plasma of electron-positron pairs and of photons can be obtained, powering the Gamma Ray Burst. ",
        "introduction": "In last years the data collected by various X-ray satellites and notably by Beppo-SAX and by Swift have clarified many crucial features of the Gamma-Ray Bursts (GRBs) so that a rather precise phenomenological description of the explosion can now be drawn. The light curve can be separated roughly in four emission periods, although some of these features can often be absent in a specific burst (for a recent review see e.g. \\cite{meszaros}).  1) Several bursts present a {\\it precursor}, namely a small signal containing only a tiny fraction of the total energy of the burst, which anticipates the main event by tens or even hundreds of seconds. The exact fraction of GRBs showing a clear evidence of a precursor is difficult to estimate due to the smallness of the signal. It has been speculated by Lazzati \\cite{lazzati} that at least 20$\\%$ of the GRBs display a precursor and this fraction could still be an underestimate due to the difficulty of detecting precursors anticipating the main event by only a few seconds. The duration of the precursor is typically of a few seconds.  2) The main event corresponds to the emission with the highest luminosity and is present in the highest energy band of the emission spectrum. The duration of the main event can vary from few seconds (here we are discussing only long bursts, having durations longer than roughly 2 s) up to hundreds of seconds. As we will show, it is possible to divide the main event into active periods whose duration can be related to the activity of the so-called inner engine, which is the source of the energy of the burst. The active periods are separated by quiescent times. The high energy photons emitted during the main event are probably associated with internal shocks of the ultrarelativistic jet produced by the inner engine. In this scenario, shells with different velocities collide among themselves producing the detected signal.  3) Swift satellite has recently provided a strong indication that a large fraction of GRBs, after the main event and an initial drop in luminosity, displays a plateau in which the luminosity drops much less rapidly. Inside the plateau some flares can also be present. The luminosity of the plateau is much smaller than that of the main event, but its duration can be much longer, order of thousands of seconds, so that the total energy released can be comparable to the energy released during the main event. The plateau can be originated by a residual activity of the inner engine, although other interpretations are possible.  4) At last the luminosity drops steadily and the so-called afterglow begins. This final prolongated emission is explained in terms of a progressive expansion and slow-down of the ultrarelativistic jet, which interacts with the interstellar medium emitting photons. This last part of the GRB emission is now the less controversial.  In this contribution we concentrate on the main event and we provide some suggestions concerning the precursor. ",
        "conclusions": ""
    },
    "0608/astro-ph0608012_arXiv.txt": {
        "abstract": "We present a complete analysis, which includes morphology, kinematics, stellar populations and N-body simulations of {\\small CGCG 480-022}, the most distant (c$z=14317~{\\rm km~s}^{-1}$) isolated galaxy studied so far in such detail. The results all support the hypothesis that this galaxy has suffered a major merger event with a companion of $\\sim0.1$ times its mass. Morphology reveals the presence of a circumnuclear ring and possibly further ring debris. The radial velocity curve looks symmetrical, whilst the velocity dispersion increases with radius reaching values that do not correspond to a virialized system. Moreover, this galaxy deviates significantly  from the Fundamental Plane and the Faber-Jackson relation. The stellar population analysis show that the ring is younger and more metal-rich, which suggest that it has undergone a fairly recent burst of star formation. Both morphological and dynamical results are in broad agreeement with our N-body simulations. ",
        "introduction": "Elliptical galaxies located in very low density environments are highly uncommon but hold important clues about galaxy formation and evolution. For example, current hierarchical models for galaxy formation predict that ellipticals in low density environments have stellar populations which are younger by 1-2~Gyr than their cluster counterparts \\citep[e.g.][]{baugh98,clemensetal06,delucia06}. In this framework, isolated ellipticals are thought to be formed in mergers of galaxy pairs or small groups of galaxies \\citep{jones00,donghia05}. In the first case, the mergers will give rise to kinematical misalignments whilst, in the second, they will produce an increasing amplitude of rotation with galactocentric radius, and also rotation around the minor axis \\citep{weilhernquist96}.  Detailed studies involving morphology, kinematics and/or stellar populations of isolated ellipticals have been restricted to nearby galaxies \\citep[c$z$$<$10000 km s$^{-1}$; e.g.][]{collobertetal06}. In this paper, we perform an exhaustive analysis of {\\small CGCG 480-022} (c$z$=14317 km s$^{-1}$), which includes morphological, kinematical, stellar population and N-body simulation properties. We present evidence supporting the hypothesis that this galaxy suffered a major merger event.  We adopt a cosmological model with H$_0$=71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m$=0.3 and $\\Omega_\\Lambda$=0.7. ",
        "conclusions": "The morphological, kinematical, N-body simulations and stellar population analysis all suggest that {\\small CGCG 480-022} could have experienced a major merger event with an alleged companion. In this section, we will focus on the evidences provided by each approach.   {\\em Morphology}. The residuals obtained when substracting a 2D model from the images show, in all bands, the presence of a ring of $r$$\\sim$5~$arcsec$. This ring, together with outer streams of material appearing in K-band images, could trace the propagating sound-waves produced by a merger \\citep[e.g.][]{hernquist87} more likely than by the action of tidal interactions or asymmetric star formation\\footnote{Even though we found a star formation burst in the ring, which would be induced by the passing wave; i.e. the varying SF history should be the effect of the ring structure, not the cause.} \\citep{colbert01}. In fact, \\citet{malincarter83} found that morphological peculiarities, such as shells or ripples, occur roughly 5 times more frequently in environments outside of rich clusters. This result was confirmed by \\citet{colbert01} who concluded that shells are much more prevalent in isolated galaxies than in group galaxies. Note that the ROSAT X-ray image of the galaxy \\citep{zimmermann01} shows two symmetric lobes in the N-S direction.  {\\em Kinematics}. The most intriguing of our results concerns the kinematics of the galaxy. First, within the errors, we find a constant $\\sigma$$\\sim$400~km~s$^{-1}$ which decreases in the interior of the ring. A similar pattern has been found in other isolated massive galaxy \\citep[{\\small NGC1700};][]{bergmann02}. However, such high $\\sigma$ does not correspond to a virialized system but, in a merger scenario, the intrinsic velocity dispersion should be affected by the collision, yielding higher values of $\\sigma$. Second, the radial velocity curve looks symmetrical. Furthermore, both $\\sigma$ and $V_{\\rm r}$ values bump at the position where the ring is located, i.e. 3$<$$|r|$$<$6~$arcsec$. As an additional test, we checked whether {\\small CGCG 480-022} follows or not the Fundamental Plane \\citep[FP;][]{djorgovskidavis87} and the Faber-Jackson \\citep{faberjackson76}, the Kormendy \\citep{kormendy77} and the Mg$_2$-$\\sigma$ \\citep{terlevich81} relations. We considered these relations in the near-IR, with data taken from \\citet{pahre98}. {\\small CGCG 480-022} deviates $3\\sigma$ from the FP, $6\\sigma$ from the Faber-Jackson relation and $4\\sigma$ from the Mg$_2$-$\\sigma$ relation; but just $0.2\\sigma$ from the Kormendy relation. Our galaxy is an outlier in those relations involving the velocity dispersion, in the sense that it {\\em should} be lower {\\em if} it corresponds to the intrinsic velocity dispersion of a virialized system. Thus, a major event affecting the dynamics of this galaxy must have occurred. A similar explanation could be proposed for the $V_{\\rm r}$ curve, which does {\\em not} correspond to a rotating system. In order to disentangle the kinematical peculiarities found in this galaxy, future 2D-spectroscopy  will be extremely helpful, specially to check if our derived kinematics depend on the slit direction.  {\\em N-body simulations}. The N-body experiments performed suggest that a possible formation mechanism is a merger between two elliptical galaxies with a mass ratio near 10:1, and meeting each other on a grazing orbit. In such an encounter the smaller system spirals inwards while the merger is proceeding finally into a torus-like structure. Such torus as seen face-on resembles a ring and the kinematics imposed to the overall field would present a discontinuity. A significant success for this model is that it reproduces without fine tuning the velocity profile observed  in {\\small CGCG 480-022}. The central zones moves with a larger velocity than the zones at increasing galactocentric distances, with respect to the systemic velocity of  the galaxy, and there are secondary maxima at the radial location of the ring in both directions from the center. The range of values of the velocity dispersion predicted by the model is commensurate with that observed, but the decline in value as the center is approached is not reproduced. A reasonable explanation for this is that more ordered motion near the galaxy center could result from an inner disk within the ring radius. The present merger simulation has not, for simplicity, included the presence of  gas, but merger models which do include gas \\citep[e.g.][]{barnes02} show a tendency for the gas to fall towards the center of the galaxy and form a disk. Also, more detailed N-body+SPH simulations including star formation and feedback processes will help to better understand the central kinematics of the system.  {\\em Stellar populations}. Following our merger hypothesis, there should be a star formation burst in the ring due to the passing sound-wave. Therefore, the stars in the ring should be younger and more metal-rich than those within inner regions, as we found. Given that we know the geometry of the ring, we can broadly estimate the crossing time of the sound-wave through the position of the ring, under the following assumptions: {\\em i)} the sound speed is constant throughout the whole system, and {\\em ii)} the system behaves as an ideal gas. Then, the mean pressure of a system of particles with a given velocity dispersion, $\\sigma$, is $p$=$\\frac{1}{3}\\rho\\sigma^2$. The propagation velocity of the perturbation corresponds to the sound speed, $c_s$, as $c_s^2$=$\\frac{{\\rm d}p}{{\\rm d}\\rho}$=$\\frac{1}{3}\\sigma^2$$\\Rightarrow$$c_s$$\\sim$180~km~s$^{-1}$. Considering that the ring radius is $\\sim$5~$arcsec$, then the ellapsed time since the formation of the perturbation is $t(s)$=4.8$\\times$10$^{-6}D/c_s$, where $D$ is the distance to the galaxy in km. Therefore, $t$ should be the stellar age difference between the ring and the centre. Given that $D$=6.2$\\times$10$^{23}$~km, then $t$=5.3 Gyr, which is in good agreement with our stellar population results (see Figure \\ref{poblaciones}). Even though approximative, these numbers do support the merger hypothesis.\\\\  A question we would like to pose is why the most distant isolated elliptical galaxy so far studied in such detail looks so peculiar. To confirm the results presented here and probe the origin of this system, deep 2D spectroscopy is strongly needed. We also note that {\\small CGCG 480-022} is not a completely isolated galaxy but has one confirmed companion at a distance of 0.25~Mpc: {\\small PGC 003298} \\citep{beers83}. The DSS image of this galaxy shows an elongated shape in the direction of {\\small CGCG 480-022}, whilst J- and K-band 2MASS images present some tail-like structures. Will {\\small PGC 003298} be the next to merge?"
    },
    "0608/astro-ph0608297_arXiv.txt": {
        "abstract": "We present observations of two LMC supernova remnants (SNRs), DEM L238 and DEM L249, with the {\\it Chandra} and {\\it XMM-Newton} X-ray satellites. Bright central emission, surrounded by a faint shell, is present in both remnants. The central emission has an entirely thermal spectrum dominated by strong Fe L-shell lines, with the deduced Fe abundance in excess of solar and not consistent with the LMC abundance. This Fe overabundance leads to the conclusion that DEM L238 and DEM L249 are remnants of thermonuclear (Type Ia) explosions. The shell emission originates in gas swept up and heated by the blast wave. A standard Sedov analysis implies about 50 $M_\\odot$ in both swept-up shells, SNR ages between 10,000 and 15,000 yr, low ($\\lesssim 0.05$ cm$^{-3}$) preshock densities, and subluminous explosions with energies of $3 \\times 10^{50}$ ergs. The central Fe-rich supernova ejecta are close to collisional ionization equilibrium. Their presence is unexpected, because standard Type Ia SNR models predict faint ejecta emission with short ionization ages. Both SNRs belong to a previously unrecognized class of Type Ia SNRs characterized by bright interior emission.  Denser than expected ejecta and/or a dense circumstellar medium around the progenitors are required to explain the presence of Fe-rich ejecta in these SNRs. Substantial amounts of circumstellar gas are more likely to be present in explosions of more massive Type Ia progenitors. DEM L238, DEM L249, and similar SNRs could be remnants of ``prompt'' Type Ia explosions with young ($\\sim 100$ Myr old) progenitors. ",
        "introduction": "Heavy elements are produced in stars and then injected into the interstellar medium (ISM) by the combined action of stellar winds and stellar explosions.  Thermonuclear and core-collapse explosions are dominant sources of heavy elements, and they govern the evolution of chemical abundances in galaxies. Thermonuclear (Type Ia) explosions inject iron-rich supernova (SN) ejecta into the ISM, while ejecta of core-collapse (CC) SNe contain large quantities of oxygen and other $\\alpha$-elements. Heavy elements are dispersed by remnants of SNe and eventually mix with the ambient ISM, forming subsequent generations of stars. The O/Fe (or more generally $\\alpha$-elements/Fe) ratio depends on the relative frequency of CC and Type Ia explosions, and on how a galaxy has evolved in the past.  It is now well established that this ratio varies among galaxies.  In particular, the $\\alpha$/Fe ratio is lower in the Magellanic Clouds (MCs) than in our Galaxy \\citep[e.g.,][]{pompeia06}. This suggests that chemical enrichment by Type Ia SNe has been relatively more important in the MCs than in our Galaxy. It is not clear whether this is caused by an increased frequency of Type Ia versus CC explosions in the MCs \\citep{tsijimoto95} or by differences in evolutionary histories between the MCs and our Galaxy \\citep{pata98}. Our poor understanding of progenitors of Type Ia explosions is a significant problem as well; the recent proposal of a new class of ``prompt'' Type Ia SNe \\citep*{mdp06,sullivan06} associated with relatively young ($\\sim 100$ Myr) stars may force reevaluation of current ideas about the poorly understood chemical evolution of the MCs.  Studies of remnants of Type Ia explosions in the MCs allow us to observe injection of Fe into the ISM directly and to study SN ejecta in detail, leading to a better understanding of Type Ia explosions and the enrichment of the MCs in heavy elements.  \\citet{tuohy82} discovered four supernova remnants (SNRs) in the Large Magellanic Cloud (LMC) with strong Balmer lines of hydrogen and weak forbidden lines of heavy elements; such Balmer-dominated remnants are believed to be of Type Ia origin.  X-ray observations with the {\\it Advanced Satellite for Cosmology and Astrophysics (ASCA)} confirmed a Type Ia origin for Balmer-dominated SNRs 0509-67.5 and 0519-69.0 by detecting strong Fe lines in their X-ray spectra \\citep{hughes95}. Their ages are 400 and 600 yr, respectively, based on the recent discovery of light echoes from these explosions \\citep{rest05}. The other two Balmer-dominated SNRs, DEM L71 and 0548-70.4, were also found to be rich in Fe \\citep{hughes98,hendrick03}.  There are several other Type Ia SNRs in the MCs, as indicated by their Fe overabundance.  The bright SNR N103B is only $\\sim 1000$ yr old, based on one detected light echo \\citep{rest05}; its origin is still under debate, with \\citet{vanderheyden02} and \\citet{lewis03} advocating a CC and Type Ia origin, respectively.  Most progenitors of Type Ia SNe are believed to belong to old stellar populations.  Because at most a modest amount of circumstellar material may be associated with the pre-explosion evolution of their low-mass progenitors, SN ejecta dissipate their kinetic energy mostly interacting with the ambient ISM.  Lack of a dense circumstellar medium for most Type Ia SNe is consistent with the absence of radio emission from SNe Ia \\citep{weiler02}.  Type Ia SN progenitors are usually located far from clumpy star formation regions and generally explode in a relatively uniform diffuse ISM.  Remnants of historical Type Ia SNe of Tycho and 1006 indeed exhibit rather smooth outer blast waves. \\citet{badenes06} successfully modeled the Tycho X-ray spectrum with SN ejecta colliding with an ambient medium of constant density, finding good agreement for delayed-detonation models of Type Ia explosions. This encouraging agreement gives credence to the conventional wisdom in which most Type Ia SNRs arise from the interaction of SN ejecta with a rather uniform ambient medium, without needing to invoke the presence of dense circumstellar material in the immediate vicinity of their progenitors. Hydrodynamic models for such an interaction have been considered in the past \\citep[e.g.,][]{dwarkadas98}, and an extensive grid of models (including synthetic X-ray spectra) has become recently available \\citep*{badenes03,badenes05,badenes06}.  We present {\\it Chandra} and {\\it XMM-Newton} observations of two moderately ($\\sim 10^4$ yr) old SNRs in the LMC, DEM L238 (\\objectname{B0534-70.5}) and DEM L249 (\\objectname{B0536-70.6}), which unexpectedly reveal the presence of Fe in their interiors, presumably produced by thermonuclear explosions of their progenitors. These $3\\arcmin$--diameter SNRs have been previously studied at optical and X-ray wavelengths \\citep{mathewson83,williams99}, but these studies have not revealed anything unusual, except for the X-ray emission enhancement in their interiors. We find that the presence of Fe-rich ejecta at their centers is not consistent with the standard Type Ia SNR models just mentioned. The presence of denser than expected material in the vicinity of their SN progenitors or in SN ejecta is required. We discuss several possible scenarios that could explain this finding. ",
        "conclusions": "Our main finding is the discovery of Fe-rich gas in the centers of the moderately old ($\\sim 10^4$ yr) SNRs DEM L238 and DEM L249. An important property of this gas is its very long ionization age; we found the interior X-ray spectra of both SNRs well fit by collisional equilibrium ionization models. Such remnants, with central Fe overabundances and long ionization timescales, have been found before. In addition to four Balmer-dominated SNRs and N103B mentioned in \\S~1, several other Fe-rich remnants have been discovered in the MCs: 0534-69.9 \\citep{hendrick03}; DEM S128, IKT5, and IKT25 \\citep{vander}; DEM L316A \\citep{nishiuchi01,williams05}; and 0454-67.2 \\citep{seward06}. Among them, DEM L316A is most similar to DEM L238 and DEM L249, as the Fe-rich gas at its center is in ionization equilibrium, and a swept-up shell is present. The SMC SNRs DEM S128 and IKT 5 also have Fe-rich gas in CIE, but their shells have not been detected by {\\it XMM-Newton}.  There is then a whole class of remnants with centrally located Fe-rich gas in CIE, and faint or undetectable X-ray shells, making up a substantial fraction of all MC Type Ia SNRs.  It is usually assumed that the presence of Fe-rich ejecta is a signature of a Type Ia explosion.  However, DEM L238 and DEM L249 exhibit a number of puzzling properties that are hard to understand in the framework of the standard Type Ia models mentioned in \\S~1. The long ionization times of the Fe-rich gas imply that the gas density has had to be high enough in the past. In the standard Type Ia models of \\citet{badenes03,badenes05,badenes06}, this can be accomplished only by increasing the ambient medium density (assumed to be constant). But the predicted shell emission in such models is orders of magnitude brighter than observed. If we use instead the blast wave properties derived from the Sedov fits to the shell emission, which imply low ISM densities, then the Fe-rich ejecta are always at low densities and ionization timescales are short. We demonstrate this by considering simple models with an exponential density profile and uniform ambient ISM \\citep{dwarkadas98}, with ejecta mass of 1.4 M$_\\odot$ composed of pure Fe, and with $E_k = 3.3 \\times 10^{50}$ ergs, $n_0 = 0.05$ cm$^{-3}$, and $t_{SNR} = 13,500$ yr derived from the Sedov analysis of DEM L238. We performed our one-dimensional simulations with the Virginia Hydrodynamics PPM (VH-1) code. This code was written and tested by the Virginia Numerical Astrophysics Group and is based on the Lagrangian-Remap version of the piecewise parabolic method developed by \\citet{colwoo84}. We employed a simplified method devised for multidimensional hydrodynamic simulations coupled with NEI calculations to follow Fe ionization states within the ejecta. We used the NEI, version 1.1 spectral code in XSPEC to calculate X-ray spectra \\citep[Fe L-shell atomic data in all NEI models in XSPEC are based on theoretical calculations of][]{liedahl}.  (The low-ionization timescales we found led us to use NEI, version 1.1 instead of the APEC-based version 2.0, because the former includes more of the lower ionization transitions.) The calculated X-ray spectra are compared with the central {\\it Chandra} X-ray spectrum of DEM L238 in Figure \\ref{fig7}. The failure of a standard Type Ia model is obvious, as the shocked Fe in the model is underionized and the ejecta X-ray emission is more than an order of magnitude too faint. An inspection of more realistic Type Ia models of \\citet{badenes03,badenes05,badenes06} confirms this conclusion. The fundamental problem is that the density of the Fe-rich ejecta is too low.  The presence of denser than expected gas at the center of DEM L238 and DEM L249 might be explained if they resulted instead from CC explosions.  Dense circumstellar medium (CSM) has been found close to progenitors of CC explosions, including SN 1987A.  The CSM will be shocked early in a SNR's evolution, and at late times the shocked CSM gas will be near CIE conditions, or maybe even recombining \\citep{im89}. After the passage through the CSM, the blast wave will enter a bubble blown by the SN progenitor during its main-sequence stage, and then it will hit a dense swept-up ISM shell surrounding the tenuous bubble. This could explain the presence of dense, optically-emitting material in Figures \\ref{fig2} and \\ref{fig2a}. The problem with this hypothesis is that Fe is not the dominant nuclear burning product of CC explosions. The observed Fe overabundance would require the operation of a selective mechanism that would enhance emission from Fe and suppress emission from abundant elements such as O. This is not likely, and we reject a CC origin for DEM L238 and DEM L249. Another possibility involves an increased Fe abundance in the LMC gas itself, but there is no evidence for such large Fe overabundances in the LMC.  Perhaps Fe is produced in thermonuclear explosions, but the ejecta structure, the structure of the ambient medium, or both are quite different than in the standard Type Ia models \\citep{dwarkadas98,badenes03,badenes05,badenes06}.  Type Ia progenitors are still poorly understood, and there is growing evidence for the existence of peculiar Type Ia SNe.  A subluminous class of Type Ia SNe was recently discovered \\citep{jha06}, exemplified by SN 2002cx, with low expansion velocities and the presence of relatively dense partially burned or unburned material at their centers.  Dense, slow-moving Fe-rich gas is likely to be prominent in the interiors of remnants of such explosions, although detailed comparisons with observations of DEM L238 and DEM L249 must await a better understanding of this peculiar class of SNe. All five known members of the class of SN 2002cx exploded in blue, star-forming spiral galaxies, making them candidates for the prompt explosions suggested by \\citet{mdp06} and \\citet{sullivan06}. This and perhaps other classes of prompt Type Ia SNe must originate from more massive progenitors than standard, much older Type Ia SNe. Such progenitors are more likely to be surrounded by circumstellar material formed by ejection of the massive progenitor envelope prior to the explosion. Several Type Ia SNe have been found to be surrounded by dense CSM, including SN 2002ic \\citep{hamuy03}, although such SNe are infrequent. Among remnants of historical supernovae, Kepler's SNR shows evidence for dense circumstellar material \\citep[e.g.,][]{blair05}, and \\citet{kinugasa99} suggested a Type Ia origin for this remnant based on a high Fe abundance derived from its {\\it ASCA} X-ray spectrum.  Finally, if the companion to a white dwarf was a massive asymptotic giant branch star at the time of the white dwarf explosion, its entire envelope could have been stripped by SN ejecta \\citep{marietta00}. At late stages of SNR evolution, an Fe-contaminated envelope is expected to appear as dense Fe-rich gas close to the center of the SNR.  An asymmetric morphology is expected in this case.  \\citet{chuken88} associated DEM L238 and DEM L249 with old (Population II) stars and Type Ia explosions, because a CC origin for these SNRs is unlikely based on their large projected distances from OB associations in the LMC.  But $\\sim 10^8$ yr old Type Ia progenitors belonging to the young stellar population would be too old to be associated with the much younger ($\\la 10^7$ yr) massive stars in OB associations. Cepheid variables are better tracers of stars with such ages, and their spatial distribution is not correlated with OB associations and is only very marginally correlated with older star clusters \\citep{batefr99}. DEM L238 and DEM L249 are located in the LMC bar with a high Cepheid density \\citep{alcock99}. It is possible that their progenitors were born in the same star formation episode as Cepheids, although older progenitors belonging to Population I (or even to Population II) cannot be excluded, as this region of the LMC consists of a mixture of stellar populations.  DEM L238 and DEM L249 might belong to the same SNR population as young Fe-rich SNRs N103B and Kepler. The circumstellar gas in optically emitting N-rich radiative shocks in Kepler is dense; \\citet{blair91} estimated preshock densities in the 100 -- 1000 cm$^{-3}$ range. This gas traces the densest parts of the clumpy circumstellar material, where radiative cooling is important. The primary non-radiative blast wave propagates through the bulk of the strongly asymmetric CSM, with preshock densities as high as $\\sim 10$ cm$^{-3}$ \\citep{blair91}. Optically emitting radiative shocks with high (180 cm$^{-3}$) preshock densities are also present in N103B \\citep{russell90}. Unlike in Kepler, the N abundance is similar to its ISM value, but densities are much higher than in the surrounding \\ion{H}{2} region N103 or in the general diffuse ISM. Such high densities are more typical of molecular clouds, but there are no known nearby molecular clouds associated with N103B \\citep{fukui99,mizuno01}. These high densities and the high X-ray luminosity of N103B suggest the presence of dense circumstellar gas, although an ISM origin for the ambient gas cannot be ruled out at this time. N103B may be associated with the nearby double star cluster NGC 1850. The main cluster is 50 Myr old, but the second cluster is very young \\citep[2--5 Myr;][]{calcas98}. Optical observations revealed an enhancement in the N103B H$\\alpha$ emission at the radial velocity of the \\ion{H}{2} region N103 produced by young stars in NGC 1850 and its surroundings \\citep{chuken88}. \\citet{chuken88} interpreted this as evidence for a blast wave interaction with ambient ISM gas, but an alternative explanation involves the CSM around a moderately old (50 Myr) SN progenitor with a radial velocity comparable to its parent cluster (and hence the \\ion{H}{2} region N103). Just like Kepler, N103B is strongly asymmetric at all wavelengths, with a compact and bright western hemisphere and a large and much fainter eastern hemisphere. The asymmetry in the ambient CSM has been invoked to explain the asymmetry of the Kepler SNR, and a similar explanation is likely for N103B. At an age of $\\sim 10^4$ yr, both Kepler and N103B could still retain an imprint of the dense ambient CSM, and resemble DEM L238 and DEM L249.  In summary, {\\it Chandra} and {\\it XMM-Newton} observations revealed Fe-rich gas at the centers of LMC SNRs DEM L238 and DEM L249. While the Fe overabundance suggests Type Ia progenitors, the standard Type Ia models cannot explain the presence of relatively dense SN ejecta with long ionization timescales.  At this time we do not understand the origin of these ejecta; one possibility involves prompt Type Ia explosions with progenitors more massive than average. The remnant of Kepler's SN and the young LMC SNR N103B may belong to this category of Type Ia SNRs.  Studies of remnants of such explosions in the MCs promise to shed light on the nature of their poorly understood progenitors. The presence of such unusual progenitors may also force a reevaluation of the current chemical evolution models of the MCs, in view of the prominent role of Type Ia explosions in their evolution."
    },
    "0608/astro-ph0608068_arXiv.txt": {
        "abstract": "We investigate the effect of the hot CNO cycle breakout reaction $^{15}$O($\\alpha$,$\\gamma$)$^{19}$Ne on the occurrence of type I X-ray bursts on accreting neutron stars.  For $\\frp \\lesssim 0.1$, where $\\frp$ is a dimensionless factor by which we multiply the $^{15}$O($\\alpha$,$\\gamma$)$^{19}$Ne reaction rate of \\citet{CF88}, our model predicts that bursts should occur only for accretion rates $\\dot{M}$ below a critical value $\\approx 0.3 \\dot{M}_{\\mathrm{Edd}}$. This agrees with observations.  For larger values of $\\frp$, including the standard choice $\\frp = 1$, the model switches to a new regime in which bursts occur all the way up to $\\dot{M} \\approx \\dot{M}_{\\mathrm{Edd}}$.  Since the latter regime disagrees with observations, we suggest that the true $^{15}$O($\\alpha$,$\\gamma$)$^{19}$Ne reaction rate is lower than usually assumed. ",
        "introduction": "Type I X-ray bursts are thermonuclear explosions that occur on accreting neutron stars in low-mass X-ray binaries. They are triggered by thermally unstable H or He burning near the stellar surface \\citep[for reviews, see][]{C04,SB03}.  For systems with accretion rates $\\dot{M} \\lesssim 0.1 \\dot{M}_{\\mathrm{Edd}}$, where $\\dot{M}_{\\mathrm{Edd}}$ denotes the mass accretion rate at which the accretion luminosity is equal to the Eddington limit, the basic physics of the burst onset is well understood to be that of the thin shell thermal instability \\citep{SH65,HvH75}, and theoretical models have been rather successful at reproducing the gross characteristics of burst observations in this regime \\citep[e.g.,][]{FHM81,FL87,CB00,NH03}.  This {\\em Letter} addresses a longstanding problem that afflicts nearly all burst models when the accretion rate $\\dot{M} \\gtrsim 0.1 \\dot{M}_{\\mathrm{Edd}}$.  Both simple one-zone burst models \\citep{FHM81,P83,B98,HCW05} and sophisticated time-dependent multi-zone models \\citep{AJ82,TWL96,FHLT03,HCW05} predict that bursts should occur for all $\\dot{M}$ up to $\\approx \\dot{M}_{\\mathrm{Edd}}$. Observations, however, indicate that bursts do not occur for $\\dot{M} \\gtrsim 0.3 \\dot{M}_{\\mathrm{Edd}}$ \\citep{vPCLJ79,vPPL88,Cetal03,RLCN06}.  Furthermore, \\citet{vPPL88} found that, for $0.1 \\lesssim \\dot{M}/\\dot{M}_{\\mathrm{Edd}} \\lesssim 0.3$, a significant fraction of the accreted plasma burns stably between consecutive bursts, leading to large values $\\gtrsim 1000$ of the parameter $\\alpha$, the accretion energy released between successive bursts divided by the nuclear energy released during a burst.  Most theoretical models, on the other hand, predict that nearly all of the accreted matter burns unstably during bursts, giving a nearly constant $\\alpha < 100$ at all $\\dot{M}$.  \\citet[][hereafter NH03]{NH03} developed a global linear stability analysis of the accreted plasma on the surface of a neutron star. They discovered a new regime of unstable nuclear burning for $0.1 \\lesssim \\dot{M}/\\dot{M}_{\\mathrm{Edd}} \\lesssim 0.3$ that they referred to as ``delayed mixed bursts.''  Their model reproduced both the occurrence of considerable stable burning preceding a burst (leading to values of $\\alpha \\gtrsim 1000$) and the absence of bursts for $\\dot{M} \\gtrsim 0.3 \\dot{M}_{\\mathrm{Edd}}$, in agreement with observations.  However, the complexity of their model made it difficult to understand the basic physics behind delayed mixed bursts and to identify the reasons why their results differed from those of other theoretical models at high accretion rates.  To remedy this, \\citet[][hereafter CN06]{CN06} constructed a simple two-zone model that helped elucidate the physics of delayed mixed bursts.  They showed that the competition between nuclear heating due to triple-$\\alpha$ reactions and hot CNO cycle H burning on the one hand, and radiative cooling via photon diffusion and emission on the other hand, drives an overstability that eventually triggers a thin-shell thermal instability and hence a delayed mixed burst.  They asserted that H burning via the temperature-independent hot CNO cycle, augmented by the extra seed nuclei produced from stable He burning, significantly lowers the temperature sensitivity of the total nuclear energy generation rate to such an extent as to suppress the thin-shell thermal instability for $\\dot{M} \\gtrsim 0.3 \\dot{M}_{\\mathrm{Edd}}$.  Therefore, H burning via the hot CNO cycle is ultimately responsible for the lower critical $\\dot{M}$ above which bursts do not occur in nature.  The above argument perhaps explains why one-zone models fail, since the models generally focus only on He burning and make large approximations with respect to H burning.  But why do detailed time-dependent multi-zone models with large reaction networks also perform poorly in relation to observations?  These models are much more sophisticated than the models of NH03 and CN06 and consequently ought to perform better, whereas in fact the latter models agree much better with observations.  CN06 hypothesized that the time-dependent multi-zone burst models may have used too large a rate for the experimentally poorly-constrained hot CNO cycle breakout reaction $^{15}$O($\\alpha$,$\\gamma$)$^{19}$Ne \\citep{W69,WW81,LWFG86,WGS99,FGWD06}.  We test this hypothesis.  We begin in \\S \\ref{themodel} with a description of the model, and we present the results of the model in \\S \\ref{results}.  We discuss the results in \\S \\ref{discussion}, and we conclude in \\S \\ref{conclusions}. ",
        "conclusions": "Using the global linear stability analysis of \\citet{CN05}, which is an expanded and improved version of the model of NH03, we have investigated the effects of the hot CNO cycle breakout reaction $^{15}$O($\\alpha$,$\\gamma$)$^{19}$Ne on the occurrence of type I X-ray bursts on accreting neutron stars at high accretion rates.  For low values of $\\frp \\lesssim 0.1$, where $\\frp$ is a dimensionless factor by which we multiply the $^{15}$O($\\alpha$,$\\gamma$)$^{19}$Ne reaction rate of CF88, the hot CNO cycle breakout reaction slightly lowers the critical accretion rate above which delayed mixed bursts occur, but otherwise the breakout reaction has little effect on the burst onset.  The predictions of these models are in good agreement with observations.  For $\\frp \\gtrsim 0.1$, a new regime of prompt mixed bursts appears at accretion rates above the rates at which delayed mixed bursts occur, and the bursts survive up to roughly the Eddington limit.  The existence of this prompt mixed burst regime up to the Eddington limit is consistent with nearly all previous theoretical models.  Our results support the hypothesis of CN06 that the discrepancies between the results of the global linear stability analysis of NH03 and the results of time-dependent multi-zone models with large reaction networks may be caused by the latter models assuming too large a strength for the $^{15}$O($\\alpha$,$\\gamma$)$^{19}$Ne reaction rate at temperatures $T \\lesssim 6 \\times 10^{8}$ K.  The fact that observations agree much better with the results of NH03 and CN06 implies that the true reaction rate is lower than the rate assumed in these multi-zone models.  Specifically, we suggest that the true $\\alpha$ width $\\Gamma_{\\alpha}$ of the 4.03 MeV state in $^{19}$Ne is lower than the $\\alpha$ width proposed by \\citet{LWFG86}.  Calculations using multi-zone models such as those of \\citet{Wetal04} and \\citet{FGWD06}, but with the $^{15}$O($\\alpha$,$\\gamma$)$^{19}$Ne reaction rate lowered in strength, need to be carried out to either confirm or refute this suggestion."
    },
    "0608/astro-ph0608542_arXiv.txt": {
        "abstract": "{Over an interval of 120 years, the extraordinary object \\object{FG\\,Sge} has been transformed from a hot post-AGB star to a very luminous cool supergiant. Theoretically, this evolution has been associated with the reignition of a helium-shell during the post-AGB stage. A series of studies of the chemical composition of the photosphere have suggested that this evolution has been associated with a dramatic increase of approximately 3 dex in the abundances of s-process elements between about 1960 and 1995. The problem with this apparent change is that it occurred at a time when the surface convection zone, which is governed by the star's effective temperature, could not have developed sufficiently to dredge processed material from the stellar interior to the surface. We have reviewed the chemical evolution of \\object{FG\\,Sge} by means of modelling the time-varying spectrum under a range of assumptions. By comparing these models with published observational data, a self-consistent picture has emerged. In particular we find that surface hydrogen has been depleted during the interval in question. In contrast, the s-process abundances have generally maintained a steady enhancement of around $1 - 2$ dex, although some modest changes may have occurred since 1950. This implies that \\object{FG\\,Sge} has not just completed dredging up freshly-produced s-process isotopes. {However, there remains a contradiction between the observed hydrogen-depletion, the age of the associated planetary nebula, and current evolutionary models for a pre-white dwarf suffering a late thermal pulse.} ",
        "introduction": "During the last century, \\object{FG\\,Sagittae} evolved from a hot central star of a planetary nebulae to become an F-type supergiant within approximately 90 years. Only two comparable objects are known; \\object{V4334\\,Sgr} (Sakurai's object) covered a similar evolution in approximately five years to become, in 1999,  a C-type giant smothered by its own dust shell (Duerbeck et al.\\ \\cite{Due02}), and \\object{V605\\,Aql} which, after a brief foray into the giant domain during the 1920's, returned to the high-temperature r\\'egime as a typical [WC]-type planetary nebula central star (Clayton \\& de Marco \\cite{Cla97}). The rarity and brevity of such events arising from a cataclysm in the final evolution of low-mass stars is such that they are of profound interest.  For these reasons, every observation of these stars is important, not so much for their contemporary significance, but for their contribution to retrospective studies of evolutionary episodes which, even for such astronomically rapid events, may last longer than most scientists' careers. In particular, the development of increasingly sophisticated analytical methods render the original data of primary importance in any historical interpretation.  The generally accepted interpretation for the redward evolution of \\object{FG Sge} is that of a late thermal pulse in a young white dwarf or planetary nebula central star (Paczynski \\cite{Pac71}; Langer, Kraft \\& Anderson \\cite{Lan74}; Sackmann \\& Despain \\cite{Sac74}; Sch\\\"onberner \\cite{Sch79}; Iben  \\cite{Ibe84}; Bl\\\"ocker \\& Sch\\\"onberner \\cite{Blo97}; Gonzalez et al. \\cite{Gon98}; Lawlor \\& MacDonald \\cite{Law03}) wherein nuclear helium-burning at the surface of the electron-degenerate carbon-oxygen core is reignited. Detailed models for these late thermal pulses have been generated by several investigators (cf.\\ Sch\\\"onberner \\cite{Sch79}; Iben \\cite{Ibe84}; Bl\\\"ocker \\cite{Blo95}), but the main features are principally the same. The energy released by the helium-shell flash forces the outer layers to expand in about 50 to a few 100 years, depending on their mass. Hydrogen-rich, helium-rich and carbon-rich material is mixed by convective overshoot from the helium-shell flash, but the surface of the star retains its pre-pulse composition. A major component of this mixed region will be a supply of newly-produced s-process elements generated by the mixing of \\iso{13}{C} and protons.  As the star expands, its surface cools allowing surface convection to occur. As the temperature bottoms-out, the convection zone digs deeper, dredging processed material to the stellar surface.  The evolution of \\object{FG Sge} in terms of effective temperature (\\Teff) and luminosity ($L$) has been established from long-term photometric studies (van Genderen \\& Gautschy \\cite{vGG95}). These data appear to support the late thermal pulse hypothesis (Bl\\\"ocker \\& Sch\\\"onberner \\cite{Blo97}; Herwig \\cite{Her01a}; Lawlor \\& MacDonald \\cite{Law03}). Further evidence is the apparent enhancement of s-process elements in the photosphere of \\object{FG\\,Sge} (Langer et al.\\ \\cite{Lan74}) as \\Teff\\ dropped below 8\\,500\\kelvin, almost exactly the point where surface convection is expected to develop (cf.\\ van Genderen \\& Gautschy \\cite{vGG95}). The dilemma, in simple terms, is that the theoretical models for late thermal pulse evolution do not predict the dredge-up of s-processed material to the surface until some considerable time after the star has reached its minimum effective temperature. Moreover, the dredge-up should also be associated with a significant change in the abundances of hydrogen, helium, carbon and oxygen. Iben (\\cite{Ibe84}) remarks that {\\it ``The heretical conclusion to which one appears to be forced is that FG\\,Sge has {\\it not} just completed dredging up freshly produced s-process isotopes. If this is so, then the fault must lie in the estimates of surface abundances.''}  Consequently, either the theoretical predictions or the observations concerning the surface abundance of late thermal pulse products is wrong, or \\object{FG\\,Sge} is not a late thermal pulse product.  More recently, a variant of the late thermal pulse phenomenon has been identified (Herwig et al. \\cite{Her99})\\footnote{In fact, both classes of thermal pulse were previously encountered, but not explicitly labelled, by Sch\\\"onberner (\\cite{Sch79}). }. The ``late thermal pulse'' (LTP) model introduced above does qualitatively well in reproducing the hydrogen-deficiency and high carbon abundances seen in hydrogen-deficient planetary nebulae central stars, but fails to reproduce the large oxygen abundances found therein. Herwig et al. (\\cite{Her99}) found, however, that a ``very late thermal pulse'' (VLTP), occurring after the H-burning shell has become inactive, was able to produce a  surface abudance pattern in good general agreement with that observed in the hydrogen-deficient post-AGB stars, primarily because he considered convective overshoot. The two models are different in one major respect. In the first, the convection zone produced in the He-rich shell at He ignition cannot extend into the hydrogen-rich envelope (Iben \\cite{Ibe76}). In the second, this convection zone is able to penetrate into the H-rich surface layers and transport protons downwards (Herwig et al. \\cite{Her99}). What is not immediately apparent from the published figures, although it is stated in the text, is that the nuclear-driven convection zones are able to propagate through to the surface before the star cools sufficiently to develop a deep opacity-driven convection layer at its surface. This means that in the VLTP model, the disappearance of surface hydrogen and the appearance of highly-processed material can occur quickly and while the effective temperature is still significantly high.  Another major difference between the LTP and VLTP models is that the latter occurs on a much shorter timescale than the former. A VLTP occurs when the envelope mass is close to $\\sim 10^{-4}$\\Msolar. Since the nuclear energy produced by protons ingested into the intershell region is enormous ($\\sim10^{39}$ erg in less than a week), expansion following a VLTP occurs on the very small thermal timescale ($\\sim1$ y) of the tiny hydrogen-rich envelope (Herwig et al. \\cite{Her97}). In contrast, a LTP occurs when the envelope mass is also $\\sim 10^{-4}$\\Msolar, but $\\sim 10^{-2}$\\Msolar of hydrogen-free intershell must also be lifted, there is no proton-ingestion, and the expansion energy comes from the He-shell flash alone. Consequently the expansion occurs on a much longer timescale ($\\sim100$ y).  The VLTP model is therefore important for understanding the observed abundance evolution of FG\\,Sge and V4334\\,Sgr (Gonzalez et al. \\cite{Gon98}, Asplund et al. \\cite{Asp97}). For example, it has been stated by several authors that, because of their respective evolution timescales, FG\\,Sge and V4334\\,Sgr are examples of LTP and VLTP evolution respectively ({\\it cf.} Herwig \\cite{Her01a}).    It is the intent of this investigation to determine whether observation can indeed be reconciled with theory, and whether FG\\,Sge has suffered either a ``late'' or a ``very late'' thermal pulse. The next sections give a more detailed discussion of specific observations, followed by a presentation of the model atmospheres and spectra used and our subsequent analysis. It considers in more detail ideas introduced by Sch\\\"onberner \\& Jeffery (2003).  The literature on \\object{FG\\,Sge} is extensive, there are currently over 394 citations in the {\\sc simbad} database between 1960 and the present (2006 July)\\footnote{There is only one citation in SIMBAD from before 1960: Hoffmeister's (\\cite{Hof44}) list of 171 new variables.}. It would be impossible and inappropriate to survey critically every one of these articles, but inadvertent omissions of material that others might consider seminal will no doubt have been made. However we recognise the contribution of all of these publications, and hence our prefatory remarks on the importance of primary data. These derive from the fact that much early spectroscopic material is no longer accessible to us. Consequently, any conclusions which run contrary to the original authors' are legitimately disputable. On the other hand, if it can be shown that a self-consistent picture of \\object{FG\\,Sge} could be realised by a reassessment of these data, this paper will have achieved its goal.   \\begin{figure} \\epsfig{file=075fig01.ps,width=85mm} % \\caption[\\object{FG\\,Sge} evolution]{The evolution of \\object{FG\\,Sge} in the $L$--\\Teff\\ diagram (van Genderen \\& Gautschy \\cite{vGG95}, Table 3: high \\Teff ). A smooth curve has been drawn through the data.} \\label{f:evol} \\end{figure}   \\begin{figure} \\epsfig{file=075fig02.ps,width=85mm} % \\caption[\\object{FG\\,Sge} \\Teff]{The run of \\te\\ for \\object{FG\\,Sge} as determined spectroscopically and photometrically. Both low ($\\times$) and high (dots) photometric \\Teff\\ scales are from van Genderen \\& Gautschy (\\cite{vGG95}: vGG95). The spectroscopic measurements are from Herbig \\& Boyarchuk (\\cite{Her68}: Her68), Langer et al.\\ (\\cite{Lan74}: Lan74), Stone (\\cite{Sto79}: Sto79), Acker et al.\\ (\\cite{Ack82}: Ack82), Montesinos et al.\\ (\\cite{Mon90}: Mon90) and Gonzalez et al.\\ (\\cite{Gon98}: Gon98). New results from the present study are also shown (Jef03). This and other figures appear in colour in the online version of this paper; parenthetical references to colour refer to the latter. } \\label{f:teff} \\end{figure} ",
        "conclusions": "We have made an extensive review of the chemical evolution of the surface of \\object{FG\\,Sge} between approximately 1960 and 1990, an epoch during which there have been published reports of dramatic changes in the abundances of many elements including particularly those created by s-process nucleosynthesis in the stellar interior. In several cases, these various reports have been shown to be seriously inconsistent with one another and with stellar evolution theory. Although it is more usual to recast the theory when observation and theory come into conflict, the observational inconsistencies are sufficiently strong that they must be examined closely.  Although we have not had direct access to the majority of the raw spectrograms, we have been able to assess the reliability of the majority of key studies, particularly those by Herbig \\& Boyarchuk (\\cite{Her68}), Langer et al.\\ (\\cite{Lan74}), Cohen \\& Phillips (\\cite{Coh80}), and Gonzalez et al.\\ (\\cite{Gon98}). Although the first three of these report some changes in the abundances of carbon and s-process elements, we find that the data presented are almost entirely consistent with a constant surface composition typical of other post-AGB stars, i.e.\\ with \\mbox{C/O $>1$} and s-process abundances enhanced by $\\sim 1 \\mbox{-} 2$ dex. The exception that the earliest data for Y and Zr indicate an increase from normal abundance to $\\sim +1$ dex remains a mystery. The very dramatic increases in s-process abundances reported in the most recent study (Gonzalez et al.\\ \\cite{Gon98}) are demonstrably wrong. They are due to the adoption of a value for \\Teff\\ which is 1\\,000\\kelvin\\ too high. A reassessment of the calibration of the photometric evolution from ultraviolet data by Montesinos et al.\\ (\\cite{Mon90}) demonstrates that the ``high'' \\Teff\\ calibration by van Genderen \\& Gautschy (\\cite{vGG95}) is most consistent with all of the photometric and spectroscopic data. Since \\object{FG\\,Sge} has now maintained its current  \\Teff\\ of around 5\\,500\\kelvin for over a decade, this may represent the cool limit in its redward evolution.   Supporting Iben (\\cite{Ibe84}), we have argued that the current surface of \\object{FG\\,Sge} does not show evidence of freshly dredged-up material in the form of carbon or s-processed elements. Contrary to Sch\\\"onberner \\& Jeffery (\\cite{Sch03}), we find that surface hydrogen has been depleted in the recent past ($\\sim40$ y). In the 1960's it had an apparently normal H content (cf. Fig. 6), whereas in 1994 there is evidence of H-depletion via a reduced background opacity (cf. Fig. 12); a conclusion based on a partial analysis of the high-resolution spectrum of Gonzalez et al. (\\cite{Gon98}) not earlier at our disposal. We do not see evidence of the very prompt mixing suggested by recent VLTP theory.  Stellar evolution theory predicts that, although a convective envelope does form as the star cools following a late thermal pulse, a deep convective envelope capable of dredging up processed material will not form until some time after the star has completed its redward evolution. Therefore, the observed hydrogen depletion, suggesting prompt envelope mixing following a VLTP, contradicts the age derived from the  associated planetary nebula, which suggests an LTP. The question is then how well the physics of the LTP and VLTP are understood, and whether the demarcation between them is as concrete as current models suggest.  \\object{FG\\,Sge} is one of those very rare objects which demonstrates stellar evolution ``while-you-wait''. As a very evolved star, its surface already shows some material which has been created by nucleosynthesis in the deep interior. Astero-archaeology is the science of sifting through such material for clues to the past evolution of the star. In this case, most of the material was brought to the surface when the star had a deep convective envelope on the asymptotic giant branch. However, it seems that \\object{FG\\,Sge} has recently entered a new period of excavation, which has recently brought fresh helium to the surface. Other elements will be added to this within the next hundred or so years, when it will be possible to observe fresh material being brought to the surface as the star begins to dig up its own past. Continuing observations will provide key tests for stellar evolution models, particularly for models of late and very late thermal pulses."
    },
    "0608/astro-ph0608404_arXiv.txt": {
        "abstract": "We report the finding of an unusual, weak precursor to a thermonuclear X--ray burst from the accreting millisecond pulsar SAX J1808.4--3658. The burst in question was observed on Oct. 19, 2002 with the Rossi X-Ray Timing Explorer (RXTE) proportional counter array (PCA). The precursor began $\\approx 1$ s prior to the onset of a strong radius expansion burst, lasted for about 0.4 s, and exhibited strong oscillations at the 401 Hz spin frequency. Oscillations are not detected in the $\\approx 0.5$ s interval between the precursor and the main burst. The estimated peak photon flux and energy fluence of the precursor are about 1/25, and 1/500 that of the main burst, respectively. From joint spectral and temporal modeling, we find that an expanding burning region with a relatively low temperature on the spinning neutron star surface can explain the oscillations, as well as the faintness of the precursor with respect to the main part of the burst.  We dicuss some of the implications of our findings for the ignition and spreading of thermonuclear flames on neutron stars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/hep-th0608211_arXiv.txt": {
        "abstract": "We investigate the vacuum expectation value of the surface energy-momentum tensor for a massive scalar field with general curvature coupling parameter obeying the Robin boundary conditions on two codimension one parallel branes in a $(D+1)$-dimensional background spacetime $\\mathrm{AdS}_{D_{1}+1}\\times \\Sigma $ with a warped internal space $\\Sigma $. These vacuum densities correspond to a gravitational source of the cosmological constant type for both subspaces of the branes. Using the generalized zeta function technique in combination with contour integral representations, the surface energies on the branes are presented in the form of the sum of single brane and second brane induced parts. For the geometry of a single brane both regions, on the left and on the right of the brane, are considered. At the physical point the corresponding zeta functions contain pole and finite contributions. For an infinitely thin brane taking these regions together, in odd spatial dimensions the pole parts cancel and the total zeta function is finite. The renormalization procedure for the surface energies and the structure of the corresponding counterterms are discussed. The parts in the surface densities generated by the presence of the second brane are finite for all nonzero values of the interbrane separation and are investigated in various asymptotic regions of the parameters. In particular, it is shown that for large distances between the branes the induced surface densities give rise to an exponentially suppressed cosmological constant on the brane. The total energy of the vacuum including the bulk and boundary contributions is evaluated by the zeta function technique and the energy balance between separate parts is discussed. ",
        "introduction": "\\label{sec:introd}  The braneworld scenario provides an interesting alternative to the standard Kaluza-Klein compactification of the extra dimensions. The simplest phenomenological models describing such a scenario are the five-dimensional Randall-Sundrum type braneworld models (for a review see \\cite{Ruba01,Maar03}% ). From the point of view of embedding these models into a more fundamental theory, such as string/M-theory, one may expect that a more complete version of the scenario must admit the presence of additional extra dimensions compactified on an internal manifold. From a phenomenological point of view, the consideration of more general spacetimes offer a richer geometrical structure and may provide interesting extensions of the Randall-Sundrum mechanism for the geometric origin of the hierarchy. More extra dimensions also relax the fine-tunings of the fundamental parameters. These models can provide a framework in the context of which the stabilization of the radion field naturally takes place. In addition, a richer topological structure of the field configuration in transverse space provides the possibility of more realistic spectrum of chiral fermions localized on the brane. Several variants of the Randall-Sundrum scenario involving cosmic strings and other global defects of various codimensions have been investigated in higher dimensions (see, for instance, \\cite{Greg00} and references therein).  Motivated by the problems of the radion stabilization and the generation of cosmological constant, the role of quantum effects in braneworlds has attracted great deal of attention \\cite{Fabi00}-\\cite{Saha06c}. A class of higher dimensional models with the topology $\\mathrm{AdS}_{D_{1}+1}\\times \\Sigma $, where $\\Sigma $ is a one-parameter compact manifold, and with two branes of codimension one located at the orbifold fixed points, is considered in Refs. \\cite{Flac03b,Flac03}. In both cases of the warped and unwarped internal manifold, the quantum effective potential induced by bulk scalar fields is evaluated and it has been shown that this potential can stabilize the hierarchy between the Planck and electroweak scales without fine tuning. In addition to the effective potential, the investigation of local physical characteristics in these models is of considerable interest. Local quantities contain more information on the vacuum fluctuations than the global ones and play an important role in modelling a self-consistent dynamics involving the gravitational field. In the previous papers \\cite% {Saha06a,Saha06b} we have studied the Wightman function, the vacuum expectation values of the field square and the energy-momentum tensor for a scalar field with an arbitrary curvature coupling parameter obeying Robin boundary conditions on two codimension one parallel branes embedded in the background spacetime $\\mathrm{AdS}_{D_{1}+1}\\times \\Sigma $ with a warped internal space $\\Sigma $. For an arbitrary internal space $\\Sigma $, the application of the generalized Abel-Plana formula \\cite{SahaRev} allowed us to extract form the vacuum expectation values the part due to the bulk without branes and to present the brane induced parts in terms of exponentially convergent integrals for the points away from the branes.  The braneworld corresponds to a manifold with boundaries and the physical quantities, in general, receive both volume and surface contributions. In particular, the contributions located on the visible brane are of special interest as they are direct observables in the theory. In Ref.~\\cite% {Saha04surf} the vacuum expectation value of the surface energy-momentum tensor is evaluated for a massive scalar field subject to Robin boundary conditions on two parallel branes in $(D+1)$-dimensional AdS bulk (for surface densities induced by quantum fluctuations of a bulk scalar field on a single brane in Rindler-like spacetimes see \\cite{Saha06c}). It has been shown that for large distances between the branes the induced surface densities give rise to an exponentially suppressed cosmological constant on the brane. In the Randall-Sundrum braneworld model, for the interbrane distances solving the hierarchy problem between the gravitational and electroweak mass scales, the cosmological constant generated on the visible brane is of the right order of magnitude with the value suggested by the cosmological observations. In the present paper we investigate similar issues within the framework of higher dimensional braneworld models with warped internal spaces.  The paper is organized as follows. In the next section we describe the corresponding bulk geometry and outline the structure of the eigenmodes and the Kaluza-Klein spectrum. An integral representation is constructed for the zeta function related to the vacuum expectation value of the surface energy-momentum tensor. This function is presented in the form of the sum of single brane and second brane induced parts. The analytic continuation of the zeta function for the model with a single brane is done in Section \\ref% {sec:1brane}. The surface energy-momentum is investigated on both sides of the branes and the renormalization procedure is discussed. In Section \\ref% {sec:2brane} we study the surface energy density induced by the presence of the second brane. Various limiting cases are considered and the cosmological constant induced on the brane is estimated. The Section \\ref{sec:enbal} is devoted to the evaluation of the total vacuum energy in the two-brane setup on the base of the zeta function technique. The relation between the bulk and surface energies is discussed. The last section contains a summary of the work. ",
        "conclusions": "\\label{sec:Conc}  Continuing our previous work \\cite{Saha06a,Saha06b}, in the present paper we have investigated the expectation value of the surface energy-momentum tensor induced by the vacuum fluctuations of a bulk scalar field with an arbitrary curvature coupling parameter satisfying Robin boundary conditions on two parallel branes in background spacetime $\\mathrm{AdS}_{D_{1}+1}\\times \\Sigma $ with a warped internal space $\\Sigma $. Vacuum stresses on the brane are the same for both subspaces and the energy-momentum tensor on the brane corresponds to the source of the cosmological constant type in the brane universe. It is remarkable that the latter property is valid also for the more general model with the mertic $g_{\\mu \\sigma }$ instead of $\\eta _{\\mu \\sigma }$ in line element (\\ref{metric}). As an regularization procedure for the surface energy density we employ the zeta function technique. By using the residue theorem we have constructed an integral representations for the partial zeta function corresponding to a given KK mode along the internal subspace. This function is presented as the sum of single brane and second brane induced parts [see formula (\\ref{zet12})]. The latter is finite at the physical point and the further analytical continuation is necessary for the first term only. We have done this in Section \\ref{sec:1brane}. As the first step we subtract and add to the integrand the leading terms of the corresponding asymptotic expansion for large values of the argument and explicitly integrate the asymptotic part. Further, for the regularization of the sum over the modes along the internal space we use the local zeta function related to these modes. By making use of the formula for the pole structure of this function, we have presented the energy density on a single brane as the sum of the pole and finite parts.These separate parts are determined by formulas (\\ref{phi2Lpf1}), (\\ref% {phi2Lpf2}). The pole parts in the surface energy density are absorbed by adding to the brane action the counterterms having the structure given by Eq. (\\ref{Counterterms}). The renormalized energy density on the corresponding surface of a single brane is determined by formula (\\ref% {epsjren}), where the second term on the right presents the finite renormalization part. The coefficients in this part cannot be determined within the model under consideration and their values should be fixed by additional renormalization conditions which relate them to observables.  Unlike to the single brane part, the surface energy density induced by the presence of the second brane contains no renormalization ambiguities and is investigated in Section \\ref{sec:2brane}. This part is given by formula (\\ref% {emt2pl3}), or equivalently by formula (\\ref{Deltepsjnew}). We have investigated the induced energy density in various asymptotic regions for the parameters of the model. In the limit $k_{D}\\rightarrow 0$ we have obtained the result for the geometry of two boundaries in the bulk $% R^{(D_{1}-1,1)}\\times \\Sigma $. For the modes along $\\Sigma $ with large KK masses, the induced energy density is exponentially small. In particular, for sufficiently small length scales of the internal space this is the case for all nonzero KK modes and the main contribution comes from the zero mode. For small interbrane distances, to the leading order, the induced energy density is given by formula (\\ref{epsjsmalldist}) and the contribution of the induced part dominates the single brane parts. When the right brane tends to the AdS horizon, $z_{b}\\rightarrow \\infty $, the induced energy density on the left brane vanishes as $e^{-2\\lambda _{\\beta }z_{b}}/z_{b}^{D_{1}/2}$ for the nonzero KK mode and like $% z_{b}^{-D_{1}-2\\nu }$ for the zero mode. In the same limit the corresponding energy density on the right brane behaves as $z_{b}^{D_{2}+D_{1}/2+1}e^{-2% \\lambda _{\\beta }z_{b}}$ for the nonzero KK mode and like $z_{b}^{D_{2}-2\\nu }$ for the zero mode. In the limit when the left brane tends to the AdS boundary the contribution of a given KK mode into the induced energy density vanishes as $z_{a}^{D+2\\nu }$ and as $z_{a}^{2\\nu }$ for the left and right branes, respectively. For small values of the AdS curvature radius corresponding to strong gravitational fields, for nonzero KK modes under the conditions $\\lambda _{\\beta }z_{a}\\gg 1$ and $\\lambda _{\\beta }(z_{b}-z_{a})\\gg 1$, the contribution to the induced energy density is suppressed by the factor $e^{-2\\lambda _{\\beta }(z_{b}-z_{a})}$. For the zero KK mode, the corresponding energy density integrated over the internal space behaves as $k_{D}^{D_{1}+1}\\exp [(D_{1}\\delta _{j}^{a}+2\\nu )k_{D}(a-b)]$ for the brane at $y=j$ and is exponentially small. In the model without the internal space we have similar behavior with $\\nu $ replaced by $\\nu _{1}$ and for a scalar field with $\\zeta <\\zeta _{D+D_{1}+1} $ the suppression is relatively weaker.  In the model under discussion the hierarchy between the fundamental Planck scale and the effective Planck scale in the brane universe is generated by the combination of redshift and large volume effects. The corresponding efective Newton's constant on the brane at $y=j$ is related to the higher-dimensional fundamental Newton's constant by formula (\\ref{GDj}) (see also Eq. (\\ref{Planckhierarchy}) for the ratio of the corresponding Planck masses)\\ and for large interbrane separations is exponentially small on the brane $y=b$. We show that this mechanism also allows obtaining a naturally small cosmological constant generated by the vacuum quantum fluctuations of a bulk scalar. For large interbrane distances the ratio of the induced cosmological constant to the the corresponding Planck scale quantity in the brane universe is estimated by formula (\\ref{LambDjest1}) \\ and is exponentially small. For the visible brane in the higher dimensional generalization of the Randall-Sundrum two-brane model, this ratio is given in terms of the effective and fundamental Planck masses by Eq.~(\\ref{Lambonb}% ). We have considered two classes of models with the compactification scale on the visible brane close to the fundamental Planck scale. For the first one the higher dimensional Planck mass and the AdS inverse radius are of the same order and in the second one a separation between these scales is assumed. In both cases the corresponding interbrane distances generating the hierarchy between the electroweak and Planck scales are smaller than those for the model without an internal space and the required suppression of the cosmological constant is obtained without fine tuning.  In Section \\ref{sec:enbal} we have considered the total vacuum energy in the region between the branes, evaluated as the sum of zero-point energies for elementary oscillators. We have shown that this energy is equal to the sum of surface and volume energies. The latter is evaluated as the integral of the bulk energy density. Similar to the case of the surface energy, for the evaluation of the total vacuum energy we have used the zeta function regularization technique. This energy is presented as the sum of the single brane and interference parts. We have extracted the pole parts of the single brane vacuum energies and have discussed the corresponding renormalization procedure. The divergences can be absorbed by adding the counterterms to the brane action which have the structure similar to that for the surface energy. The interference part in the vacuum energy is given by formula (\\ref% {DeltaEhigh}) and is regular for all nonzero interbrane distances. We have investigated the asymptotic behavior of this part in various limiting cases. For sufficiently small interbrane distances the total vacuum energy is dominated by the interference part given by formula (\\ref{DeltEsmalldist}). In particular, from this formula it follows that in the case $\\sigma =1$ any minimum of the vacuum energy is local and the fixation of the interbrane distance by the corresponding effective potential is stable only locally. Further, we have shown that the induced vacuum densities and vacuum effective pressures on the branes satisfy the standart energy balance equation (\\ref{enbalance}) with the inclusion of the surface terms."
    },
    "0608/astro-ph0608626_arXiv.txt": {
        "abstract": "We present the formalism and numerical results for torsional oscillations of relativistic stars endowed with a strong dipole magnetic field, assumed to be confined to the crust. In our approach, we focus on axisymmetric modes and neglect higher-order couplings induced by the magnetic field.  We do a systematic search of parameter space by computing torsional mode frequencies for various values of the harmonic index $\\ell$ and for various overtones, using an extended sample of models of compact stars, varying in mass, high-density equation of state and crust model.  We show that torsional mode frequencies are sensitive to the crust model if the high-density equation of state is very stiff. In addition, torsional mode frequencies are drastically affected by a dipole magnetic field, if the latter has a strength exceeding roughly $10^{15}$G and we find that the magnetic field effects are sensitive to the adopted crust model. Using our extended numerical results we derive empirical relations for the effect of the magnetic field on torsional modes as well as for the crust thickness. We compare our numerical results to observed frequencies in SGRs and find that certain high-density EoS and mass values are favored over others in the non-magnetized limit. On the other hand, if the magnetic field is strong, then its effect has to be taken into account in attempts to formulate a theory of asteroseismology for magnetars. ",
        "introduction": "\\label{sec:Intro}   There is growing evidence that Soft Gamma Repeaters (SGRs) could be magnetars experiencing starquakes that are connected (through the intense magnetic field) to gamma ray flare activity. Magnetars are thought to be neutron stars with very strong magnetic fields, greatly exceeding $\\sim 5 \\times 10^{13}$ G \\citep{DT1992}. The flare activity consist of an initial short peak in the hard part of the spectrum, followed by a decaying softer part (tail) which last for hundreds of seconds. Recent analysis of detailed observations from several SGRs \\citep{Israel2005,SW2005,WS2006} revealed that the decaying part of the spectrum exhibits a number of oscillations with frequencies in the range of a few tenths of Hz to a few hundred Hz. There are three events detected up to now which can be associated with crust oscillations of magnetars. The first event was detected already in 1979 from the source SGR 0526-66 \\citep{Mazets1979,Barat1983}, the second in 1998 from SGR 1900+14 \\citep{Hurley1999}, while the third and most energetic one was observed in December 2004 from the source SGR 1802-20 \\citep{Terasawa2005, Palmer2005}. Analysis of the tail oscillations of SGR 1806-20 revealed the presence of oscillations at approximately 18, 26, 29, 92.5, 150, 626.5, and 1837 Hz \\citep{Israel2005, WS2006,SW2006}, for SGR 1900+14 the detected frequencies were 28, 54, 84, and 155 Hz \\citep{SW2005}, while for the case of SGR 0526-66 there was only one frequency identified at 43.5 Hz \\citep{Barat1983}.   Magnetar starquakes may be driven by the evolving intense magnetic field, which accumulates stress and eventually leads to crust fracturing.  The excited seismic oscillations are known to be of two types: {\\em shear modes}, which are polar-type oscillations and {\\em torsional modes}, which are axial-type oscillations. The latter are thought to be more easily excited during a fracturing of the crust, since they only involve oscillations of the velocity, with the shear modulus of the crust acting as the necessary restoring force. Slowly rotating stars still maintain a nearly spherical shape and the velocity field of torsional oscillations is then divergence-free, with no radial component. Torsional modes ({\\em t}-modes) are labeled as $_\\ell t_n$, where $\\ell$ is the angular index, while the index $n$ corresponds to the number of radial nodes in the eigenfunctions of the overtones for a specific $\\ell$. Shear and torsional modes have been studied mainly in Newtonian theory, see e.g. \\citet{HC1980,McDermott1988,Carroll1986, Strohmayer1991a,Lee2006}, while there are only a few studies of torsional modes in general relativity \\citep{Schumaker1983, Leins1994, Messios2001,SA2006}.  In the Newtonian, non-magnetized limit \\citet{HC1980} found that the period $_\\ell P_0 $ of the fundamental torsional modes depends mostly on the radius of the star, $R$, the speed of shear waves, $u_s$, and the angular index $\\ell$, via \\begin{equation} {}_\\ell P_0 \\approx 2\\pi[\\ell(\\ell+1)]^{-1/2} R/u_s, % \\end{equation} where  $u_s=(\\mu/\\rho)^{1/2}$, with $\\mu$ and $\\rho$ being the shear modulus and the density, respectively. Furthermore, it was found that the period of overtones of index $n$ is essentially independent of the angular index $l$ (provided $l$ does not greatly exceed $n$) and is basically determined by the crust thickness $\\Delta r$: \\begin{equation} {}_\\ell P_n \\approx 2n^{-1}\\Delta r/u_s. \\end{equation} From the above two relations it follows that the relative crust thickness, $\\Delta r/R$, is given by \\begin{equation} \\frac{\\Delta r}{R} \\approx \\frac{\\pi n}{\\sqrt{\\ell (\\ell+1)}} \\frac{_\\ell f_0}{_\\ell f_n}, \\label{eq:mode_spacing} \\end{equation} which is independent of the details of the equation of state in the crust and basically only depends on the ratio of the frequency of the overtones to the fundamental frequency, see also \\citet{SA2006}. It thus becomes obvious that the successful identification of both the fundamental torsional frequencies and their overtones could allow for the determination of the crust thickness, which in turn could lead to information on the high-density part of the equation of state. In such a determination, of course, general relativistic and magnetic field effects on the frequencies of torsional modes should properly be taken into account, which is one of the aims of the present paper.  Relativistic effects have been found to significantly increase the fundamental $\\ell=2$ torsional mode period by roughly 30\\% for a typical $M=1.4M_\\odot$, $R=10{\\rm km}$ model. \\citet{Duncan1998} derived the following empirical formula for the period of the fundamental $\\ell=2$ torsional mode of a relativistic star \\begin{equation} _\\ell P_0 ({\\rm ms}) \\approx 33.6 R_{\\rm 10}\\frac{0.87+0.13M_{\\rm 1.4} R^{-2}_{\\rm 10}}{(1.71-0.71M_{\\rm 1.4} R^{-1}_{\\rm 10})^{1/2}}, \\label{eq:Duncan} \\end{equation} where $R_{\\rm 10}=R/(10 {\\rm km})$ and $M_{\\rm 10}=M/(1.4M_\\odot)$. In this relation, two effects have been taken into account, the (relativistic) gravitational redshift and the (Newtonian) acceleration of gravity $g=GM/R^2$ at the surface. More details can be found in \\cite{Duncan1998}, see also \\cite{Messios2001}. In addition, the formula (\\ref{eq:Duncan}) agrees well with the fully relativistic empirical relation found by \\citet{Messios2001} \\begin{equation} {}_{\\ell} P_0 ({\\rm ms}) \\approx 34 R_{\\rm 10}. \\end{equation} The latter is an extension of the Newtonian Hansen \\& Cioffi formula, based on a particular $1.4 M_\\odot$ model. On the other hand, the effect of a strong magnetic field, $B$, on torsional modes was considered by \\citet{Duncan1998} and \\citet{Messios2001}. Assuming that the shear modulus $\\mu$ is augmented by the magnetic field tension, $B^2/4\\pi$, one can derive the following estimate for the period of a torsional mode \\begin{equation} P \\approx P^{(0)}\\left[1+(B/B_\\mu)\\right]^{-1/2}, \\label{Pb} \\end{equation} where $P^{(0)}$ is the period in the limit of vanishing magnetic field and $B_\\mu=\\sqrt{4\\pi\\mu}$. However, \\citet{Messios2001} found that if the structure of the magnetic field is close to monopolar in the crust, the period of the fundamental $\\ell=2$ torsional mode does not follow the above relation. This underlines the importance of properly taking into account the structure of the magnetic field, so that Equation (\\ref{Pb}) must be generalized to particular magnetic field multipoles.  An essential ingredient in computing torsional mode frequencies is the assumption one makes about the shear modulus $\\mu$. Here, we adopt the zero-temperature limit of the approximate formula derived by \\citet{Strohmayer1991b} \\begin{equation} \\mu= 0.1194 \\frac{n_i (Ze)^2}{a}, \\label{eq:mu} \\end{equation} where $n_i$ is the ion number density, $a^3=3/4\\pi n_i$ is the average ion spacing, and $+Ze$ is the ion charge. Notice that this approximate formula for the shear modulus is derived on the assumption of a bcc crystalline state, which is averaged over all directions. For the detailed composition of the crust, we consider a modern EOS by \\citet{DH2001} (hereafter DH) and, for comparison, an older EOS by \\citet{NV1973} (hereafter NV). The two EOSs for the crust differ both in composition as well as in the value of the density at the base of the crust. As the crust thickness depends on the compactness of the star, we explore the allowed parameter space for the unknown high-density EOS by selecting several tabulated high-density EOSs that yield neutron star models with significantly different macroscopic properties.           In the present work, we derive the perturbation equations for torsional oscillations of magnetized stars in general relativity following the approach by \\citet{Messios2001} and neglecting spacetime perturbations (Cowling approximation). We specialize the magnetic field to a dipole structure, which allows us to reduce the system of equations to only one dimension. A large number of torsional modes (fundamental and overtones) is obtained for selected magnetar models differing in the high-density EOS, the crust EoS and mass. Our study is valid for moderate magnetic field strengths, up to roughly $10^{16}$ G, assuming that the magnetic field is confined to the crust and neglecting the possible presence of a thin fluid ocean.  The reason for this is twofold: on one hand, for significantly stronger magnetic fields the distortion of the equilibrium shape of the star as well the magnetic coupling to higher harmonics should also be taken into account; on the other hand, the Alfv\\'en velocity $u_A\\equiv B/(4\\pi\\rho)^{1/2}$ $\\approx 3\\times 10^7$ cm/s $B_{15} \\, \\rho_{14}^{-1/2}$ becomes comparable to the speed of shear waves in the crust when the magnetic field reaches values of 10$^{15}$-10$^{16}$G, which implies that global magnetosonic waves will be strongly coupled to shear waves and the torsional oscillations may no longer be confined to the crust. The effect of including global magnetosonic waves even at smaller magnetic field strengths should also be studied in more detail, see \\citep{GSA2006,SKSV06}.  The article is structured as follows: in the next two sections we describe the general-relativistic ideal MHD equations, the background stellar configuration and the way that we have derived the dipole geometry of the magnetic field. In the fourth section we derive the perturbation equations for torsional oscillations of magnetized stars in the Cowling approximation, while in the fifth section we present the obtained frequencies for selected magnetar models along with fitting formulae for the influence of the magnetic field in the oscillation spectrum. The article closes with a summary and discussion.  Unless otherwise noted, we adopt units of $c=G=1$, where $c$ and $G$ denote the speed of light and the gravitational constant, respectively, while the metric signature is $(-,+,+,+)$. ",
        "conclusions": "We have derived the formalism for computing torsional oscillations of relativistic stars endowed with a strong dipole magnetic field, confined to the crust. Our equations are valid in the Cowling approximation (no spacetime perturbations) and we have neglected the effect of the magnetic field on the equilibrium configuration and on the coupling of torsional modes to $\\ell \\pm 2$ terms and to global magnetosonic modes. Under these approximations, our formalism allows us to obtain an estimate of the magnetic field effects on torsional modes, up to moderate values of the magnetic field strength (up to a few times $B_\\mu=4 \\times 10^{15}$G).  These moderate values of the magnetic field strength are appropriate for models of magnetars, for which there is strong evidence that they are at the heart of the SGR phenomenon.  We have done a systematic search of parameter space by computing torsional mode frequencies for various values of the harmonic index $\\ell$ and for various overtones, using an extended sample of models of compact stars. These models vary in mass, high-density equation of state and crust model, uniformly covering the allowed mass vs. radius parameter space.  Our numerical results have shown that torsional mode frequencies are sensitive to the crust model if the high-density equation of state is very stiff (such as EoS L). In addition, torsional mode frequencies are drastically affected by a dipole magnetic field, if the latter has a strength exceeding roughly $10^{15}$G. The effect of the magnetic field is surprisingly sensitive to the adopted crust model. Using our extended numerical results we have derived empirical relations for the effect of the magnetic field on torsional modes as well as for the crust thickness.  We compare our numerical results to observed frequencies in SGRs and find that certain high-density EoS and mass values are favored over others in the non-magnetized limit. On the other hand, if the magnetic field is strong, then its effect has to be taken into account in attempts to formulate a theory of asteroseismology for magnetars. This topic, as well as the inclusion of global magnetosonic modes will be discussed in a separate publication \\citep{SKSV06}."
    },
    "0608/astro-ph0608556_arXiv.txt": {
        "abstract": "The Very Large Array was used to observe the ultracool rapidly rotating M9 dwarf TVLM 513-46546 simultaneously at 4.88 GHz and 8.44 GHz.  The radio emission was determined to be persistent, variable and periodic at both frequencies with a period of $\\sim 2$ hours.  This periodicity is in excellent agreement with the estimated period of rotation of the dwarf based on its $v$ sin $i$ of $\\sim 60$ km s$^{-1}$.  This rotational modulation places strong constraints on the source size of the radio emitting region and hence the brightness temperature of the associated emission.  We find the resulting high brightness temperature, together with the inherent directivity of the rotationally modulated component of the emission, difficult to reconcile with incoherent gyrosynchrotron radiation.  We conclude that a more likely source is coherent, electron cyclotron maser emission from the low density regions above the magnetic poles.  This model requires the magnetic field of TVLM 513-46546 to take the form of a large-scale, stable, dipole or multipole with surface field strengths up to at least 3kG.  We discuss a mechanism by which broadband, persistent electron cyclotron maser emission can be sustained in the low density regions of the magnetospheres of ultracool dwarfs.  A second non-varying, unpolarized component of the emission may be due to depolarization of the coherent electron cyclotron maser emission or alternatively, incoherent gyrosynchrotron or synchrotron radiation from a population of electrons trapped in the large-scale magnetic field. ",
        "introduction": "Magnetic activity in cool main sequence stars is usually diagnosed through the detection of coronal X-ray and radio emission, and via the emission from spectral lines such as chromospheric H$\\alpha$. In the Sun's case, it is accepted that such magnetic activity is powered by magnetic fields generated by a so-called $\\alpha\\Omega$ dynamo at the base of the convective layer where differential rotation is strongest \\citep{parker75}.  This dynamo mechanism is also thought to operate in a wide spectral range of cool stars from F-type to early M-type dwarfs and is particularly successful in accounting for the well-established rotation-activity relationship observed over this spectral range \\citep{pallavicini81, noyes84}.  \\par For low mass stars with spectral type $\\gtrsim$ M3, the nature of magnetic activity is less well understood.  Here lies the boundary at which low mass stars become fully convective, no longer possessing the radiative layer necessary to sustain the $\\alpha\\Omega$ dynamo.  However, M dwarf flare stars, very often with spectral type later than M3, are associated with intense levels of magnetic activity and are also characterized by large magnetic filling factors, with kG magnetic fields covering a large proportion of the stellar disk \\citep{saar94, johnskrull96}.  Indeed, the fraction of objects exhibiting emission associated with magnetic activity, particularly high levels of quiescent H$\\alpha$ emission, is seen to increase steadily for spectral types $\\gtrsim M3$ reaching 70\\% by about spectral type M8 \\citep{gizis00, west04}.  Recently, magnetic mapping of the stellar surface of an M4 dwarf has directly confirmed the presence of large-scale, axisymmetric, poloidal fields, including a strong dipolar component, despite the absence of surface differential rotation \\citep{donati06}.  Clearly an extremely efficient alternative dynamo operates in low mass, fully convective stars.   \\par Beyond spectral type $\\sim$ M8, the realm of ultracool stars and brown dwarfs, a sharp drop occurs in the levels of quiescent H$\\alpha$ and X-ray emission observed, along with a complete breakdown of the activity-rotation relationship, with many fast rotators ($v$ sin $i$ $>$ 15 km s$^{-1}$) showing no evidence of quiescent H$\\alpha$ or X-ray emission \\citep{neuhauser99, gizis00, mohanty03, fleming03, west04}.  This reduction in activity may be due to the reduced ionization levels in the atmospheres of these cooler dwarfs \\citep{mohanty02}.  Magnetic fields become decoupled from the surrounding atmosphere and are no longer able to generate the necessary magnetic stress required for the dissipation of heat to sustain a chromosphere and corona.   However, this does not preclude the efficient operation of a dynamo or the generation of large-scale kG magnetic fields on ultracool dwarfs.  In fact, despite the reduction in the levels of quiescent emission, there have been detections of H$\\alpha$ and X-ray flares from a number of ultracool dwarfs \\citep{reid99, gizis00, rutledge00, liebert03, fuhrmeister04, stelzer04, rockenfeller06a} including a T dwarf \\citep{burgasser00} which point to the presence of strong magnetic fields.  \\citet{rockenfeller06b} have recently used multiband photometric monitoring to confirm a periodicity of 3.65 $\\pm 0.1$ hours in observations of the M9V dwarf 2M1707+64.  Moreover, these authors associate this periodicity with the rotational period of the dwarf and rule out the presence of dust clouds as a source of this variability. Instead, they attribute this periodicity to the presence of magnetically induced cool spots on the surface of the star, further evidence for the presence of strong surface magnetic fields on ultracool dwarfs.   \\par Based on the sharp drop in quiescent H$\\alpha$ and X-ray emission and the empirical relationships established between the observed X-ray and radio emission for a wide spectral range of low mass, main sequence stars \\citep{gudel93, benz94}, it was expected that any radio emission from ultracool stars and brown dwarfs would be below the sensitivity of the current generation of radio telescopes.  Nonetheless, \\citet{berger01} unexpectedly detected persistent, highly variable radio emission from the M9 field brown dwarf LP 944-20 at levels which, based on the lower limit for the quiescent X-ray flux established by the Chandra observation of \\citet{rutledge00}, violated the G\\\"udel-Benz relations by more than 4 orders of magnitude.  Follow up VLA observations by \\citet[hereafter B02]{berger02} yielded three further candidates. All detected sources with measured rotation rates were fast rotators, with confirmed \\textit{v} sin \\textit{i} ranging up to 60 km s$^{-1}$, despite quiescent levels of H$\\alpha$ and X-ray emission being absent, greatly suppressed, or not yet observed.  B02 postulated that the rotation-activity relationship may hold for radio emission in ultracool stars and brown dwarfs despite breaking down in other wavelength regimes.  \\citet{burgasser05} also conducted a survey of southern late M and L dwarfs using the Australia Telescope Compact Array (ATCA) and confirmed two detected sources.  One of these sources, DENIS 1048-3956, was seen to emit an extremely bright ($\\sim 20$ mJy), narrow bandwidth, $100\\%$ circularly polarized flare, with the inferred brightness temperature $\\gtrsim 10^{13} K$ requiring a coherent emission mechanism, most likely an electron cyclotron maser which would require the presence of kG magnetic fields.  \\citet[hereafter B05]{berger05} have also confirmed a periodicity of $\\sim$ 3 hours in the radio emission from the L3.5 candidate brown dwarf 2MASS J00361617+1821104 (hereafter 2M 0036+18) in observations separated by three years, citing three possible sources of the periodicity, (1) orbital motion of a companion (2) rotation of 2M 0036+18 or (3) weak periodic flaring.  Finally, a large scale survey of late M, L and T dwarfs has yielded three more confirmed radio detections and highlights the trend that $L_{\\mathrm{rad}}/L_{\\mathrm{bol}}$ increases with later spectral type, contrary to what is observed at H$\\alpha$ and X-ray wavelengths \\citep{berger06}.   \\par Observations thus far have established that persistent radio emission at the bottom of the main sequence is a reality, despite the sharp drop in activity observed in X-ray and H$\\alpha$ emission and confirming the probable existence of large-scale, stable magnetic fields on ultracool stars and brown dwarfs.  In this paper we report on VLA observations of a previously detected M9 dwarf, TVLM 513-46546 (hereafter TVLM 513), undertaken to further investigate the mechanism responsible for this anomalous radio emission.  These observations provide insights into the nature, topology and strength of the associated magnetic fields and more fundamentally, the dynamo mechanism responsible for their generation. ",
        "conclusions": "\\par We report on VLA observations showing rotational modulation of the radio emission from the M9 dwarf TVLM 513, which places strong constraints on the source size of the emitting region, and hence brightness temperature of the emission.  The resulting high brightness temperatures along with the implicit directivity of the rotationally modulated component of the radio emission are difficult to reconcile with incoherent gyrosynchrotron radiation from a population of mildly relativistic electrons.  We discuss a new model for the radio emission from ultracool dwarfs based on the assumption that the magnetic field on these objects takes the form of a large scale dipole or multipole powered with kG strength fields at the surface.  In regions over the magnetic poles, plasma can escape along open field lines forming density cavities where the electron cyclotron frequency is much greater than the plasma frequency.  If a component of the population of trapped electrons was continuously accelerated into such density cavities by magnetic field-aligned electric fields, possibly due to the presence of electrostatic double layers, such density cavities may be the source of steady-state, broadband electron cyclotron maser emission.  We also discuss the possible presence of a non-varying, unpolarized component of the radio emission, which may be due to depolarization or mode conversion of the maser emission in the density cavity. Alternatively, this component of the radio emission may be incoherent gyrosynchrotron or synchrotron radiation.  In particular, trapped particles mirroring on closed field lines in the large scale magnetic field can form high density radiation belts where electrons can be accelerated to ultrarelativistic energies emitting synchrotron radiation.  \\par Coherent maser emission can account for the presence of high levels of circular polarization and the high brightness temperatures of the radio emission from ultracool dwarfs.  The periodicity and high degree of variability from both TVLM 513 and 2M J0036+18 can be attributed to the inherent directivity of the radio emission coupled with the rapid rotation of the dwarf.   We differentiate between this broadband, persistent maser emission and the extremely bright $100\\%$ circularly polarized flare observed from the M8.5 dwarf DENIS 1048-3956 (Burgasser \\& Putman 2005), typical of flares observed from earlier type M dwarfs.   Such flares are almost always narrow-banded, $100\\%$ circularly polarized and can display structure over very short millisecond timescales.  We speculate that both the broadband, quiescent emission and narrowband, flare emission are generated in the same source regions by the same mechanism, the electron cyclotron maser, albeit with fundamental differences in the nature by which the emission is generated.  For example, the smooth broadband emission may be associated with constant particle acceleration due to quasi-stable electric fields across the source region, whereas the narrowband flaring may be due to more impulsive acceleration events.  This is analogous, once again, to Jupiter's decametric emission, which is characterized by smooth, broadband bursts with $\\Delta \\nu \\sim \\nu$ (referred to as L bursts), and more impulsive, narrowbanded bursts with $\\Delta \\nu \\ll \\nu$, which can reach much higher brightness temperatures (referred to as S bursts) \\citep{zarka98}.  \\par The configuration of the magnetosphere may also be a contributing factor to the ubiquitous rapid rotation of ultracool dwarfs.  It is unlikely that such a large-scale stable magnetic field would undergo reconnection events in the increasingly neutral atmospheric plasma of ultracool dwarfs, minimizing momentum loss due to dissipation of magnetic energy, and thus leading to much longer spin down times than observed on earlier type dwarfs. This model for the radio emission from ultracool dwarfs can also account for the fact that $L_{\\mathrm{rad}}/L_{\\mathrm{bol}}$ is maintained, and even seen to increase for late M and L type dwarfs, despite the sharp drop in magnetic activity detected at other frequencies.  The lower plasma density with later spectral type which inhibits the generation of H$\\alpha$ and X-ray emission, leads to conditions where more efficient generation of radio emission by the electron cyclotron maser instability can occur.  \\par The mechanism by which H$\\alpha$ and X-ray flares are generated on ultracool dwarfs is still unclear.  \\citet{chabrier06} postulate that the magnetic stresses necessary for such flares are generated in the higher temperature conductive layers of the star and rise to the surface, dissipating magnetic energy in the upper atmospheric layers.  Such a model is supported by the observations of \\citet{fuhrmeister04} who report on a H$\\alpha$ flare from the M8.5 dwarf DENIS 1048-3956, the source of the large coherent radio flare discussed in $\\S 1$. They interpret broadening on the blue side of the spectral line during the flare as possible evidence that the flare was associated with a rising cloud of emission ejected from the dwarf.  Such ejections may also be the source of the trapped plasma mirroring in the large-scale magnetic field responsible for the generation of the radio emission.  \\par The generation of coherent emission in density cavities associated with large scale dipolar or multipolar fields should be considered as a possible source of both quiescent and flaring radio emission in a wide range of astrophysical objects, including T Tauri stars \\citep{smith03}, RS CVn binaries \\citep{osten04}, Algol-type binaries \\citep{mutel98} and chemically peculiar stars \\citep{trigilio00}.  Moreover, the absence of high levels of circular polarization and narrowband structure does not preclude the possibility that such emission may be coherent.  \\par The recent confirmation of a strong, large-scale axisymmetric field on a rapidly rotating M4 dwarf, with a strong dipolar component \\citep{donati06}, highlights the possibility that electron cyclotron maser emission from polar density cavities may be a viable source of radio emission from dMe stars.  Indeed, \\citet{bingham01} have suggested maser emission from a shell distribution of electrons as a possible source of the 100\\% circularly polarized emission from UV Ceti.  Considering the ubiquity of highly polarized, narrowband, coherent bursts from such objects, it is quite possible that broadband, slowly varying maser emission from the same source regions responsible for the flares may also contribute to the quiescent component of the emission.  Indeed, such objects have much longer periods of rotation than later spectral type dwarfs such as TVLM 513 and 2M J0036+18 resulting in much slower rotational modulation of the highly beamed emission, and hence lower variability.  An important factor in the determination of the emission mechanism in dMe stars has been the direct measurement of the extent of stellar radio coronae, and hence constraint of the brightness temperature of the associated emission, afforded by very long baseline interferometry (VLBI).  Some VLBI observations of single dMe stars have revealed large extended radio coronae up to a few stellar radii in diameter, with brightness temperatures and degrees of circular polarization in line with incoherent emission such as gyrosynchrotron or synchrotron radiation \\citep{alef97, benz98, pestalozzi00}.  However, other observations of single dMe stars have revealed unresolved flaring and quiescent emission with brightness temperatures, $T_{\\mathrm{B}} > 10^{10}$ K, and very high degrees of circular polarization \\citep{benz91, benz95}.  In these cases a coherent emission mechanism such as electron cyclotron maser is the more plausible candidate.  Further VLBI observations may clarify the degree to which both coherent and incoherent mechanisms contribute to the radio emission from dMe stars, and whether the prevalence of an individual mechanism is frequency dependent, activity dependent or related to the magnetospheric conditions on each individual star.    \\par Further high sensitivity observations of ultracool dwarfs previously detected at radio frequencies should be undertaken to investigate the presence of periodicity and to confirm if this periodicity is indeed due to the rotation of the dwarf.  Direct detection of the magnetic fields on these objects may also be feasible.  Assuming a gyrosynchrotron source would imply a magnetic field strengths of the order of a few hundred gauss to 1kG, however, the coherent process, requires considerably higher field strengths of several kilo-gauss which could potentially be detectable using the FeH band around 1$\\mu$m \\citep{reiners06}. Although this technique is still in its infancy, the above authors show that that are sufficient FeH lines available even for the faster rotators whose observation and modelling could have an important bearing on the question of which radio process is at work in these ultracool dwarfs.  \\par Confirmation of the generation of broadband coherent emission from ultracool dwarfs would have immense implications for our understanding of stellar magnetic activity and the nature of the dynamo mechanism responsible for the generation of the magnetic fields driving this activity.  It would also highlight the possibility that a similar mechanism for the generation of coherent radio emission from large-scale dipolar or multipolar fields may be a operating in a wide range of astrophysical objects, ranging from planets to brown dwarfs and cool stars.  Comparison with more exotic phenomena such as coherent pulsar radio emission is also compelling."
    },
    "0608/astro-ph0608283_arXiv.txt": {
        "abstract": "We study the reionization histories where ionizing UV photons are emitted from decaying particles, in addition to usual contributions from stars and quasars, taking account of the fact that the universe is not fully ionized until $z \\sim 6$ as observed by Sloan Digital Sky Survey. Likelihood analysis of the three-year data from the WMAP (Wilkinson Microwave Anisotropy Probe) severely constrains the decaying particle scenario.In particular, the decaying particle with relatively short lifetime is not favored by the polarization data. ",
        "introduction": "Ionization history after recombination was not known for many years, until recent development of observations. The Sloan Digital Sky Survey (SDSS) obtained the first evidence of the reionization epoch \\cite{SDSS}. The Gunn-Peterson test for the observed spectra of quasars revealed that there existed neutral hydrogen at the redshift $z\\sim 6$, which implies that the completion of the reionization of the universe was seen at $z \\sim 6$, although the amount of the neutral hydrogen at $z\\gtrsim 6$ may not be very large to explain the observed Gunn-Peterson trough.   On the other hand, the Wilkinson Microwave Anisotropy Probe (WMAP) observation found the earlier reionization than $z\\sim 6$ in terms of the large optical depth \\cite{WMAP1,WMAP3}. The first-year data favored very large optical depth $\\tau_{op} = 0.17 \\pm 0.04$ \\cite{WMAP1}. The conventional reionization by UV photons emitted from usual early formed stars and quasars, could only reach to the optical depth of $\\sim 0.05$~\\cite{LB,stars3} if it is consistent with the SDSS data. Of course, using the nonstandard initial mass function and unrealistic escaping UV fraction from galaxies, one could achieve a large optical depth \\cite{stars1}. However, the UV fluxes from stars and quasars increase with time and full reionization takes place before $z\\sim 6$. Thus, one has to consider a more complex ionization history which may be realized, for example, by the use of different photon emission processes between Population II and III stars~\\cite{stars2}.  An alternative line is to consider another source of UV photons. It could be decaying particles \\cite{KKS,KK, decay}. In Refs.~\\cite{KKS,KK}, we considered decaying particles in addition to conventional stars and quasars and found that the decaying particle scenario is consistent with the first year WMAP data for large parameter space.  Recently, WMAP team released three year data, in which the value of the optical depth is lowered: $\\tau_{op} = 0.09 \\pm 0.03$ \\cite{WMAP3}. It changed towards the conventional value of $\\sim 0.05$, but some extra UV photons are still necessary. It would be possible to have such a small amount of extra UV photons before $z\\sim 6$ using less unrealistic assumptions for stars and quasars. In this article, however, we stick to the decaying particle scenario to assess its capability.  The structure of the paper is as follows. In the next section, we explain how we follow the evolution of the ionization history including the decaying particles. In Sect.~III, we evaluate the $\\chi^2$ for each model parameters, and construct likelihood contours for the parameters of decaying particle scenario. At the same time, the detail of the power spectra of temperature and polarization anisotropies in the cosmic microwave background (CMB) radiation are shown, and we discuss how it discriminates the preferred model parameters in the same section. The final section is devoted to our conclusions. ",
        "conclusions": "We have considered the ionizing UV photons emitted from decaying particles, in addition to the usual contributions of stars and quasars. The stars and quasars are responsible for the quick rise of the inonization fraction at $z\\sim 6$ observed by SDSS as the Gunn-Peterson trough. On the other hand, the decaying particles cause partial ionization at $z > 6$.  The decaying particle scenario could explain the large optical depth, and was consistent with the first-year WMAP data. However, the three-year WMAP data do not favour the decaying particle scenario so much. Notice that our results apply to most kinds of decaying particles which emit photons with branching ratio $B_\\gamma$ provided that one regards the effective abundance as $B_\\gamma \\Omega_\\phi$.  The ionization by the decaying partcles with $\\tau_{op} \\lesssim 0.17$ does indeed seem to be consistent within $2\\sigma$ with TT power spectrum. The peculiar aspect of the three-year WMAP data is in the polarization; in particular, the polarization data do not allow the earlier reionization which is a distinct feature of the ionization history from the UV emission from particle decays, and hence the longer lifetime is favoured. Since more power at low $\\ell$ is favoured by the polarization anisotropy, one may ambitiously claim that a better fit may be obtained for the decaying particle scenario with increasing decay rate."
    },
    "0608/gr-qc0608088_arXiv.txt": {
        "abstract": "We present high-angular resolution observations of the circumstellar disk around the massive Herbig Be star R~Mon (M$_*$$\\sim$8~M$_\\odot$) in the continuum at 2.7mm and 1.3mm and the $^{12}$CO 1$\\rightarrow$0 and 2$\\rightarrow$1 rotational lines. Based on the new 1.3mm continuum image we estimate a disk mass (gas+dust) of 0.007~M$_\\odot$ and an outer radius of $<$150~AU. Our CO images  are consistent with the existence of {\\it a Keplerian rotating gaseous disk around this star}. Up to our knowledge, this is the most clear evidence for the existence of Keplerian disks around massive stars reported thus far. The mass and physical characteristics of this disk are similar to those of the more evolved T Tauri stars and indicate a shorter timescale for the evolution and dispersal of circumstellar disks around massive stars which lose most of their mass before the star becomes visible. ",
        "introduction": "Herbig Ae/Be stars (HAEBE) are intermediate mass (M$\\sim$2--10~M$_\\odot$) pre-main sequence objects. A big theoretical and observational effort has been done in recent years for the understanding of the disk occurrence and evolution in HAEBE. For Herbig Ae and late-type Be stars (spectral type later than B7) (hereafter HAE), it is generally accepted the presence of disks similar to those associated with T Tauri stars (TTs). The association of the more massive early Be stars (HBE) with disks is more uncertain. Evidence for the existence of dusty and gaseous circumstellar disks around some HBE exists at optical, NIR and mid-IR wavelengths (Meeus et al. 2001,Vink et al. 2002, Millan-Gabet et al. 2001, Acke et al. 2005). Furthermore, disks around HBE seem to have different geometry from those around HAE and TTs. While most HAE and TTs have flared disks, HBE seem to have (if any) geometrically flat disks (see e.g. Acke et al. 2005). However, the direct detection of circumstellar disks around HBE at millimeter wavelengths has remained elusive until recently.  We carried out a high sensitivity search for circumstellar disks around HBE with the aim of studying the frequency and time-scales of disks in these massive stars. We observed a sample of HBE in the continuum at mm- and cm- wavelengths using the PdBI and VLA (Fuente et al. 2001; Fuente et al. 2003, hereafter Paper~I) and detected two dusty disks around MWC~1080 and R Mon with total (gas+dust) masses of $\\sim$0.003~M$_\\odot$ and $\\sim$0.01~M$_\\odot$ respectively. These were the first detections of dusty disks in HBE.  R~Mon is a very young HBE with a T Tauri companion separated by 0.69$''$. Because of its youth, R Mon is not directly visible in the optical but appears as a resolved conical reflection nebula in scattered light. At infrared wavelengths, R Mon appears as a point source located 0.06$\"$$\\pm$0.02 south from the optical peak. Close et al. (1997) measured an extinction of A$_v$=13.1 mag towards the star that they interpreted as due to an optically thick disk of radius, r$_{out}$=100~AU. Our previous VLA and PdBI images of R~Mon  at 1.3cm, 0.7cm and 2.7mm showed a compact source located $\\sim$0.45$\"$ south from the optical position (Paper~I). Based on these interferometric observations we derived a mass of 0.01~M$_\\odot$ for the dusty disk. R Mon hosts the most massive disk detected in HBE so far and it is the best candidate to study disks around massive stars. In this paper, we present observations at higher angular resolution and sensitivity of the R~Mon disk in the continuum at 3mm and 1.3mm and in the $^{12}$CO rotational lines. Our $^{12}$CO observations are consistent with the existence of a gaseous disk in Keplerian rotation around the star. ",
        "conclusions": "R~Mon is the first disk around a HBE detected in molecular lines and so far, our unique opportunity to investigate the physical structure and kinematics of the gas reservoir in the disks associated with these stars. Our most interesting result is that the disk is in Keplerian rotation around the star. Keplerian rotation has also been found in most of the TTs and HAE studied so far and indicates a similar formation mechanisms for the stars in the range 1--8 M$_\\odot$. The values obtained for the indexes $q$ and $p$ as well as that for the  turbulent broadening are also similar to those found in TTs and some HAE (see e.g. Dutrey et al. 2004). R~Mon is a very young object (age$\\sim$10$^5$ yr) which is still deeply embedded in the molecular cloud. On the contrary, TTs are evolved objects (age$>$10$^6$ yr) which are not any more associated with the parent molecular core (age$>$10$^6$ yr). The similarity between the disks of such different objects suggests a shorter timescale for the evolution of the disks associated with HBE. The low occurrence of disks associated with HBE also supports this conclusion. In Paper~I we showed that the disk masses in HBe stars are at least an order of magnitude lower than in HAe stars. In fact, the ratio M$_d$/M$_*$ is roughly constant and equal to 0.04 for stars with spectral type A0-M7 and M$_d$/M$_*$ $<$ 0.001 in HBe stars. This indicates a rapid dispersal of the disk material in HBE.  We have compared the R~Mon disk with those associated with massive protostars in order to have a deeper insight into the evolution of disks around HBE. Evidences for circumstellar disks have been found in massive protostars. The surface density in these disks is well described with $p>$2 (Cesaroni et al. 2005). This means that most of the gas is concentrated at small radii. Recently, Schreyer et al. (2006) detected a 1 M$_\\odot$ disk around MWC~490,  a 8--10 M$_\\odot$ star which is in a transition stage to HBE. In this case, the surface density varies with p$=$1.5 and the disk radius is 1400 AU. The values of $p$ and $r_{out}$ are similar to those derived in the disk around R~Mon. However the mass of the MWC~490 disk is more than an order of magnitude larger than that of the R~Mon disk. These results indicate that the disks around HBE flatten and lose a large fraction of their mass ($\\sim$90\\% )  before the pre-main sequence phase ($<$10$^5$ yr). The flattening of the disks can be due to the rapid grain growth. The low value of $\\beta$ measured in R~Mon is a clear evidence of the presence of large grains. The grain growth causes the optical depth of the disk to drop and allows the UV radiation to penetrate deep into the circumstellar disk and photo-evaporate the disk external layers (Dullemond \\& Dominik 2004).  Thus, we can propose an evolutionary sequence in which the disks associated with HBE start with a flaring shape but become flat during the pre-main sequence and lose most of their mass ($>$90\\%) before the star becomes visible ($<$10$^5$ yr). The results of this Letter are consistent with this scenario which, however, needs to be confirmed."
    },
    "0608/astro-ph0608006_arXiv.txt": {
        "abstract": "We present $Spitzer$ observations of thirteen $z\\sim6$ quasars using the Infrared Array Camera (IRAC) and Multiband Imaging Photometer for $Spitzer$ (MIPS). All the quasars except SDSS J000552.34$-$000655.8 (SDSS J0005$-$0006) were detected with high S/N in the four IRAC channels and the MIPS 24$\\mu$m band, while SDSS J0005$-$0006 was marginally detected in the IRAC 8.0$\\mu$m band, and not detected in the MIPS 24$\\mu$m band. We find that most of these quasars have prominent emission from hot dust as evidenced by the observed 24$\\mu$m fluxes. Their spectral energy distributions (SEDs) are similar to those of low-redshift quasars at rest-frame 0.15$-$3.5 $\\mu$m, suggesting that accretion disks and hot-dust structures for these sources already have reached maturity. However, SDSS J0005$-$0006 has an unusual SED that lies significantly below low-redshift SED templates at rest-frame 1 and 3.5 $\\mu$m, and thus shows a strong near-IR (NIR) deficit and no hot-dust emission. Type I quasars with extremely small NIR-to-optical flux ratios like SDSS J0005$-$0006 are not found in low-redshift quasar samples, indicating that SDSS J0005$-$0006 has different dust properties at high redshift. We combine the $Spitzer$ observations with X-ray, UV/optical, mm/submm and radio observations to determine bolometric luminosities for all the quasars. We find that the four quasars with central black-hole mass measurements have Eddington ratios of order unity. ",
        "introduction": "High-redshift quasars provide direct probes of the distant early universe where the first generation of galaxies and quasars formed. In the last few years, more than 20 luminous quasars at $z>5.7$ have been discovered by the Sloan Digital Sky Survey \\citep[SDSS;][]{yor00}; the most distant one is at $z=6.42$ \\citep[e.g.][]{fan00,fan01,fan03,fan04,fan06}. The discovery of these luminous objects at $z\\sim6$ reveals the existence of supermassive black holes (BHs) with masses higher than $10^9$M$_\\sun$ \\citep[e.g.][]{bar03,ves04} in the first Gyr of cosmic history. These quasars provide a unique high-redshift sample to answer a series of challenging questions: How did the first billion-solar-mass BHs appear less than 1 Gyr after the Big Bang? How did they grow with time? Did the physical structure of quasars/AGNs evolve with time? Is the emission mechanism in quasars/AGNs the same at $z\\sim6$ as at $z\\sim0$? How were quasars and starburst activity related in the earliest galaxies? What is the role the quasar activity played in early galaxy evolution?  Understanding quasars requires observations from X-ray to radio wavelengths, each spectral region revealing different aspects of quasar emission and probing different regions of the active nuclei. One of the main results from the studies of $z\\sim6$ quasars is the apparent lack of strong evolution in their rest-frame UV/optical and X-ray properties. Their emission-line strengths and UV continuum shapes are very similar to those of low-redshift quasars \\citep[e.g.][]{bar03,pen03,fan04}, the emission line ratios indicate solar or supersolar metallicity in emission-line regions as found in low-redshift quasars \\citep[e.g.][]{ham99,die03,fre03,mai03}, and the optical-to-X-ray flux ratios and X-ray continuum shapes show little evolution with redshift \\citep[e.g.][]{str05,vig05,ste06,she05,she06}. These measurements show that quasar accretion disks and emission-line regions are formed on very short time scales and their properties are not sensitive to the cosmic age.  However, it is not known whether this lack of evolution in high-energy spectral energy distributions (SEDs) at $z\\sim6$ extends to the rest-frame IR wavelength range. According to classical AGN unification models \\citep{ant93}, the accretion disk is surrounded by a dusty torus. Much of the emission from quasars/AGNs is re-processed by the dust and is re-emitted at IR wavelengths, where the quasar/AGN SEDs peak. The hottest dust lies within a few pc and produces near-IR (NIR) radiation, while warm and cool dust can extend to a few kpc and dominates the mid-IR (MIR) and far-IR (FIR) emission \\citep[e.g.][]{pol00,nen02,sie05}. The individual contributions from AGN activity and star formation to the heating of warm/cool dust are poorly known \\citep[e.g.][]{wil01}. On one hand, the shapes of the MIR-to-FIR SEDs indicate that AGN activity may dominate the heating of the warm/cool dust \\citep[e.g.][]{pol00,haa03}. On the other hand, the radio-to-FIR correlation for star-forming galaxies \\citep{con92} still holds for most IR-luminous quasars out to the highest redshifts \\citep[e.g.][]{car01,car04}, which suggests that the dust could be heated by starbursts in quasar host galaxies rather than quasar central engines. However, it is generally believed that emission from hot dust with temperature of $\\sim$1000K is directly powered by central active nuclei \\citep[e.g.][]{rie81,pol00,haa03}, and thus is closely related to quasar activity.  The $Spitzer\\ Space\\ Telescope$ \\citep[$Spitzer$;][]{wer04} allows us, for the first time, to explore the rest-frame NIR range for high-redshift quasars and to constrain the evolution of hot dust in quasar environments. \\citet{hin06} have observed thirteen $z>4.5$ quasars using the Infrared Array Camera \\citep[IRAC;][]{faz04} and Multiband Imaging Photometer for $Spitzer$ \\citep[MIPS;][]{rie04}. They find that the SEDs of these high-redshift quasars at rest wavelength 0.6--4.3 $\\mu$m do not significantly differ from those of quasars with similar luminosity at low redshifts. In this paper, we report on $Spitzer$ observations of thirteen $z\\sim6$ quasars discovered by the SDSS. All the quasars were observed with IRAC. Ten of them were also observed in the MIPS 24$\\mu$m band, while the 24$\\mu$m fluxes of the other three were taken from \\citet{hin06}.  In $\\S$2 of this paper, we present our high-redshift quasar sample and the $Spitzer$ photometry of the thirteen quasars. In $\\S$3 we show their rest-frame NIR SEDs and study hot dust at $z\\sim6$. We calculate bolometric luminosities and accretion rates for these quasars in $\\S$4 and give the summary and discussion in $\\S$5. Throughout the paper we use $\\lambda_0$ ($\\nu_0$) to denote rest-frame wavelength (frequency), and use a $\\Lambda$-dominated flat cosmology with H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m}$ = 0.3 and $\\Omega_{\\Lambda}$ = 0.7 \\citep[e.g.][]{spe03,spe06}. ",
        "conclusions": "We have carried out IRAC and MIPS 24$\\mu$m photometry for thirteen $z\\sim6$ quasars. All the quasars except SDSS J0005$-$0006 were detected with high S/N in the IRAC and MIPS 24$\\mu$m bands, while SDSS J0005$-$0006 was marginally detected in the IRAC 8.0$\\mu$m band, and not detected in the MIPS 24$\\mu$m band. The sample of the quasars is used to study the properties of IR SEDs and hot dust in high-redshift quasars. We find that the SEDs of most quasars follow low-redshift SEDs at the probed wavelengths. However, two quasars, SDSS J0005$-$0006 and SDSS J1411+1217, have unusual SEDs that lie significantly below the prediction of low-redshift SED templates at 24 $\\mu$m and/or 8 $\\mu$m, showing a strong IR deficit. A simple model shows that most of the quasars have substantial hot dust, while the two IR-weak quasars do not show any hot-dust emission. We combine the $Spitzer$ observations with X-ray, UV/optical, mm/submm and radio observations to determine bolometric luminosities for the high-redshift quasars. We find that these quasars have Eddington ratios of order unity.  It has been revealed that $z\\sim6$ quasars show a lack of evolution in their rest-frame UV/optical and X-ray SEDs. Our $Spitzer$ observations show that, for most $z\\sim6$ quasars, NIR-to-MIR SEDs also do not differ significantly from those at low redshift. This suggests that accretion disks, emission-line regions and dust structures in most high-redshift quasars have reached maturity very early on. However, we found two quasars with a strong IR deficit. As shown in Figure 4, NIR-weak Type I quasars like SDSS J1411+1217 are very rare; and Type I quasars with extremely weak NIR fluxes like SDSS J0005$-$0006 are not found in a large sample at low redshift, but exist in a small sample of thirteen quasars at $z\\sim6$, suggesting that some quasars at high redshift may have different dust properties.  In the local universe, most dust in the ISM is produced by low and intermediate-mass AGB stars, which develop 0.5$-$1 Gyr after the initial starburst. At $z\\sim 6$, the age of the universe is less than 1 Gyr, so quasar host galaxies are very young, with their first star formation likely to have occurred less than half a Gyr earlier. The two NIR-weak quasars SDSS J0005$-$0006 and SDSS 1411+1217 could be too young to have formed dust tori around them, or perhaps, dust properties in such young systems differ from those of lower-redshift quasars. Thus the unusual IR SEDs of the two quasars may be a reflection of different dust properties in very young host galaxies at high redshift. \\citet{mai04} found that the dust extinction curve in quasar SDSS J1048+4836 ($z=6.20$) is different from that observed at $z<4$ (which is SMC-like, \\citet{hop04}), implying a grain size distribution more similar to that expected from dust produced by supernovae. While other dust production mechanisms \\citep[e.g.][]{elv02} are also possible, it is currently unknown what their emission spectra are. SDSS J0005$-$0006 is especially interesting. If it has no hot dust, the observed SED at 1 $\\mu$m$<\\lambda_0<3.5$ $\\mu$m comes from the disk only. This provides a strict constraint on disk models.  Hot dust in SDSS J0005$-$0006 and SDSS J1411+1217 could also be hidden by dust tori. According to AGN unification models, the accretion disk is surrounded by a dust torus. The hottest dust, with temperature more than 1000K, lies within a few pc, while cool dust with temperature of a few tens of kelvins can extend to a few kpc. As the angle between the torus axis and the line of sight increases, more hot dust could be hidden by the cool dust torus. Thus the hot-dust radiation in the NIR would be suppressed, and the $\\lambda_0\\sim1\\mu$m dip could be shifted toward longer wavelengths \\citep{haa03}. Intriguingly, the emission line widths in the two NIR-weak quasars are the narrowest in our sample of $z\\sim 6$ quasars. This suggests that the weak NIR SEDs in two sources could be caused by an obscuration with hot dust hidden by cooler dust tori. On the other hand, both objects show normal UV continuum and UV-to-X-ray flux ratio, and are thus not type II obscured objects in the normal sense, but could be intermediate objects in which the hot dust contribution is somewhat reduced. Further deep $Spitzer$ observations are needed to constrain the SED shapes at $1\\mu m<\\lambda_0<4\\mu m$, and NIR spectra are also needed to detect possible narrow line components."
    },
    "0608/astro-ph0608230_arXiv.txt": {
        "abstract": "Astrometric measurements for 25 red dwarf systems are presented, including the first definitive trigonometric parallaxes for 20 systems within 10 pc of the Sun, the horizon of the RECONS sample.  The three nearest systems that had no previous trigonometric parallaxes (other than perhaps rough preliminary efforts) are SO 0253+1652 (3.84 $\\pm$ 0.04 pc, the 23rd nearest system), SCR 1845-6357 AB (3.85 $\\pm$ 0.02 pc, 24th), and LHS 1723 (5.32 $\\pm$ 0.04 pc, 56th).  In total, seven of the systems reported here rank among the nearest 100 stellar systems.  Supporting photometric and spectroscopic observations have been made to provide full characterization of the systems, including complete $VRIJHK_s$ photometry and spectral types.  A study of the variability of 27 targets reveals six obvious variable stars, including GJ 1207, for which we observed a flare event in the $V$ band that caused it to brighten by 1.7 mag.  Improved parallaxes for GJ 54 AB and GJ 1061, both important members of the 10 pc sample, are also reported.  Definitive parallaxes for GJ 1001 A, GJ 633, and GJ 2130 ABC, all of which have been reported to be within 10 pc, indicate that they are beyond 10 pc.  From the analysis of systems with (previously) high trigonometric parallax errors, we conclude that parallaxes with errors in excess of 10 mas are insufficiently reliable for inclusion in the RECONS sample.  The cumulative total of new additions to the 10 pc sample since 2000 is now 34 systems --- 28 by the RECONS team and six by other groups. This total represents a net increase of 16\\% in the number of stellar systems reliably known to be nearer than 10 pc. ",
        "introduction": "\\label{sec:intro}  Trigonometric parallax determinations provide one of the most fundamental measures of the cosmos.  Their beauty lies in their simplicity: they are based only on geometric techniques and are the straightforward result of the Earth orbiting the Sun.  Trigonometric parallaxes help define solar neighborhood membership, allow for accurate masses calculations in binary systems, provide distances to fundamental stellar clusters, and supply benchmarks for stellar population studies in the Milky Way and nearby galaxies.  With concerted effort, trigonometric parallaxes with accuracies of $\\sim$1 milliarcsecond are possible, yielding distances good to 10\\% even at 100 pc, and allow astronomers to create an accurate three-dimensional map of the nearby Galaxy.  Much like early explorers' discoveries created a more complete map of our Earth, accurate trigonometric parallaxes create a more detailed picture of the solar neighborhood, filling in the blank regions of nearby space.  \\citet{1995yCat.1174....0V} compiled decades of ground-based work by dedicated astrometrists in {\\it The General Catalogue of Trigonometric Stellar Parallaxes} (also known as the ``Yale Parallax Catalog'', hereafter, YPC).  The YPC includes 15994 total trigonometric parallaxes for 8112 stars measured using photographic plates and CCDs published before the end of 1995.  The Hipparcos space mission, launched in August 1989, provided $\\sim$118,000 parallaxes \\citep[ hereafter HIP]{1997A&A...323L..49P,1997yCat.1239....0E}, and was essentially complete for stars with $V = 7.3-9.0$, while reaching to observe a few dozen stars even fainter than $V$$\\sim$13.  However, Hipparcos could not measure most nearby white, red, and brown dwarfs, which comprise more than 80\\% of the solar neighborhood population, and many of these systems have not yet had their distances measured via ground-based trigonometric parallax efforts.  Here we present results for nearby star systems observed during our Cerro Tololo Inter-american Observatory Parallax Investigation (CTIOPI).  Twenty systems are new members of the Research Consortium on Nearby Stars (RECONS) sample, meeting the RECONS requirements that they have trigonometric parallaxes, $\\pi_{trig}$, larger than 100 mas with errors less than 10 mas (see section~\\ref{sec:astr.reconsdefn}). We provide a brief dossier for each new systems, including $VRIJHK_s$ photometry and spectral types.  We also provide improved $\\pi_{trig}$ for two additional important systems previously known to be within 10 pc, GJ 54 AB and GJ 1061, and for three systems that have reported $\\pi_{trig}$ larger than 100 mas, but which, in truth, lie beyond the 10 pc horizon.  In total, we report 27 new trigonometric parallax measurements for 25 systems. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{The Solar Neighborhood Population} \\label{sec:disc.population}  After the publication of the YPC in 1995 and Hipparcos results in 1997, there were 215 systems in the RECONS sample.  Between the time of those publications and the first parallax results from CTIOPI (2005), only eight new systems had reliable trigonometric parallaxes measured that placed them within 10 pc (see Table 1), and the nearest, GJ 1061, was by our group \\citep{1997AJ....114..388H}.  CTIOPI has already reported eight additional new 10 pc systems meeting the RECONS criteria --- five in \\citet{2005AJ....129.1954J}, two in \\citet{2005AJ....130..337C}, and one in Costa et al.~(2006, AJ in press).  In this paper, we add an additional 20 systems to the sample of stars within 10 pc with high quality parallaxes.  Counting additions in this paper, since the publication of YPC and HIP results, the total number of new RECONS systems stands at 36.  The 34 systems added since 2000 constitute a 16\\% increase in the number of systems within 10 pc in six years.  Figure~\\ref{fig:henry2} shows an HR diagram of the RECONS sample, plotting $M_V$ vs. $V-K$ for all systems with parallaxes meeting the RECONS criteria (see Section~\\ref{sec:astr.reconsdefn}).  All multiples have been deconvolved into individual components as well as can currently be done at both the $V$ and $K$ bandpasses.  Small circles indicate objects known to be within 10 pc after publication of the YPC and HIP catalogs, as well as new discoveries by other groups. Solid points represent RECONS discoveries.  Labels are given for prominent or unusual stars.  Triangles mark the four known subdwarfs within 10 pc.  AU Mic and AT Mic AB (GJ 803 and GJ 799 AB, respectively) form a wide triple system that is apparently quite young.  The system appears to be comoving with the $\\beta$ Pic group and has an age of only $\\sim$20 Myr.  The position of all three stars above the main sequence supports the hypothesis that the system is young, and the stars have not yet descended onto the main sequence. Further support came recently, when AU Mic was found to be the nearest (9.9 pc) star with a dust disk directly observed at optical and near-infrared wavelengths \\citep{2004Sci...303.1990K}.  The three known L dwarfs within 10 pc --- 2MA 0036+1821, 2MA 1507-1627, and DEN 0255-4700 --- are shown, but none of the eight known T dwarfs or four extrasolar planets are plotted.  All of the systems reported in this paper are mid- to late-type M dwarfs (spectral types listed in Table 2).  The three brightest new stars are labeled --- LP 771-095 A, LHS 3746, and GJ 1207 --- as well as many of the fainter discoveries from the RECONS group and others. Note in particular the $M_V$ of 24.4 for DEN 0255-4700, for which a parallax and photometry were determined during our CTIOPI 1.5m program (Costa et al.~2006, AJ in press).  We believe that this is the faintest object outside the Solar System for which $M_V$ has been determined.  The dominance of the nearby stellar population by red dwarfs is obvious from Figure~\\ref{fig:henry2}.  For objects in the refereed literature, including this paper, the complete count of current objects meeting the RECONS criteria is: 18 white dwarfs, 4 A types, 6 F, 21 G, 44 K, 240 M, 3 L, 8 T, and 4 extrasolar planet candidates. Counts include the Sun and the adoption of a dividing line between spectral types K and M at $M_V$ $=$ 9.00.  Thus, of the 348 total objects, the 240 M dwarfs comprise 69\\% of the nearby population. Elimination of the presumably non-stellar L, T, and extrasolar planet portions brings the fraction to 72\\% of all stars, and this fraction will continue to grow as more nearby red dwarfs are revealed.  \\subsection{Hidden Multiples} \\label{sec:disc.multiples}  To confirm or reveal unresolved multiple systems, we make a comparison between the accurate trigonometric distances presented here and in previous CTIOPI papers and the photometric distances estimated via the suite of $M_K$-color relations described in \\citet{2004AJ....128.2460H}.  This technique provides distance estimates good to 15\\%, measured by running RECONS stars with accurate trigonometric parallaxes back through the suite of relations, under the assumption that the stars are single.  Photometric distance estimates for the systems discussed in detail in this paper are given in column 14 of Table 2, and the number of relations used for each target (maximum of 12) is given in column 15.  Figure~\\ref{fig:henry3} compares the photometric distances with trigonometric distances for objects within 20 pc, as reported by CTIOPI in the four papers with parallaxes to date.  The distance cutoff of 20 pc has been selected because that is the distance limit of our 20-20-20 sample (within 20 pc, orbital periods less than 20 years, at least one component with a mass less than 20\\% of the Sun's; for details see \\citet{1999ApJ...512..864H}), which includes red dwarf binaries observed to measure high quality masses.  Presumed single main sequence stars are shown as open circles and known close multiples with combined photometry are shown with filled squares.  The solid line represents identical photometric and trigonometric distances.  Dotted lines trace distances differing by 20\\% --- stars above the outlined region are often subdwarfs; stars below are often unresolved multiples because they are predicted photometrically to be closer than they actually are.  In two cases, SCR 1845-6357 AB and LP 771-095 BC, the multiples are within the 20\\% region, as is expected because of the relatively large magnitude differences between the components.  Additional close multiples in which secondaries contribute minimal light to the system would not be revealed with this method.  From left to right (increasing trigonometric distance), the known close multiples recovered below the 20\\% region are G 041-014 ABC, GJ 54 AB, GJ 2005 ABC, G 089-032 AB, SCR 0630-7643 AB, and GJ 2130 BC.  The three best candidates for previously unnoticed multiples are shown as heavy open squares --- from left to right, GJ 555, ER 2, and LTT 6933.  GJ 555 is a marginal case, while the latter two are almost certainly close multiples with small flux ratios.  \\subsection{The Big Picture} \\label{sec:disc.future}  To date, we have observed more than 400 stars for parallax during CTIOPI.  This paper is the fifth in our series that provides accurate parallaxes for nearby stellar systems --- one from Ianna's plate parallax effort at Siding Spring Observatory and four from the ongoing CCD effort at CTIO.  With the 20 new systems described here, we have now reported high quality parallaxes for a total of 29 new RECONS systems.  We anticipate that future papers will yield additional nearby systems because the pool of stars nearer than 10 pc without accurate parallaxes is by no means drained.  We emphasize that the stellar companions discussed here are not at CTIOPI's detection threshold --- the ASPENS program will search for substellar and possibly planetary mass companions.  Red and white dwarf systems within 10 pc and south of $\\delta$ $=$ 0, including all but one (WT 460) of the new RECONS members discussed here, are being observed intensely to reveal any possible long term astrometric perturbations.  Initial results imply that we can detect companions with masses as low as 10 M$_{Jup}$ for stars with good observational coverage and strong reference star sets.  Each new system found within 10 pc is a member of the fundamental sample of objects that is used to: (1) determine accurate luminosity and mass functions in the solar neighborhood, (2) evaluate the multiplicity of stellar populations, and (3) calculate the total amount of mass found in stars and substellar objects.  Because of their proximity, each system is a promising target for detailed astrophysical studies such as astroseismology, variability investigations, direct radii measurements, and metal abundance evaluations.  Aside from astrophysical efforts, these systems have great astrobiological potential because of their proximity.  They are prime targets to be scrutinized as possible planetary hosts, and any planets found can be explored as nearby harbors of life."
    },
    "0608/astro-ph0608140_arXiv.txt": {
        "abstract": "In this paper we describe the performance of the 9 instrumented IceCube strings and 16 surface IceTop stations taking data at the Geographical South Pole after 2 deployment seasons.  We will focus on the description of the array and on the construction and data analysis status. The expected full array performance is discussed and compared to other results for the relevant physics studies using high energy neutrinos for astrophysical neutrino searches. ",
        "introduction": "\\label{sec1}  A considerable effort is ongoing to broaden our knowledge of the universe through the observation of high energy neutrinos. It is reasonable to believe, given also the increasing hints coming from gamma astronomy experiments, that the most powerful accelerators of the universe are sites of proton acceleration and interactions on targets of matter and photons. Part of the energy could be transfered to neutrinos from meson decays in astrophysical sources of high energy cosmic rays, such as Active Galactic Nuclei (AGNs) and GRBs (Gamma-Ray Bursters), shock waves in SuperNova Remnants or micro-quasar binary systems exhibiting radio jets. The operation of previous detectors and the underlying connections between TeV-$\\gamma$ production at sources indicate that $\\nu$ fluxes are small and the detection of a positive signal may need years of operation of cubic-km scale detectors. Promising $\\nu$ sources are those where strong $\\gamma$ absorption is likely in the source, such as the micro-quasar systems LS 5029 and LSI +61 303 (which is below IceCube horizon), from which periodical TeV photon signals have been observed by HESS (Aharonian et al., 2006a) and MAGIC (Albert et al., 2006), respectively. Whilst the $\\gamma$-ray emissions from these sources can be weakened by electromagnetic cascading processes, neutrinos would travel undisturbed to us producing up to a few tens of events in a km$^3$ detector per year (Aharonian et al., 2006b).  The state of the art of the detection of diffuse $\\nu_{\\mu}$ fluxes is summarized in Fig.~\\ref{fig1} (on the left). The Waxman \\& Bahcall (WB) upper limit derived from measured $10^{18}$ eV cosmic rays (Waxman \\& Bahcall , 2000) represents the maximum $\\nu$ emission from optically thin extra-galactic sources. AMANDA-II recently presented the most stringent limit (G. Hill et al., 2006) and the Baikal experiment in a siberian lake is competitive in setting limits for diffuse fluxes using cascade-like events (Aynutdinov et al., 2006). IceCube can observe fluxes well below this WB upper limit already in 1\\,yr of active operation (Ahrens et al., 2004). The most significant results for neutrino point-like source searches are currently produced by the AMANDA-II detector: an average sensitivity for $E^{-2}$ neutrino fluxes of $E^2 dN/dE \\sim  6 \\cdot 10^{-8}$ GeV cm$^{-2}$ s$^{-1}$ has been achieved  in 1001 days of operation during 2000-4. In a sample of 4282 $\\nu$ events (about 4/day) no evidence for a significant excess of $\\nu$s has been found. The largest fluctuation of the atmospheric $\\nu$ background has a significance of 3.74$\\sigma$, that occurs 69\\% of the times on 100 randomized data samples. Upper limits for a few sources of interest are given in Tab.~\\ref{tab1}. A search for the years 2000-4 is going to be published soon (Ackermann et al., 2006a). IceCube capabilities are shown in Fig.~\\ref{fig1} (on the right): the atmospheric $\\nu$ event rate will be of the order of $10^5$/yr. The angular resolution (median of the angle between the $\\nu$ and the reconstructed $\\mu$ track)  will improve compared to the $2-3^{\\circ}$ of AMANDA-II better than $\\sim 1^{\\circ}$. Notice that the results shown in Fig.~\\ref{fig1} for IceCube from (Ahrens et al., 2004) are obtained using AMANDA reconstruction adapted to the different geometry. We are currently working on a reconstruction based on the full knowledge of the pulse waveforms (see Sec.~\\ref{sec2}) from which we expect a considerable improvement. Moreover, it should be noted that the sensitivities for IceCube shown in Fig.~\\ref{fig1} include the contribution of a high prediction for the high energy charm contribution. As discussed in (Ahrens et al., 2004) lower predictions will result in a less conservative estimate.  \\begin{table}[t!] \\begin{center} \\caption{\\label{tab1}Upper limits to $E^{-2}$ neutrino fluxes for $E_{\\nu} > 10$~GeV and measured and expected background events for a few of the monitored neutrino source candidates in 1001 days of AMANDA-II live-time in 2000-4.} \\bigskip \\begin{tabular}{llll} \\hline Source Name & N. obs. & N. exp. & Upper limit (90\\% c.l.) \\\\ &              &                & $10^{-9}$ cm$^{-2}$ s$^{-1}$\\\\ \\hline Crab   & 10   & 6.74 & 10.2\\\\ Mrk 501   &  8              & 6.39 & 8.5\\\\ LSI +61 303   &  5             & 4.81 & 7.4\\\\ 1ES 1959+650   &  5 &4.77              & 7.8\\\\ \\hline \\end{tabular} \\end{center} \\end{table} \\begin{figure}[t!] \\vspace{10truecm} \\special{psfile=diffuse_numu_models.eps  voffset=-2 hoffset=-3 hscale=34.0 vscale=35 angle=0} \\special{psfile=nulimits.eps voffset=-2 hoffset=190 hscale=34.0 vscale=44 angle=0} \\caption[h]{{\\bf On the left:}  The red full circles is the atmospheric $\\nu$ flux measurement using 2000 AMANDA-II data (Ackermann et al., 2005) with the central values for different  calculations shown as solid red lines (for a review see Honda and Gaisser, 2005). The AMANDA-II upper limits for E$^{-2}$ $\\nu_{\\mu}$ diffuse fluxes (thick solid horizontal line) and for other models: limits = solid curved lines (G. Hill et al., 2006), models = dashed curved lines from AGN jet model in  (Mannheim et al, 2001) and AGN core model in (Stecker, 2005). The Baikal limit for all flavor $\\nu$s (upper line)  is shown  (Aynutdinov et al., 2006), divided by 3 to account for equi-partition of neutrino flavors due to oscillations. From bottom to top, the dotted horizontal lines are: sensitivity for 3~yrs of IceCube operation (Ahrens et al. 2004), Waxman \\&Bahcall upper limit (Waxman \\& Bahcall, 2000)  and 3~yrs ANTARES sensitivity (Montaruli et al., 2005). {\\bf On the right:} upper limits on $E^{-2}$ $\\nu_{\\mu}$ fluxes (90\\% c.l.) vs declination in the Southern hemisphere from 1001 d of 2000-4 AMANDA-II data for selected sources (full red circles) and sensitivity (dashed line between circles) (Ackermann et al., 2006a). The IceCube sensitivity (lower dotted line) is from (Ahrens et al., 2004) for 1 yr. The MACRO upper limits (open squares) are from (Ambrosio et al., 2001). The ANTARES preliminary sensitivity for 1 full year of data is shown in the Northern hemisphere (Montaruli et al., 2005).} \\label{fig1} \\end{figure}  Neutrino telescopes are aimed at the detection of high energy $\\nu$s ($> 100$ GeV) through the Cherenkov light produced by $\\nu$ induced charged products. Neutrinos of all flavors could reach us from sources, after long baseline oscillations transform the original $\\nu_{\\mu}$ and $\\nu_{e}$ fluxes from meson decays in sources. Muon induced $\\nu$s point back to their sources for energies $> 1$\\,TeV, since the kinematical muon-neutrino angle becomes small, and the long muon tracks allow precise reconstruction through the arrival times of detected Cherenkov photons. This, together with the advantage of the target mass increase due to the muon range, makes 3 dimensional arrays of PMTs located in the depths of the ice or in deep sea or lake water, telescopes capable at pointing back to sources with resolutions at degree level or better. Cascades induced by neutral current neutrino interactions or by tau and electron neutrino charged current interactions, though with much reduced pointing capabilities, allow for searches of diffuse fluxes thanks to the harder spectral properties of neutrinos from shock acceleration mechanisms compared to neutrinos produced in atmospheric cascades. These atmospheric neutrinos represent from one side a background for astrophysical searches, from the other a calibration tool and a signal for high energy interaction studies. Similarly the atmospheric muons are a dangerous background, as well a tool for understanding the detector with huge statistics. Their discrimination requires good reconstruction capabilities to single out upward-going events that could be only induced by $\\nu$s crossing the Earth, or potential to identify events with the interaction vertex in or close to the interaction region. ",
        "conclusions": "After 2 deployment seasons, IceCube is currently taking data in the configuration of 9 strings and 16 surface stations. Based on last season experience, the Collaboration expects to deploy up to 14 strings per season and complete IceCube in the 2010-11 season. Software development and physics analyses are ongoing and we expect first results very soon."
    },
    "0608/astro-ph0608189_arXiv.txt": {
        "abstract": "Color--magnitude diagrams (CMDs) are presented for the {\\it Spitzer} SAGE (Surveying the Agents of a Galaxy's Evolution) survey of the Large Magellanic Cloud (LMC). IRAC and MIPS 24 \\mic \\ epoch one data are presented. These data represent the deepest, widest mid--infrared CMDs of their kind ever produced in the LMC. Combined with the 2MASS survey, the diagrams are used to delineate the evolved stellar populations in the Large Magellanic Cloud as well as Galactic foreground and extragalactic background populations.  Some 32000 evolved stars brighter than the tip of the red giant branch are identified. Of these, approximately 17500 are classified as oxygen--rich, 7000 carbon--rich, and another 1200 as ``extreme'' asymptotic giant branch (AGB) stars. Brighter members of the latter group have been called ``obscured'' AGB stars in the literature owing to their dusty circumstellar envelopes. A large number (1200) of luminous oxygen--rich AGB stars/M supergiants are also identified. Finally, there is strong evidence from the 24 \\mic \\ MIPS channel that previously unexplored, lower luminosity oxygen--rich AGB stars contribute significantly to the mass loss budget of the LMC (1200 such sources are identified). ",
        "introduction": "The Spitzer Space Telescope (Spitzer), with its rapid wide field mapping capability and sensitivity, allows us to survey the thermal emission produced by mass loss from individual evolved stars in the nearest galaxies. \\citet{mex06} presented the {\\it Spitzer} SAGE (Surveying the Agents of a Galaxy's Evolution) survey of the Large Magellanic Cloud (LMC) along with a detailed description of its goals, expected extent in depth and coverage and some preliminary results. The SAGE survey is designed to enable studies of the life--cycle of baryonic matter, as traced by dust emission, in the LMC. In this work, we focus on one corner stone of this life--cycle, the dusty evolved star population, which is returning matter to the interstellar medium. Coupled with detailed star formation history (SFH) studies at optical wavelengths, SAGE provides the opportunity to link specific LMC populations to their dust content as a function of key parameters such as metallicity and age. Such analyses can, in turn, be used to validate predictive models of infrared emission in high redshift galaxies. In the present paper, we focus on the necessary first step of identification and characterization of the dusty evolved star populations.  SAGE provides a global view of the dust producing stellar populations in the LMC that will contribute to our understanding of evolved stars. Essentially all significant, dusty, mass-losing stars in the LMC will be detected by SAGE as it provides the deepest, widest survey in the near-- and mid-infrared wavelengths (3--24 microns) produced to date for this nearby galaxy. These wavelengths are particularly well suited to the investigation of the late stages of stellar evolution because once beyond the near--infrared (five microns), emission from circumstellar dust can become the dominant source of emission over the radiation from the photosphere. By eight microns, many LMC sources show appreciable emission from dust, and the MIPS 24 micron band is extremely sensitive to cool envelopes with excess dust emission even for sources with little evidence of strong dust excess at shorter wavelengths.  SAGE builds on and leverages a wealth of past observations of evolved stars in the LMC. Optical \\citep[MCPS; see][and http://ngala.as.arizona.edu/dennis/mcsurvey.html]{zar04} and near--infrared surveys, for example, DENIS \\citep{denis} and 2MASS \\citep{skrut06}, have provided a global view of the LMC at shorter wavelengths. See \\citet{cioni06} and \\citet{niko00} for LMC data from these two surveys, respectively. The Infrared Astronomy Satellite \\citep[IRAS,][]{neug84} revolutionized the study of evolved stars by observing objects with luminous dusty envelopes at long wavelengths. This survey has provided samples of stars in the LMC to be observed in greater detail up to the present, including spectroscopic observations with {\\it Spitzer}; see \\citet{zijl06} and \\citet{mark05}. However, IRAS was only sensitive enough to observe the brightest sources in the LMC \\citep{schwer89} from which several hundreds of mass--losing evolved star candidates could be identified \\citep{loup97}. The successful Midcourse Space Experiment \\citep[MSX,][]{price01} was about four times more sensitive and provided a full LMC survey \\citep{egan01}. SAGE should be roughly 1000 times more sensitive than MSX when complete \\citep[][the present data are roughly 400 times more sensitive]{mex06}. While near--infrared and optical surveys of the LMC have identified lower mass evolved stars, only SAGE has the depth to detect such stars at wavelengths necessary to explicitly explore their dust properties. The ISOCAM Magellanic Cloud Mini--Survey \\citep[see][]{cioni03} has provided a preliminary mid--infrared view of the LMC over a small field of view (0.3 square degrees) near the depth of SAGE.  The SAGE sensitivity and areal coverage of the entire LMC will allow a detailed quantitative derivation of the global mass loss budget from all stellar populations when combined with existing and future mid--infrared spectroscopic observations of the asymptotic giant branch (AGB) stars and supergiants; for example, see \\citet{vloon99}, \\citet{vloon05}, \\citet{mark05}, and \\citet{zijl06}. Several studies have derived detailed SFHs in the LMC from a variety of field locations and clusters \\citep{holtz99,olsen99} using principally optical color--magnitude diagrams (CMDs) and sophisticated stellar models and fitting techniques \\citep[e.g.][]{harris01}. Recently, \\citet{cioni06} made a global SFH calculation using the 2MASS and DENIS surveys. Their new models still have difficulties producing the correct mix of observed carbon--rich AGB stars (C--rich or C--stars) versus oxygen--rich AGB stars (O--rich or M stars), but their extension to these wavelengths of the general technique is a great advance. With SAGE, we will be able to associate the mass--loss and chemical properties of the AGB stars in the LMC with their known evolutionary state from existing SFH studies. This coupling can then be used to investigate, with precision, the evolutionary status of galaxies which lie at much larger distances through the development and testing of detailed population synthesis models \\citep[for example, see][]{mouhlan03}.  In this paper, we continue the analysis begun by \\citet{mex06} by presenting CMDs of the survey data for the first of two epochs. We lay the ground work for the detailed analysis to come when the full SAGE survey is available by dissecting the observed CMDs and indicating where various LMC and fore and background populations reside in these diagrams.  The present paper covers approximately 49 square degrees in the IRAC bands and MIPS 24 \\mic \\ band. ",
        "conclusions": "The Hess diagrams and CMDs presented in Figures~\\ref{j3}--\\ref{824} can be used to assess the relative importance to the mass loss budget in the LMC. Detailed model calculations will be made once the full SAGE data set is in hand {\\it and} there are appropriate template spectra to assess the stellar envelope dust properties as a function of location in these diagrams. The SAGE team and others have existing and planned {\\it Spitzer} and ground based spectroscopic follow--up programs underway. Some of the brightest mass--losing sources in the LMC have already been analyzed \\citep[see, for example,][]{zijl06}. These sources are prolific, but rare mass--losing objects. The full impact of the less luminous stars on the AGB has yet to be quantified.  The effect of dust in shaping the infrared CMDs is clear at a glance from Figures~\\ref{j3}--\\ref{824} and considering the 2MASS $H-K$ $vs.$ $K$ CMD \\citep{niko00}. In the latter, even at $Ks$ (2.1 \\mic), the tail of extreme AGB stars decreases in brightness for redder colors. This suggests the primary effect at short wavelengths is circumstellar extinction. In the $J-$[3.6] $vs.$ [3.6] diagram of Figure~\\ref{j3}, the extreme star branch has essentially constant brightness. By 8 \\mic, the sequence is increasing in brightness clearly indicating that the circumstellar envelopes are exhibiting successively more excess emission. The [8.0]$-$[24] $vs.$ [24] CMD shows that the majority of the reddest objects are the fainter objects we have identified with background galaxies. However, at the brightest and reddest this sequence seems to merge with the luminous stellar sources (i.e. the ``extreme'' AGB and SG stars). Indeed a survey of the SIMBAD catalog shows a number of these objects to be PNe and other bright, stellar, IRAS sources (one Wolf--Rayet star and several objects associated with HII regions, as well as obscured AGB stars). Of the 236 sources with [8.0]$-$[24] $>$ 3.0 mag and [24] $<$ 7.0 mag, 10 are known PNe and 13 are classified as IRAS or IR sources \\citep[of which ][discuss four as obscured AGB stars]{loup97}, The majority have no other counterparts in the SIMBAD data base.  This general picture is confirmed by the {\\it Spitzer} IRS sources plotted in Figure~\\ref{824}. The stars identified as ``C'' lie among the ``extreme'' stars where many C--rich objects are expected\\footnote{One object, LI--LM~603, had anomalous IRAC/MIPS colors. In this case, we substituted the IRS spectrophotometry.}. The ``S'', or silicate, sources lie among these, but also slightly to the red and at brighter magnitudes where the M supergiants ({\\it green} points) and luminous AGB stars predominate. The IRS PNe lie at bright magnitudes and red colors as do several massive YSOs as discussed in \\S~3.2.2. Clearly SAGE will provide a host of new mass losing objects to follow up in detail.  While there are differences in detail, the location of the basic LMC AGB and supergiant sequences correspond remarkably well to those predicted from Galactic models (moved to the distance of the LMC, distance modulus$=$18.5) for the analogous objects \\citep{wain92,coh93,coh94,coh95}. In particular, this includes the general location, extent, and slope of the ``extreme'' AGB stars in Figures~\\ref{j3} and \\ref{j8}, as well as the position of the most luminous supergiants and O--rich objects above the extreme stars. The predicted positions of the Galactic objects such as PNe and HII regions are also in rough agreement with the objects identified above in the [8.0]$-$[24] $vs. $[24] CMD (Figure~\\ref{824}).  Figure~\\ref{j3} shows the most dominant {\\it Spitzer} population by number is the red giants. The RGB is the most densely populated sequence with a tip magnitude of TRGB([3.6])$=$ 11.85. There are approximately 650000 stars on the RGB below the TRGB in the epoch one data set. The peak number density (2400 stars in a 0.0475$\\times$0.0625 square magnitude pixel) occurs at $J-$[3.6], [3.6] $=$ 0.8, 15.6.  The RGB number density drops rapidly to a value $\\sim$ 25$\\%$ of its value just to the faint side of the tip. Approximately 74$\\%$ of the stars in the $J-$[3.6] $vs.$ [3.6] CMD above the TRGB are O--rich, C--rich, and SG types. Roughly 12$\\%$ of stars above the TRGB are foreground giants and dwarfs with the remainder being mostly blue LMC supergiants according to the discussion in \\S3.  The AGB stars and supergiants can be divided up among those with strong mass loss (the so called ``extreme'' stars) and those whose colors suggest weaker mass loss or less dusty envelopes. Of the latter, most of the objects can be classified as O--rich or C--rich based on the $J-K$ $vs.$ $K$ classification of \\citet{cioni06}. The red tail of C--stars (literally where Figure~\\ref{j3} shows red points) stretches to approximately $J-$[3.6] $=$ 3.1 where the density of objects drops significantly toward redder colors (see Figure~\\ref{j3}). This color limit is a useful indicator of the extreme or obscured AGB stars as it corresponds roughly to the point in the $J-K$ $vs.$ $K$ CMD where the AGB stars become fainter at redder color (exhibiting strong circumstellar extinction). The C/O star ratio (not including the ``extreme'' stars) is 0.39 (see Table~1). This value is somewhat lower than given by \\citet{cioni06}; however, the data analized by \\citet{cioni06} includes stars at over a larger area than considered here and for which the number of C--stars is significant. A full analysis of the C/O star ratio with position will follow in a subsequent paper.   According to the models of \\citet{mar03}, the position of the C--stars in the red tail \\citep[discuss this in terms of $J-K$]{mar03} is due to cooler temperatures resulting from their molecular opacity differences compared to the O--rich stars. The opacity difference is a direct result of the changing molecular equilibrium as dredge up changes the ratio of C/O (and hence the important molecular constituents in the C--rich phase compared to the O--rich phase). The colors are thus ``photospheric'' in nature, and not until the obvious break in the CMDs (e.g., $J-$[3.6] $=$ 3.1) does dust significantly affect the colors. The majority of stars in the red tail of Figure~\\ref{j3} are optically visible C--stars with objective prism identifications given by \\citet{bm90} and \\citet{kdm01}. About 25$\\%$ of the sources identified as C--stars here have no previous identification in the SIMBAD database (but are 2MASS sources).  Comparison of the [3.6]$-$[8.0] $vs.$ [8.0] and [8.0]$-$[24] $vs.$ [24] CMDs suggests the relative importance of mass loss can be seen among the different AGB stars. The {\\it green} SGs lie to the blue of the {\\it yellow} extreme stars in the former diagram, but generally to the red in the latter. The SGs thus appear to have cooler, perhaps more extended, dust envelopes. The {\\it red} ``visible'' C--stars generally have little or no 24 \\mic \\ excess, whereas a subset of the {\\it blue} O--rich AGB stars exhibit increasing amounts of excess emission and presumably mass loss. This sequence is approximately at constant [8.0] magnitude ($\\approx$ 10.5). This group of AGB stars is well above the TRGB in the shorter wavelength CMDs and represents the brightest tip of stars classified as O--rich just at the transition to the ``visible'' C--star locus in the $J-$[8.0] $vs.$ [8.0] CMD (Figure~\\ref{j8}). The group of stars is clearly visible as a slight enhancement in density in the Hess diagram (Figure~\\ref{j8}) above and to the red of the thin finger of AGB stars which rise above the TRGB (i.e. at $J-$[8.0], [8.0] $\\approx$ 1.5, 10.5).  Fainter O--rich AGB stars are not detected at high signal--to--noise in this preliminary data set at 24 \\mic. Even so, this is striking evidence for mass loss in lower luminosity stars than has generally been explored in the LMC.  Van Loon et al. (1999) derived mass loss rates from a sample of 57 stars in the LMC chosen by their infrared excess. Their analysis suggests lower luminosity AGB stars have mass loss rates \\apge 10$^{-7}$ M$_\\odot$ yr$^{-1}$. The SAGE 24 \\mic \\ data reach somewhat below the luminosities probed by van Loon et al., and the epoch one data will be combined with epoch two photometry, allowing SAGE to reach even fainter sources at 24 \\mic \\ (0.4 mag). The fact that the low luminosity mass losing sources we have identified above (Figure~\\ref{824}, ``F'') span a range of 24 \\mic \\ excess from no excess to $\\sim$ 2 mag suggests they represent the faintest population of dusty sources with significant mass loss (10$^{-7}$ \\apge M$_\\odot$ yr$^{-1}$ \\apge 10$^{-8}$) in the LMC. Thus, the full SAGE catalog should detect the all the important mass losing sources in the LMC. Adding up the first epoch SAGE sources identified here as ``extreme'' AGB stars, SG/luminous AGB, and the present lower luminosity O--rich AGB stars, we find approximately 10$\\%$ of the evolved stars above the TRGB are likely significant mass--loss sources.  In Figure~\\ref{sed}, the average spectral energy distribution (SED) is plotted for these low luminosity AGB sources. The sample of stars is arbitrarily divided into three bins of varying [8.0]$-$[24] color ([8.0]$-$[24] $\\le$ 0.67, 0.67 $\\le$ [8.0]$-$[24] $<$ 1.34, 1.34 $\\le$ [8.0]$-$[24] $<$ 2.0) to show the increasing excess emission along the sequence at constant 8 \\mic \\ brightness. There are 632, 439, and 75 stars in these bins, respectively. To contribute to a given bin, a star is required to have valid photometry for each 2MASS, IRAC, and MIPS (24 \\mic) band; hence, there are more stars plotted in Figure~\\ref{824} than represented in Figure~\\ref{sed} since the former includes objects for which there is no required $H$, [4.5], or [5.8] photometry.  The sources in all bins have very similar SEDs up to 8 \\mic, but then exhibit increasing 24 \\mic \\ flux indicative of successively more cool dust emission which we suggest is due to increasing mass loss among otherwise very similar objects.  Careful scrutiny of Figure~\\ref{sed} shows that the stars with the least amount of 24 \\mic \\ excess are the brightest at shorter wavelengths, while those with the most excess are fainter at the shorter wavelengths. The cool dust envelopes are likely providing additional extinction at the shorter wavelengths. While not shown in the figure, arbitrarily scaling the SEDs to match in the near infrared shows that they are nearly identical until about 6.8 \\mic. At this point, a clear but small excess begins to show at 8 \\mic. This fact, and the fact that no fainter population of obvious mass losing stars is evident in the much deeper CMD shown in Figure~\\ref{j8} provides support for our claim that SAGE will detect all the important mass losing sources in the LMC.  Could these objects be chance alignments of ``normal'' AGB stars and background (dusty) galaxies at 24 \\mic? To estimate the possible number of chance alignments, consider the SWIRE number counts given above. The IRAC--24 \\mic \\ search radius was 1$''$. The total area covered by {\\it all} 24 \\mic \\ sources is then 28000 sources $\\times$ the area per 24 \\mic \\ source $=$ 0.007 square degrees. Multiplying this by the number of SWIRE galaxies per square degree (253) gives two chance alignments per 28000 24 \\mic \\ sources)."
    },
    "0608/astro-ph0608695_arXiv.txt": {
        "abstract": "Neutrino-neutrino interactions can lead to collective flavour conversion effects in supernovae and in the early universe.  We demonstrate that the case of ``bipolar'' oscillations, where a dense gas of neutrinos and antineutrinos in equal numbers completely converts from one flavour to another even if the mixing angle is small, is equivalent to a pendulum in flavour space.  Bipolar flavour conversion corresponds to the swinging of the pendulum, which begins in an unstable upright position (the initial flavour), and passes through momentarily the vertically downward position (the other flavour) in the course of its motion.  The time scale to complete one cycle of oscillation depends logarithmically on the vacuum mixing angle.  Likewise, the presence of an ordinary medium can be shown analytically to contribute to a logarithmic increase in the bipolar conversion period.  We further find that a more complex (and realistic) system of unequal numbers of neutrinos and antineutrinos is analogous to a spinning top subject to a torque.  This analogy easily explains how such a system can oscillate in both the bipolar and the synchronised mode, depending on the neutrino density and the size of the neutrino-antineutrino asymmetry. Our simple model applies strictly only to isotropic neutrino gasses. In more general cases, and especially for neutrinos streaming from a supernova core, different modes couple to each other with unequal strength, an effect that can lead to kinematical decoherence in flavour space rather than collective oscillations. The exact circumstances under which collective oscillations occur in non-isotropic media remain to be understood. ",
        "introduction": "\\label{sec:introduction}  The ``refractive index'' caused by a medium has a strong impact on neutrino flavour oscillations~\\cite{Wolfenstein:1977ue, Mikheev:1986gs, Mikheev:1986wj, Fuller1987, Notzold:1987ik, Kuo:1989qe}. This matter effect is a standard ingredient for neutrino oscillations in laboratory experiments and in astrophysical environments.  However, when neutrinos themselves form a significant ``background medium'' as in the early universe or in core-collapse supernovae, the oscillation equations become nonlinear, sometimes resulting in surprising collective phenomena. Based on Pantaleone's key observation that neutrinos as a background medium produce a flavour off-diagonal refractive index~\\cite{Pantaleone:1992eq}, the behaviour of dense neutrino gases was investigated in a series of papers by Samuel, Kosteleck\\'y and Pantaleone~\\cite{Samuel:1993uw, Kostelecky:1994ys, Kostelecky:1993yt, Kostelecky:1993dm, Kostelecky:1995dt, Kostelecky:1995xc, Samuel:1996ri, Kostelecky:1996bs, Pantaleone:1998xi}. Two classes of collective effects were identified in these papers: we shall call them ``synchronised'' and ``bipolar'' oscillations, respectively.  Synchronised oscillations occur when the neutrino-neutrino interaction potential is large compared to the ordinary oscillation frequencies in vacuum or in a medium, and a sufficiently large asymmetry exists between the neutrino and antineutrino distributions.  As a result, all neutrinos and antineutrinos oscillate with the same frequency that is a certain average of the ordinary oscillation frequencies. In the spin-precession analogy of flavour oscillations, the flavour polarisation vectors of all neutrino modes form one big spin that precesses in a weak ``external magnetic field.'' This big spin is held together by the strong ``internal magnetic field'' formed in flavour space by the strong neutrino-neutrino interaction~\\cite{Pastor:2001iu}. Various qualitative and quantitative aspects of synchronised oscillations and applications to neutrino flavour oscillations in the early universe~\\cite{Lunardini:2000fy, Dolgov:2002ab, Wong:2002fa, Abazajian:2002qx} were studied a few years ago, motivated by the question if neutrinos with chemical potentials reach flavour equilibrium in the early universe before the epoch of big-bang nucleosynthesis.  Much less attention has been paid to bipolar flavour conversions, but very recently it has been recognised that this peculiar phenomenon likely plays a crucial role for supernova neutrino oscillations~\\cite{Duan:2005cp, Duan:2006an, Duan:2006jv}. In the simplest case, bipolar oscillations occur in a dense gas of equal numbers of neutrinos and antineutrinos of the same flavour. For a suitable mass hierarchy, even a small mixing angle can cause a complete conversion of both neutrinos and antineutrinos to the other flavour. We stress that a non-vanishing vacuum mixing angle is pivotal; bipolar oscillations, although a quasi self-induced effect, will not occur if the vacuum mixing angle strictly vanishes.  The relevant conditions for bipolar conversions probably occur for neutrinos streaming off a collapsed supernova core where one expects a hierarchy of average energies $\\langle E_{\\nu_e}\\rangle<\\langle E_{\\bar\\nu_e}\\rangle <\\langle E_{\\nu_x}\\rangle$ with $\\nu_x=\\nu_\\mu,\\bar\\nu_\\mu,\\nu_\\tau,\\bar\\nu_\\tau$~\\cite{Keil:2002in}. Assuming equal luminosities for all species, the number flux of, say, $\\nu_\\mu$ and $\\bar\\nu_\\mu$ is each smaller than that of $\\nu_e$ and $\\bar\\nu_e$ respectively. For the inverted mass hierarchy one would then expect bipolar oscillations driven by the ``atmospheric'' neutrino mass difference and the small mixing angle $\\Theta_{13}$.\\footnote{We concentrate here on 13-mixing because $\\nu_\\mu$ and $\\nu_\\tau$ streaming off a supernova core have equal spectra, behave equally in ordinary matter, and are known to be nearly maximally mixed so that the 23-mixing is not relevant in this context. The ``solar'' mass difference is much smaller and the corresponding hierarchy is normal so that {12}-mixing is expected to lead to less prominent effects at larger distances from the supernova core. In general, however, one would need to perform a three-flavour treatment.}  Bipolar flavour conversion occurs when the neutrino-neutrino interaction energy, $\\mu=\\sqrt{2}G_{\\rm F}n_\\nu$, exceeds a typical vacuum oscillation frequency $\\omega=\\Delta m^2/2E$.  Furthermore, the neutrino-antineutrino asymmetry must not be too large in a sense to be quantified later.  Once these conditions are met, bipolar oscillations take place for astonishingly small values of the mixing angle and are nearly unaffected by the presence of an ordinary background medium, even if it is much denser than the neutrino gas~\\cite{Duan:2005cp}. These oscillations are also unrelated to a Mikheyev-Smirnov-Wolfenstein (MSW) resonance of the ordinary background medium.\\footnote{Of course, it has long been recognised that the high density of neutrinos near a supernova neutrino sphere could lead to refractive effects comparable to those of the ordinary medium, and nonlinear effects could play an important role~\\cite{Pantaleone:1994ns, Qian:1994wh, Sigl:1994hc}.  The conclusion of these early works appears to be that the neutrino-neutrino term causes a small shift of the oscillation parameters where an MSW resonance occurs and hence a small correction to the ordinary matter effect. Subsequent numerical studies of nonlinear effects in the supernova hot bubble region actually revealed flavour conversion for surprisingly small neutrino mass differences, but a connection to the bipolar oscillation mode was not made~\\cite{Pastor:2002we, Balantekin:2004ug}.}  The notion of ``bipolar oscillations'' originates from the claimed numerical observation that the flavour polarisation vectors of all neutrinos and of all antineutrinos align together to form two ``block spins'' that evolve separately.  Indeed, the analytic descriptions of Refs.~\\cite{Kostelecky:1995dt, Samuel:1996ri} were based on the study of a system of one neutrino and one antineutrino polarisation vector that were taken to represent, respectively, the complete ensemble of neutrinos and antineutrinos.  We find that this description of bipolar oscillations is incorrect and actually has never been explicitly demonstrated in the literature.  In fact, each mode of the neutrino and antineutrino ensemble evolves differently and the term ``bipolar'' is a misnomer. On the other hand, the behaviour {\\it is\\/} bipolar in the sense that neutrinos and antineutrinos oscillate in ``opposite directions'' and thus form two separate cohorts, even if they do not form two block spins. Therefore, we use the term ``bipolar'' to describe this collective phenomenon.  The newly recognised role of bipolar oscillations represents a change of paradigm for supernova neutrino oscillations.  Inspired by these exciting developments we turn to an analytic study of bipolar oscillations with the aim of providing a simple qualitative and quantitative understanding of the salient features of this puzzling effect. Analytic solutions for certain cases have already been provided in the literature~\\cite{Kostelecky:1995dt, Samuel:1996ri}. However, a judicious choice of variables enables us to write the equations of motion in the form of an ordinary pendulum. This picture allows one to grasp the salient features of the bipolar phenomenon at a single glance.  Moreover, it allows one to calculate explicitly the dependence of the bipolar oscillation period on the vacuum mixing angle and on the density of an ordinary background medium.  When the fluxes of neutrinos and antineutrinos are not equal (as expected for supernova neutrinos), the equations for the flavour pendulum remain the same, but the asymmetric initial conditions imply the presence of an inner angular momentum (i.e., spin) of the pendulum: the system is equivalent to a spinning top subject to a torque. If the top is not spinning, we simply recover the motion of an ordinary spherical pendulum (i.e., bipolar behaviour of a symmetric $\\nu$-$\\bar\\nu$ system).  Otherwise the motion is more complicated. If the spin is sufficiently large, the top precesses in the force field exerting the torque (i.e., synchronised oscillations). If the spin is too small, the top wobbles or even completely turns over (i.e., bipolar behaviour of an asymmetric system).  In the most general case, however, different neutrino modes have different energies and, in a non-isotropic medium such as that encountered by neutrinos streaming off a supernova core, couple to each other with different strengths (``multi-angle case''). In this situation kinematical decoherence rather than collective oscillations can obtain and is an unavoidable outcome certainly in the simplest system of equal numbers of neutrinos and antineutrinos. However, in the more realistic case of unequal neutrino and antineutrino fluxes in conjunction with a slowly varying effective neutrino density, collective oscillations often still obtain. The exact criteria that determine if kinematical decoherence or collective oscillations occur remain to be understood.  We begin in Sec.~\\ref{sec:oscillator} with the simplest bipolar system consisting of one polarisation vector for the neutrinos and one for the antineutrinos and establish the equivalence to a flavour pendulum. We then study analytically the impact of an ordinary background medium in Sec.~\\ref{sec:matter}. In Sec.~\\ref{sec:density} we allow the neutrino density to vary and show that this effect is crucial for the nearly complete flavour conversion in supernovae. In Sec.~\\ref{sec:asymmetric} we consider a system with an initial $\\nu$-$\\bar\\nu$ asymmetry and show its equivalence to a spinning top subject to a torque. We apply our insights to explain the salient features of flavour conversion of the neutrinos streaming off a supernova core. In Sec.~\\ref{sec:manymodes} we consider a system of many modes and discuss the conditions under which it is equivalent to the simple bipolar system.  In Sec.~\\ref{sec:quantum} we discuss the possibility of flavour conversion initiated by quantum fluctuations, in the absence of flavour mixing.  A summary of our findings is given in Sec.~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have studied bipolar neutrino oscillations, i.e., the flavour evolution of an ensemble initially consisting of equal numbers of $\\nu_e$ and $\\bar\\nu_e$. We have shown that the classical equations of motion can be cast in the form of a pendulum in flavour space. The surprising bipolar conversion effect observed for the inverted mass hierarchy corresponds to the pendulum starting in a near upright position, with the excursion angle growing exponentially until the pendulum makes an almost complete swing.  Conversely, if it starts in a nearly vertical downward position (i.e., normal mass hierarchy), the system behaves as a harmonic oscillator.  For the inverted case, the time for a complete swing is given by the pendulum's oscillation period times a factor depending on the logarithm of the initial excursion angle which is nearly identical with twice the vacuum mixing angle. Therefore, the bipolar conversion is delayed by the logarithm of the small vacuum mixing angle. Likewise, we derived an analytic solution for the case when ordinary matter is present and showed that it affects the bipolar conversion time also only logarithmically.  If the vacuum oscillation frequencies are different for different modes, one cannot represent the ensemble by two ``block spins.'' However, in a dense $\\nu$-$\\bar\\nu$ gas our model remains unaffected except that the vacuum oscillation frequency is replaced by an average over all modes. When the coupling strength between different modes varies, as would be realistically expected in a non-isotropic ensemble, yet other forms of behaviour appear. In particular, the different modes can kinematically decohere in flavour~space.  Indeed, our results suggest---and it has been numerically observed in simulations~\\cite{Duan:2006an,Duan:2006jv}---that partial flavour decoherence, instead of a simple swapping of flavours, is a possible feature in a multi-coupling/multi-angle system.  However, the same simulations also suggest that the flavour evolution of neutrinos streaming off a supernova core is qualitatively approximated by a single-angle treatment, and that collective behaviour appears to be the more generic outcome.  A partial explanation for this, at least in the inverted hierarchy case, is provided by our observation that the single-angle system with a decreasing neutrino density and a non-zero neutrino-antineutrino asymmetry never becomes properly bipolar. Rather, it evolves such that the potential and kinetic energy of the pendulum remain equipartitioned, signifying that the system remains at the edge of the bipolar condition throughout its evolution and is hence better immuned to decoherence.  The only apparent case of practical interest for this discussion is flavour conversion of neutrinos streaming off a supernova core where collective flavour transformations play an important role. Close to the neutrino sphere, the oscillations are synchronised up to a few tens of kilometres, then they enter the bipolar regime, and finally, beyond 100--200 km, ordinary oscillations occur~\\cite{Duan:2005cp, Duan:2006an, Duan:2006jv}. Ordinary matter effects modify the oscillations in the usual way both in the synchronised regime and far away where collective effects are irrelevant, whereas in the intermediate regime of bipolar oscillations ordinary matter has no significant impact. This counter-intuitive situation was conjectured and numerically observed in Refs.~\\cite{Duan:2005cp, Duan:2006an, Duan:2006jv}. In our model of a flavour pendulum the impact of ordinary matter can be calculated analytically.  Note that while the bipolar behaviour extends over a large range in radius outside the neutrino sphere, we have explicitly assumed in our treatment that the system consists of only two flavours because only one mass splitting is of importance. However, it could well be that the solar mass difference $\\Delta m_{\\rm solar}^2 \\sim 8 \\times 10^{-5} \\, {\\rm eV}^2$ cannot be ignored in the whole region. If this is the case a full three-flavour description must be employed, and new phenomena might arise.  In any case, collective neutrino oscillations, unsuppressed by ordinary matter, in the region a few tens of kilometres above the neutrino sphere will likely change the picture of supernova flavour oscillations and observable consequences in various ways.  Taking the atmospheric mass hierarchy to be inverted, the ``single-angle approximation'' predicts a swapping of the $\\bar\\nu_e$ and $\\bar\\nu_\\mu$ as well as of the $\\nu_e$ and $\\nu_\\mu$ fluxes with a possible impact on r-process nucleosynthesis~\\cite{Pantaleone:1994ns, Qian:1994wh, Sigl:1994hc, Qian:1993dg}, energy transfer to the stalling shock wave~\\cite{Fuller1992}, and the possibility to observe shock-wave propagation effects in neutrinos~\\cite{Schirato:2002tg, Takahashi:2002yj, Lunardini:2003eh, Fogli:2003dw, Tomas:2004gr, Dasgupta:2005wn, Fogli:2004ff, Choubey:2006aq, Fogli:2006xy}.  Nothing new happens in the case of a normal mass hierarchy, so that one still expects observable effects such as Earth matter effects in the neutrino signal from the next galactic supernova~\\cite{Dighe:1999bi, Lunardini:2001pb, Dighe:2003be, Chiu:2006, Mirizzi:2006xx}. However, this conclusion assumes the validity of the single-angle treatment. Partial or complete kinematical decoherence, caused by the multi-modal nature of the system, will affect the flavour composition of the neutrinos passing the bipolar region even in the normal hierarchy.  The one case that probably remains unaffected is the prompt deleptonisation burst where initially the $\\nu_e$ flux is strongly enhanced relative to $\\nu_{\\mu}$, $\\bar\\nu_{\\mu}$, $\\nu_{\\tau}$ and $\\bar\\nu_{\\tau}$, while the $\\bar\\nu_e$ flux is strongly suppressed~\\cite{Kachelriess:2004ds}.  In this case, the bipolar condition is not fulfilled and one expects ``ordinary'' synchronised oscillations."
    },
    "0608/astro-ph0608376_arXiv.txt": {
        "abstract": "We study the kinematics of gaseous disks in triaxial dark matter halos using the closed-loop orbit solutions in non-axisymmetric potentials. The orbits are in general non-circular and, for given triaxiality, their ellipticity depends on the ratio of escape to circular velocities, $V_{\\rm esc}^2/V_c^2$. This ratio increases steeply towards the center for cold dark matter (CDM) halo density profiles, implying that even minor deviations from spherical symmetry may induce large deviations from circular orbits in the velocity field of a gaseous disk, especially near the center. This result suggests that caution should be exercised when interpreting constraints on the presence of density cusps in the dark halo derived from the innermost velocity profile. Simulated long-slit rotation curves vary greatly in shape, depending primarily on the viewing angle of the disk and on its orientation relative to the principal axes of the potential. ``Solid-body'' rotation curves---typically interpreted as a signature of a constant density core in the dark matter distribution---are often obtained when the slit samples velocities near the major axis of the closed loop orbits. Triaxial potentials imprint specific symmetries in 2D velocity fields, generally inducing ``twists'' in the isovelocity contours and anti-symmetric patterns in opposite quadrants. We suggest that triaxial halos may be responsible for the variety of shapes of long-slit rotation curves of low surface brightness (LSB) galaxies, as well as for the complex central kinematics of LSBs, which are sometimes ascribed to the presence of ``radial motions'' in the gas. We argue that LSB rotation curves might be reconciled with the structure of CDM halos once the effects of halo triaxiality on the dynamics of gaseous disks are properly taken into account. ",
        "introduction": "\\label{sec:intro}  It is commonly believed that the inner regions of low surface brightness (LSB) galaxies are ideal probes of the inner structure of dark matter halos.  Although estimates of the stellar mass-to-light ratio remain somewhat uncertain, the baryonic component contribution to the mass budget in these galaxies is generally thought to be small.  Under this assumption, dynamical tracers of the potential such as rotation curves are expected to cleanly trace the dark matter distribution in LSBs. This provides important astrophysical clues to the nature of dark matter, since the spatial distribution of dark material in these highly non-linear regions is expected to be quite sensitive to the physical properties of the dark matter.  LSB rotation curves may thus be contrasted directly with theoretical predictions of the inner structure of halos, and there is now an extensive body of work in the literature that reports substantial disagreement between the shape of LSB rotation curves and the circular velocity curves of simulated cold dark matter (CDM) halos \\citep[see, e.g.,][]{FLORES94,MOORE94,MCGAUGH98,DEBLOK01}.  Some of these rotation curves are fit better by circular velocity curves arising from density profiles with a constant density ``core'' rather than by the ``cuspy'' density profiles commonly used to fit the structure of CDM halos \\citep[hereafter NFW]{NFW96, NFW97}. This discrepancy adds to a growing list of concerns regarding the consistency of CDM with observational constraints on the scale of individual galaxies \\citep[see, e.g.,][]{SELLWOOD00} that has prompted calls for a radical revision of the CDM paradigm on small scales \\citep[see, e.g.,][]{SS00}.  Before accepting the need for radical modifications to CDM it is important to note a number of caveats that apply to the LSB rotation curve problem.  For instance, many of the early rotation curves where the disagreement was noted were significantly affected by beam smearing in the HI data \\citep{SWATERS00}. The observational situation has now improved substantially thanks to higher-resolution rotation curves obtained from long-slit and two-dimensional H$\\alpha$ observations \\citep[see, e.g.,][]{DEBLOK02,SWATERS03,SIMON03,SIMON05}. We shall restrict our analysis to these newer datasets in what follows.  We also note that, strictly speaking, the observational disagreement is with the fitting formulae used to parameterize the structure of simulated CDM halos (usually the profile proposed by NFW), rather than with the structure of simulated halos themselves. Although the fitting formulae provide a simple and reasonably accurate description of the mass profile of CDM halos, the radial range over which they have been validated often does not coincide with the scales where the disagreement has been identified. This is important, since small but significant deviations between the NFW profile and simulated halos have been reported as the mass and spatial resolution of the simulations has increased \\citep{MOORE98,GHIGNA00,FUK97,FUK01}.  In particular, the central slope of the density profile has been the subject of much debate.  The NFW density profile is given by $\\rho(r) \\propto (r/r_s)^{-1} (1+r/r_s)^{-2}$, where $r_s$ is a scale radius. This profile approaches a power-law $\\rho(r) \\propto r^{-1}$ ``cusp'' near the center. \\cite{MOORE98} and \\cite{GHIGNA00} argued that halos simulated with higher resolution have steeper central density profiles, $\\rho(r) \\propto r^{-1.5}$.  However, \\cite{HAYASHI04} and \\cite{NAVARRO04} show that the density profile of halos becomes shallower towards the center and in fact may not actually converge to a well-defined asymptotic inner power law.  At present, the simulation results of all groups are in good agreement, and the logarithmic slope of the density profile is $\\simeq -1.2$ or shallower at the innermost resolved radius \\citep{DIEMAND04}.  Although there is no broad consensus yet regarding how deviations in simulated halo profiles from fitting formulae like NFW may affect the comparison with observed rotation curves \\citep[see, e.g.,][]{HAYASHI04}, the fact that the deviations worsen towards the centre, in addition to the observed scatter in halo profiles, advise against using extrapolations of simple fitting formulae such as the NFW profile to assess consistency with observation.  Finally, it must be emphasized that the ``cusp vs. core'' problem arises when comparing rotation speeds of gas in LSB disks to spherically-averaged circular velocities of dark matter halos.  Given that CDM halos are expected to be significantly non-spherical, some differences between the two are to be expected.  Previous studies have shown that simulated halos are triaxial objects with shapes ranging from oblate to prolate \\citep{DAVIS85,BARNES87,FRENK88,WARREN92,JING95,THOMAS98,JING02}.  The angular momentum of the halo tends to be aligned with the minor axis, therefore one expects galactic disks to lie in a plane close to the plane defined by the intermediate and major axes of the halo.  It is therefore important to take into account deviations from spherical symmetry in the structure of CDM halos in order to make predictions regarding the rotation curves of gaseous disks that may be compared directly to observation.  In fact, asymmetries in the velocity fields of gas disks have been observed \\citep[e.g.,]{COURTEAU03}, although in normal galaxies they are usually attributed to non-circular motion induced by bars or spiral arms.  We address this issue here by exploring the closed loop orbits within simple non-axisymmetric potentials whose mass profiles are consistent with cold dark matter halos.  We follow the formalism developed by \\cite{GERHARD86} and \\cite{BT87} in the context of triaxial bulges and bar potentials, and by \\cite{FRANX94} and \\cite{SCHOEN97} to probe non-axisymmetry in spiral galaxies. We focus our analysis on the shape and inner slope of the rotation curves inferred for disks in triaxial halos by simulating long-slit observations of their velocity fields.  We find that velocities along the slit may deviate significantly from the circular velocity, a result that may account for the variety of LSB rotation curve shapes and that may reconcile the ``cores'' inferred for some LSBs with the cuspy mass profile of CDM halos. ",
        "conclusions": "\\label{sec:disc}  Given the variety of rotation curve shapes that arise when considering gaseous disks in non-axisymmetric halo potentials, the outlook for reconciling dark matter cusps with LSB rotation curves is promising. In particular, the fact that departures from circularity are expected to increase towards the center, and that relatively large effects are possible even in the presence of minor asphericity appear rather encouraging.  However, it would be premature to argue that the problem has been fully solved. Although from the analysis presented here it appears that, on general grounds, one may expect at least some rotation curves to be consistent with cored profiles, it is unclear at this point whether the frequency (and sizes) of apparent ``cores'' is consistent with this interpretation. This is likely to depend critically on the triaxiality of the potential in the very inner regions, as well as on the response of the halo to the assembly of the disk and its orientation relative to the halo's principal axes. All of these issues are still quite uncertain.  It is therefore important to build a more compelling case for the triaxial interpretation of cores in LSB rotation curves, so as to render it falsifiable. Are the shapes of CDM halos in fact consistent with the perturbed potentials we have explored?  We are currently investigating this using a sample of high-resolution galaxy-sized halos.  Are there any corroborating traits that may be used to confirm or exclude the hypothesis that halos surrounding LSBs are indeed triaxial?  Baryonic matter in the form of a bar could have a similar effect on gas motions as a triaxial dark matter halo \\citep{RHEE04}.  How can one best verify the triaxial halo interpretation in two-dimensional velocity maps?  Identifying a clean and unambiguous indication of triaxiality, such as the unusual isovelocity ``twists'' and symmetry patterns shown in Figure~\\ref{figs:rotcur}, will be as important as the success of aspherical halos in reproducing the rich variety of shapes of LSB rotation curves.  Only if this is accomplished shall we be able to conclude that LSB rotation curves do not preclude the presence of dark matter density cusps, thereby freeing the CDM paradigm of one vexing challenge on small scales."
    },
    "0608/astro-ph0608050_arXiv.txt": {
        "abstract": "We calculate coherent neutrino and antineutrino flavor transformation in the supernova environment, for the first time including self-consistent coupling of intersecting neutrino/antineutrino trajectories. For neutrino mass-squared difference $|\\delta m^2|=3\\times 10^{-3}\\,\\mathrm{eV}^2$ we find that in the normal (inverted) mass hierarchy the more tangentially-propagating (radially-propagating) neutrinos and antineutrinos can initiate collective, simultaneous medium-enhanced flavor conversion of these particles across broad ranges of energy and propagation direction. Accompanying alterations in neutrino/antineutrino energy spectra and fluxes could affect supernova nucleosynthesis and the expected neutrino signal. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608099_arXiv.txt": {
        "abstract": "We show that the star known as 2MASS J05162881+2607387 (hereafter J0516) is a double-lined eclipsing binary with nearly identical low-mass components.  The spectroscopic elements derived from 18 spectra obtained with the High Resolution Spectrograph on the Hobby-Eberly Telescope during the Fall of 2005 are $K_1=88.45\\pm 0.48$ km s$^{-1}$ and $K_2=90.43\\pm 0.60$ km s$^{-1}$, resulting in a mass ratio of $q=K_1/K_2= 0.978 \\pm 0.018$ and minimum masses of $M_1\\sin^{3}i=0.775\\pm 0.016\\,M_{\\odot}$ and $M_2\\sin^{3}i=0.759\\pm 0.012\\,M_{\\odot}$, respectively.  We have extensive differential photometry of J0516 obtained over several nights between 2004 January-March (epoch 1) and 2004 October-2005 January plus 2006 January (epoch 2) using the 1m telescope at the Mount Laguna Observatory.  The source was roughly 0.1 mag brighter in all three bandpasses during epoch 1 when compared to epoch 2.  Also, phased light curves from epoch 1 show considerable out-of-eclipse variability, presumably due to bright spots on one or both stars.  In contrast, the phased light curves from epoch 2 show little out-of-eclipse variability.  The light curves from epoch 2 and the radial velocity curves were analyzed using our ELC code with updated model atmospheres for low-mass stars.  We find the following: $M_1=0.787 \\pm 0.012\\,M_{\\odot}$, $R_1=0.788 \\pm 0.015\\,R_{\\odot}$, $M_2=0.770 \\pm 0.009\\,M_{\\odot}$, and $R_2=0.817 \\pm 0.010\\,R_{\\odot}$.  The stars in J0516 have radii that are significantly larger than model predictions for their masses, similar to what is seen in a handful of other well-studied low-mass double-lined eclipsing binaries.  We compiled all recent mass and radius determinations from low-mass binaries and determine an empirical mass-radius relation of the form $ R(R_{\\odot})=0.0324 + 0.9343M(M_{\\odot}) +0.0374M^2(M_{\\odot}) $. ",
        "introduction": "\\begin{table*}[t] \\caption[Mass and Radius of other double-lined binaries] {Mass and Radius of sample double-lined eclipsing binaries plotted in Figure \\ref{mr}. \\label{tab_mr}} \\begin{tabular}{lccl}    \\hline \\hline Name &  Mass (\\Msun) & Radius (\\Rsun) & Reference \\\\  \\hline V818 Tau B & $0.7605 \\pm 0.062$ & $0.768 \\pm 0.010$ & \\citet{tor02} \\\\ RX J0239.1-1028 A & $0.736 \\pm 0.009$ & $0.735\\pm 0.018$ &\\citet{rib05}\\\\ RX J0239.1-1028 B &  $0.695 \\pm  0.006$ &   $0.710 \\pm 0.016$&\\citet{rib05}\\\\ GU Boo A & $0.610 \\pm  0.007$ &  $0.623 \\pm 0.016$ &\\citet{lop04}\\\\ GU Boo B & $0.599 \\pm  0.006$ &  $0.620 \\pm   0.020$& \\citet{lop04} \\\\ YY Gem AB & $0.5992 \\pm 0.0047$ &  $0.6191 \\pm  0.0057$ &\\citet{tor02} \\\\ NSVS0103 A & $0.540 \\pm 0.002$ & $0.527 \\pm 0.002$ & Lopez-Morales et al. (in prep)\\\\ NSVS0103 B & $0.498 \\pm 0.002$ & $0.505 \\pm 0.002$ & Lopez-Morales et al. (in prep)\\\\ TrES-Her0-07621 A & 0.493$\\pm$0.003 & 0.453$\\pm$0.060 & Creevey et al. (2005)\\\\ TrES-Her0-07621 B & 0.489$\\pm$0.003 & 0.452$\\pm$0.050 & Creevey et al. (2005)\\\\ BW3 V38 A & 0.44$\\pm$0.07 & 0.51$\\pm$0.04 & \\citet{mac04}\\\\ BW3 V38 B & 0.41$\\pm$0.09 & 0.44$\\pm$0.06 & \\citet{mac04}\\\\ CU Cnc A & $0.4333 \\pm 0.0017$ &  $0.4317 \\pm  0.0052$&\\citet{rib03}\\\\ CU Cnc B & $0.3890 \\pm  0.0014$ &  $0.3908 \\pm  0.0094$&\\citet{rib03}\\\\ CM Dra A & 0.2307$\\pm$0.0010 & 0.2516$\\pm$0.0020 & Lacy (1977) and Metcalfe et al. (1996) \\\\ CM Dra B & 0.2136$\\pm$0.0010 & 0.2347$\\pm$0.0019 & Lacy (1977) and Metcalfe et al. (1996) \\\\ \\hline \\end{tabular} \\end{table*}  Binary stars offer the best opportunity for accurate measurements of the radii and masses of stars, measurements essential to verify stellar evolution theory and to determine the properties of other diverse objects such as white dwarfs, neutron stars, black holes, and extra-solar planets. These measurements rely on dynamical constraints from observed radial and rotational velocities, and geometric constraints from photometric time series observations. Using computer models of binary stars we can derive physical parameters (e.g. masses, radii, etc.). Understanding the structure and evolution of stars is a basic goal of stellar astronomy, and is  required in most other branches of astronomy. Critical tests of evolution theory for stars other than the Sun can be made on a small set of eclipsing binary stars [see, Pols et al.\\ 1997; Schr\\\"oder et al.\\ 1997]. In general, when accurate tests are available, the results of stellar evolution models compare favorably to data for main sequence stars with masses greater than one solar mass \\citep{pol97}. In contrast, evidence has been growing that the models for stars on the lower main sequence have problems when confronted with precise data from eclipsing binaries. For example, \\citet{tor02} showed that all available evolutionary models underestimate the radii of the components of the M-star YY Gem by about 20$\\%$ and overestimate the effective temperatures by 150 K or more. Similar discrepencies are found in V818 Tau \\citep{tor02}, in CU Cnc \\citep{rib03}, and GU Boo \\citep{lop04}.  See Table \\ref{tab_mr} for a compilation of recent mass and radius determinations derived from low-mass eclipsing binaries.    The disagreement between the models and the data for these  binaries is troubling since models for low mass stars are used to estimate the ages for open clusters and individual T Tauri stars by placing them in an HR diagram.  Since the number of well-studied low-mass binaries is still relatively small, observations of additional low mass binaries would be extremely useful. The star known as 2MASS J05162881+2607387 (hereafter J0516) was discovered to be an eclipsing binary by Schuh et al.\\ (2003), who noticed its variability during the course of an extensive campaign to monitor a pulsating white dwarf.  The eclipse period was determined to be 1.29395(25) days.  On the basis of a moderate resolution spectrum and photometric colors, Schuh et al.\\ (2003)  determined an effective temperature of 4200 K for the primary, corresponding to a spectral type of K7V.  Assuming that the orbital period is the same as the eclipse period, Schuh et al.\\ (2003) modelled the light curve and derived a mass ratio of $M_2/M_1 \\approx 0.11$, a radius ratio of $R_2/R_1 \\approx 1$, and a temperature ratio of $T_2/T_1 \\approx 0.6$.  Based on their modelling results, Schuh et al.\\ (2003) suggested that J0516 consists of a late K-type star paired with a brown dwarf.  Schuh et al.\\ (2003) also briefly considered the possibility that the orbital period of J0156 was twice the eclipse period (i.e.\\ $2\\times 1.29395 = 2.58790$ days). In that case, the primary and secondary eclipses are identical in their phased light curve.  Our spectroscopic results clearly indicate that the orbital period of J0516 is in fact 2.58791 days.  We discuss below our  observations of J0516, models of the light and velocity curves, and implications of our results.   \\begin{figure*}[t] \\center \\includegraphics[angle=0,scale=.55]{f1.eps} \\caption{ A $10\\times 10$ arcminute $I$-band image of the field of J0516 showing the seven comparison stars (Table 2). The position of J0516 is indicated by the hash marks. \\label{comp}} \\end{figure*} ",
        "conclusions": "We have derived fairly accurate values of the masses and radii of the two stars in J0516. The masses are determined to 1.5\\% precision and the radii are formally determined to 1.9\\% precision, although there may be a systematic error of a few percent on the radius determinations. This high level of precision compares favorably to the precision achieved on the small number of other low-mass eclipsing binaries mentioned in Section 1. For low-mass stars with well determined masses and radii, the observed radii are consistently larger than what evolutionary models predict for their masses. The stars in J0516 are no exception. Figure \\ref{mr} shows the stars in J0516 and those of other well-studied low-mass binaries in a mass-radius plot.  We also show in Figure \\ref{mr}  theoretical mass-radius relationships with ages of 50 Myr, 100 Myr, and 300 Myr taken from the models of Baraffe et al.\\ (2002). No single model passes through all of the observations, and all of the stars are above the 300 Myr model. We formed an empirical mass-radius relationship by fitting a parabola using weights on both the mass and radius through the other stars and the Sun (excluding J0516) and found the following equation $$ R(R_{\\odot})=0.0324 + 0.9343M(M_{\\odot}) +0.0374M^2(M_{\\odot}). $$ Compared to this empirical relationship, the J0516 primary agrees quite well, whereas the J0516 secondary is slightly too large.   Why is the secondary of J0516 so large for is mass? Ribas (2005) raises the possibility that the high level of stellar activity observed in the well-studied double-lined low-mass eclipsing binaries is what causes these stars to have radii that are larger than what is predicted.  As we discussed earlier, J0516 is much more active than the other binaries, and the secondary's deviation from the mass-radius relationship is the most extreme. It is not clear, however,  if higher levels of activity would automatically lead to larger  radii.  The Li I line is often used as an age indicator. An upper limit to the equivalent width of the Li I line at 6708~\\AA\\ of 20 m\\AA\\ translates roughly to an abundance of $A({\\rm Li})=0.0$, which translates into a lower limit on the age of the binary of about 150 Myr (Steinhauer 2003; Deliyannis  private communication). The evolutionary models of Baraffe et al.\\ (2002) indicate that a $0.8\\,M_{\\odot}$ star takes about 70 Myr to contract and reach its ``normal'' main sequence radius.  Taken at face value, the lower limit on the age of 150 Myr from the lack of lithium indicates that J0516 is well past its pre-main sequence phase, and that the unusually large radius of the secondary cannot be explained by youth.  The only practical way to establish the age of J0516 would be to associate it with a cluster or moving group.  To do this, one would need good  proper motions and distance for J0516 and a large number of stars in the nearby field.  We know of no measurement of the proper motion of J0516 [for example, no proper motion measurement is given in the USNO-B catalog (Monet et al.\\ 2003)]. The distance can be estimated from our model, but the result depends on the extinction and on the assumed temperatures of the stars. Schuh et al.\\ (2003) determined a color excess of $E(B-V)=0.9\\pm 0.2$ mag from their model of the spectral energy distribution. Assuming $A_V=3.1E(B-V)$ and $A_K=0.114A_V$ (Cardelli, Clayton, \\& Mathis 1989), the $K$-band extinction is $0.318\\pm 0.017$ mag. The absolute $K$ magnitude of the binary was found using filter-integrated surface brightnesses computed from {\\sc NextGen} models by France Allard (private communication),   assuming the radii given in Table \\ref{tbl} and temperatures of $4200\\pm 200$ K. The average apparent magnitude in the $K$-band determined by the 2MASS survey is $K=13.113\\pm 0.039$ (see Schuh et al.\\ 2003). We find a distance of $d=753 \\pm 34$ pc. This distance is substantially larger than the distances to the other well-studied low-mass binaries mentioned earlier [e.g.\\ $d=140\\pm 8$ pc in the case of GU Boo \\citep{lop04}].   J0516 is clearly a very important system in that it allows for further observational constraints for the lower main sequence.  Owing to the enhanced variability, it would be worthwhile to obtain additional photometry to better define the ``quiescent'' phase and to better establish the eclipse profiles in quiescence.  In addition, moderate resolution but high signal-to-noise spectroscopy would be useful in order to further investigate the apparent lack of an H$\\alpha$ feature in either star and to check for possible correlations between the profile of the H$\\alpha$ feature (if any) and the activity level."
    },
    "0608/gr-qc0608017_arXiv.txt": {
        "abstract": "The general relativistic modifications to the resistive state in superconductors of second type in the presence of a stationary gravitational field are studied. Some superconducting devices that can measure the gravitational field by its red-shift effect on the frequency of radiation are suggested. It has been shown that by varying the orientation of a superconductor with respect to the earth gravitational field, a corresponding varying contribution to AC Josephson frequency would be added by gravity. A magnetic flux (being proportional to angular velocity of rotation $\\Omega$) through a rotating hollow superconducting cylinder with the radial gradient of temperature $\\nabla_r T$ is theoretically predicted. The magnetic flux is assumed to be produced by the azimuthal current arising from Coriolis force effect on radial thermoelectric current. Finally the magnetic flux through the superconducting ring with radial heat flow located at the equatorial plane interior the rotating neutron star is calculated. In particular it has been shown that nonvanishing magnetic flux will be generated due to the general relativistic effect of dragging of inertial frames on the thermoelectric current. ",
        "introduction": "Although the effects of gravitational field on the electromagnetic properties of superconductors have been widely discussed by a number of authors (see, for example, \\cite{dw}-\\cite{a99}) there exist unconsidered problems which are of physical and astrophysical interest. To the best of our knowledge theoretical investigation of the effects of general relativity on electrodynamics of laboratory II-type superconductors has not been done. This subject is very interesting, on one hand, from point of view to detect weak gravitational effects in the solar system since superconductors provide sensitive and accurate measurements.  On other hand it is important for astrophysics of magnetized relativistic compact objects since the core of neutron stars forms a matter in the superconducting state of II-type~\\cite{bpp69} arising from the realistic estimation that the superconducting protons coherence length ($2-6 fm$) for stellar matter is typically much smaller than the London screening length ($100-300 fm$). Magnetic flux vortices are the result of the type-II superconductivity in the inner crust and core region. The magnetic field of the neutron star is confined into individual vortices of flux $\\Phi_0=\\pi\\hbar c/e=2\\times 10^{-7}G\\cdot cm^2$. The number density of such vortices is $n_v=B/\\Phi_0\\approx 10^{19} B_{12}cm^{-2}$, where $B_{12}=B/10^{12} G$. The knowledge of the electromagnetic properties of flux vortices in gravitational field may be relevant for understanding some observable phenomena of pulsars~\\cite{sls00}. However, we should mention here that in recent papers (see, for example, \\cite{bmz04}) it was discussed inconsistency of the standard picture of the neutron star core, composed of a mixture of a neutron superfluid and a proton type-II superconductor with observations of long period precession in isolated pulsars.  The present work is the sequel of our previous investigation~\\cite{a99,a98} on gravitational and rotational effects on electromagnetic properties of conductors and superconductors. In our previous paper~\\cite{a99} we did not consider the electromagnetic properties of type-II superconductors in the gravitational field. One of the purpose of the paper is to study gravitational effects on type-II superconductivity. The second purpose of the paper is to examine a question on generation of magnetic field in superconductors arising from the effects of gravity or inertia. Since the mechanism of the generation of neutron star magnetic fields is under big astrophysical interest, we apply the discussed mechanism to the interior of slowly rotating neutron star in the simple toy model.  The paper is organized as follows. In the section~\\ref{sup2} we consider the general relativistic modifications to the electromagnetic properties of II-type superconductor when the Lorentz force created by transport current flowing in the superconductor exceeds the pinning force and the vortices started to move laterally to the current (This stage is called as resistive one due the appearance of the resistance from the energy dissipation.). This generalizes the resisitive state in II-type superconductors when the gravitational field is present. In addition, in subsection~\\ref{redshift} the weak general relativistic contribution on the frequency of radiation from II-type superconductor with transport current in an applied magnetic field is calculated. More concretely we propose here an experiment where the AC Josephson effect in a superconductor of II-type has a combined electric and gravitational origin.  The magnetic flux through superconducting hollow cylinder arising from interplay between thermally generated electric current and uniform rotation is obtained in section~\\ref{rotation}. In section~\\ref{star} new possible mechanism for magnetic field production inside a rotating relativistic star is discussed. The last section~\\ref{conclusion} summarizes our results.  We use here a space-like signature $(-,+,+,+)$ and a system of units in which $c=1$ (unless explicitly shown otherwise for convenience). Greek indices are taken to run from $0$ to $3$. ",
        "conclusions": "\\label{conclusion}  Here we investigated the effect of the gravitation on the electromagnetic properties of II-type superconductors and the effects of uniform rotation on superconductors with a heat flow. In particular we have shown that gravity can affect the alternating voltage in II-type superconductor. The proposal is presented which uses gravity to create an additional emf and consequently may  change the superconducting Josephson effect frequency. The frequency is becoming a function of the gravitational field. Thus the results on the electromagnetic properties of II-type superconductor embedded in gravitational field could lead to a way towards measuring of the tiny general relativistic effects in the earth conditions.  We have also derived the expression for magnetic flux through rotating hollow type-I superconductor of cylindrical shape when a temperature gradient is applied in radial direction between the inner and outer boundaries of the sample. In principle, the magnetic field due to the rotational effects on the electric current in the superconducting state should be observable in this type of laboratory experiment, although its magnitude is less than flux quantum for many orders.  The mechanism of the generation of magnetic field inside rotating neutron star due to the frame dragging effect has been considered. A possibility of magnetic field production has been shown in the toy model, where for the simplicity  superconducting matter has been considered as the disc in the equatorial plane."
    },
    "0608/astro-ph0608116_arXiv.txt": {
        "abstract": "I revisit the question of the adiabaticity of initial conditions for cosmological perturbations in view of the 3--year WMAP data. I focus on the simplest alternative to purely adiabatic conditions, namely a superposition of the adiabatic mode and one of the 3 possible isocurvature modes, with the same spectral index as the adiabatic component.  I discuss findings in terms of posterior bounds on the isocurvature fraction and Bayesian model selection. The Bayes factor (models likelihood ratio) and the effective Bayesian complexity are computed for several prior ranges for the isocurvature content. I find that the CDM isocurvature fraction is now constrained to be less than about $10\\%$, while the fraction in either the neutrino entropy or velocity mode is below about $20\\%$. Model comparison strongly disfavours mixed models that allow for isocurvature fractions larger than unity, while current data do not allow to distinguish between a purely adiabatic model and models with a moderate (\\ie, below about $10\\%$) isocurvature contribution.  The conclusion is that purely adiabatic conditions are strongly favoured from a model selection perspective. This is expected to apply in even stronger terms to more complicated superpositions of isocurvature contributions. ",
        "introduction": "The detailed nature of the initial conditions for cosmological perturbations is one of the open questions in cosmology. The exquisite precision of the WMAP measurement of the first acoustic peak location in the cosmic microwave background (CMB) temperature power spectrum ($\\ell = 220.7 \\pm 0.7$, see \\cite{Hinshaw:2006ia}) is a strong indication in favour of adiabatic initial conditions, which predict for the first peak $\\ell \\approx 220$. The alternative possibility of cold dark matter (CDM) isocurvature initial conditions excites a sine wave (rather than the cosine excited by adiabatic conditions) in the photon--baryon plasma, resulting in a first acoustic peak displaced by half a period to $\\ell \\approx 330$, see e.g.\\ \\cite{Trotta:2004qj,Durrer:2004fx}. Furthermore, the ratio of the Sachs--Wolfe plateau for $\\ell \\lsim 50$ to the height of the peak is very different for the two modes.  A few years ago, Bucher and collaborators introduced two new isocurvature modes, called ``neutrino density'' (or, more appropriately, ``neutrino entropy'') and ``neutrino velocity'' modes \\citep{Bucher:1999re}. They are characterized by a non--vanishing initial entropy perturbation in the neutrino sector and by a non--vanishing difference in the neutrino to photon velocity, respectively. A superposition of the adiabatic and the three isocurvature modes (cold dark matter, neutrino entropy and neutrino velocity) constitutes the most general initial conditions for the perturbations, at least if the Universe is radiation dominated in its early phase \\citep{Trotta:2004qj}. A baryon isocurvature mode is observationally indistinguishable from a CDM isocurvature one \\citep{Bucher:2000hy,Gordon:2002gv} and can thus be neglected without loss of generality.  Allowing for the most general type of initial conditions has two effects on cosmological parameter extraction from CMB measurements. First, the extra parameters associated with the initial conditions introduce severe degeneracies which limit our ability to reconstruct the cosmology \\citep{Trotta:2001yw,Trotta:2002iz}, even though this can fortunately be remedied by using polarization information \\citep{Bucher:2000hy,Trotta:2004sm}. Secondly, it becomes difficult to constrain the type of initial conditions, i.e.\\ the amount of isocurvature contributions allowed on top of the predominantly adiabatic mode \\citep{Moodley:2004nz}.  Recent works have investigated general isocurvature admixtures in the initial conditions \\citep{Beltran:2005gr,Moodley:2004nz,Bean:2006qz}. In this work I focus on the simplest alternative to a purely adiabatic power spectrum, namely a superposition with one totally (anti--)correlated isocurvature mode at the time with the same spectral index as the adiabatic one. This is partly motivated by models for the generation of initial conditions such as the curvaton (see \\eg~\\cite{Gordon:2002gv,Lyth:2001nq} and references therein), where this kind of scenario arises as a generic prediction. A second justification comes from the model selection approach used in the second part of this work. In comparing the simplest (\\ie, purely adiabatic) scenario with a more complex one, it makes sense to start by adding a minimal number of extra parameters, and see whether the extended model is justified by the data. This model selection perspective has been recently advocated by \\cite{Beltran:2005xd,Trotta:2005ar}.  This paper is organized as follows: in section \\ref{sec:isoc} we introduce the parameterization of the initial condition parameters space we are considering, while in section \\ref{sec:bayes} we review some concepts of Bayesian statistics and in particular the model selection approach. We present our results in terms of parameters constraint and model comparison outcome in section \\ref{sec:res} and offer our conclusions in \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have submitted the question of the type of initial conditions for cosmological perturbations to renewed scrutiny in the light of WMAP 3--year data. We have focused on the simplest and well motivated alternative to a purely adiabatic model, namely an admixture of one totally (anti--)correlated isocurvature mode at the time, with the same spectral tilt as the adiabatic one.  We have derived posterior bounds on the isocurvature fraction from WMAP 3--year data combined with other CMB measurements and SDSS. We have constrained the isocurvature fraction in the CDM mode to be less than about $10\\%$, while the maximum allowed neutrino isocurvature contribution (either density or velocity) is about $20\\%$.  Bayesian model selection tends to favour purely adiabatic initial conditions with strong odds ($>20:1$) when compared to models predicting isocurvature fractions larger than unity. For such models -- having a large prior range on the isocurvature fraction -- we have shown that the data can support 7 parameters, but that only 6 of them are required, with no need to include isocurvature modes from a model selection point of view. These findings confirm the conclusions of~\\cite{Kunz:2006mc}. However, mixed models that limit the isocurvature contribution to less than about $10\\%$ cannot presently be ruled out. We have shown that the constraining power of the data for this class of models is insufficient, and therefore we must hold our judgement until better data becomes available. These findings are however dependent on the parameterization chosen for the isocurvature fraction, that in this work is motivated by the curvaton scenario. The question of prior selection will be further addressed in a future publication.  It is reasonable to expect that the same conclusion would apply in even stronger terms to the case of more complicated models, \\eg~those involving a superposition of different isocurvature modes at the same time, or with arbitrary correlations among them. In fact, more complicated models (such as the class considered by \\cite{Bean:2006qz}) ought to be even more disfavoured because of their larger volume of wasted parameter space. At present, Occam's razor is perfectly compatible with the simplest possibility, namely purely adiabatic initial conditions.  We think that this model comparison approach can be a useful complement to parameter constraints analysis, and that it can offer valuable guidance in building models for the generation of primordial perturbations.  {\\em Acknowledgments} R.T. is supported by the Royal Astronomical Society through the Sir Norman Lockyer Fellowship. I am grateful to Andrew Liddle and David Parkinson for comments. I acknowledge the use of the package \\texttt{cosmomc}, available from \\texttt{cosmologist.info} and the use of the Legacy Archive for Microwave Background Data Analysis (LAMBDA). Support for LAMBDA is provided by the NASA Office of Space Science."
    },
    "0608/astro-ph0608446_arXiv.txt": {
        "abstract": "We present first results of a spectroscopic survey targeting $K$-selected galaxies at $z=2.0-2.7$ using the GNIRS instrument on Gemini-South. We obtained near-infrared spectra with a wavelength coverage of 1.0--2.5~$\\rm\\mu$m for 26 $K$-bright galaxies ($K<19.7$) selected from the MUSYC survey using photometric redshifts. We successfully derived spectroscopic redshifts for all 26 galaxies using rest-frame optical emission lines or the redshifted Balmer/4000\\AA\\ break. Twenty galaxies have spectroscopic redshifts in the range $2.0<z<2.7$, for which bright emission lines like \\ha\\ and [O\\,{\\sc iii}] fall in atmospheric windows. Surprisingly, we detected no emission lines for nine of these 20 galaxies. The median $2\\sigma$ upper limit on the rest-frame equivalent width of \\ha\\ for these nine galaxies is $\\sim$10\\,\\AA. The stellar continuum emission of these same nine galaxies is best fitted by evolved stellar population models. The best-fit star formation rate (SFR) is zero for five out of nine galaxies, and consistent with zero within $1\\sigma$ for the remaining four.  Thus, both the \\ha\\ measurements and the independent stellar continuum modeling imply that 45\\% of our $K$-selected galaxies are not forming stars intensely. This high fraction of galaxies without detected line emission and low SFRs may imply that the suppression of star formation in massive galaxies occurs at higher redshift than is predicted by current CDM galaxy formation models. However, obscured star formation may have been missed, and deep mid-infrared imaging is needed to clarify this situation. ",
        "introduction": "Observations imply that stellar populations of high-mass galaxies were formed at higher redshift than those of low-mass galaxies \\citep[e.g.,][]{co96,ju05}. Recent hierarchical CDM models are able to produce dead, massive galaxies at high redshift by incorporating feed-back from active galactic nuclei \\citep[AGN; e.g.,][]{cr06,bo06,dl06,ho06}. In order to observationally determine when and how star formation in massive galaxies was suppressed it is necessary to identify and study these objects out to the highest redshifts.  Recently, massive and apparently dead galaxies have been identified at $z>1.5$ \\citep[e.g.,][]{mc04,sa05,da05,la05,re05,re06,pa06}. Beyond $z=2$ most studies rely on photometric redshifts and broad-band colors to identify these galaxies. However, as dust and age have similar effects on the broad-band spectral energy distribution (SED), the star formation rates (SFRs) are often not well constrained. Furthermore, their photometric redshifts are not well calibrated, as only a few at $z>2$ have spectroscopic redshifts \\citep{da05}. Thus spectroscopic redshifts and independent stellar population diagnostics are needed to determine the prevalence of ``red and dead'' galaxies beyond $z=2$.  Spectroscopic confirmation of non-star-forming galaxies at $z>2$ is complicated due to their faint rest-frame UV emission and lack of nebular emission lines. Deep near-infrared (NIR) spectroscopy provides the best option to confirm such galaxies at $z>2$, as their relatively bright rest-frame optical luminosity allows direct detection of the stellar continuum. The optical continuum shape, and in particular the Balmer/4000 \\AA\\ break, can be used to derive redshifts for galaxies without emission lines, and provides independent constraints on stellar populations \\citep{kr06}.  To study a high-redshift spectroscopic sample that is not biased towards galaxies with bright emission lines, we are conducting a NIR spectroscopic survey of $K$-selected galaxies with photometric redshifts $z\\sim2.3$. Here we report on a surprising result of our survey: the large fraction of galaxies with no detected emission lines. Throughout the letter we assume a $\\Lambda$CDM cosmology with $\\Omega_{\\rm m}=0.3$, $\\Omega_{\\rm \\Lambda}=0.7$, and $H_{\\rm 0}=70$~km s$^{-1}$ Mpc$^{-1}$. All broadband magnitudes are given in the Vega-based photometric system. ",
        "conclusions": "We find that nine out of 20 $K$-selected galaxies at $2.0<z_{\\rm spec}<2.7$ have no detected \\ha\\ emission (\\wha\\,$\\lesssim$~10\\AA) and are best fit by stellar population models with low specific SFRs ($\\sim\\,0.004\\rm\\,Gyr^{-1}$), implying a fraction of galaxies with strongly suppressed star formation of $45^{+18}_{-12}$\\,\\%.  The quoted uncertainty is derived assuming Poisson statistics, and does not include the following systematic errors and caveats. First, our $z_{\\rm phot}$ selection criterion could introduce biases, as systematic errors in photometric redshift may correlate with SED type\\footnote{This effect may be small, as the distribution of specific SFRs of the six objects at $z_{\\rm spec}<2$ is similar to that of the 20 remaining objects}. Also, the $K$-band selection criterion could bias our sample, as for starburst galaxies strong emission lines can contribute significantly to the $K$-band flux \\citep{er06a}. We do not expect this bias to be strong, as the median contribution of the emission lines to the $K$-band is only 0.04 mag for the eleven emission line galaxies. Third, incompleteness may play a role, as we observed only $\\sim$20\\% of the galaxies that meet the selection criteria. We note, however, that according to RS- and KS-tests our $K$-selected sample has a similar distribution of rest-frame $U-V$-colors as the large mass-limited sample ($>10^{11} M_{\\odot}$) by \\cite{vd06}, when applying the same $K$-magnitude cut. Furthermore, the assumption that \\ha\\ emission is only due to star formation may lead us to underestimate the fraction.  Finally, and perhaps most importantly, we may miss star formation with very high rest-frame optical extinction. Although most local ULIRGs have high integrated \\wha\\ \\citep[$W_{\\rm H\\alpha+[N\\,{\\sc ii}]} \\sim 90$\\AA,][]{li95}, for some the star-burst regions are almost completely obscured (e.g., Arp 220 has $W_{\\rm H\\alpha+[N\\,{\\sc ii}]} = 18$\\AA) and these objects might be misinterpreted in our analysis. Available 24 $\\mu$m imaging on DRG samples shows that 30-50\\% have no MIPS counterpart \\citep{we06,pa06,re06}, and \\citet{re06} finds that red galaxies with low MIPS fluxes typically have low specific SFRs. To resolve this issue it is necessary to combine our spectra with deep mid/far-infrared imaging to detect hidden star formation.  It is interesting to compare our fraction of galaxies with strongly suppressed star formation to previous studies\\footnote{The fractions of galaxies with suppressed star formation among DRGs and $K$-selected galaxies are similar in our sample}. Our result is consistent with the fraction found by \\cite{la05}, as they identified three ``red and dead'' galaxies from a sample of 11 DRGs in the HDF-S, and with the study of \\cite{re06}, who find that seven out of 24 DRGs at $1.5<z<2.6$ are not detected at 24 $\\mu$m, and have an average low specific SFR of 0.05 Gyr$^{-1}$. However, our fraction is significantly higher than the fractions found by \\cite{da04} and \\cite{pa06}. Using the $BzK$ criterion, \\cite{da04} identify all their 11 $K$-selected galaxies at $2.0<z<2.7$ as star-forming galaxies. The difference may be partly explained by different definitions, as $\\sim$4 of our 9 galaxies with suppressed star formation galaxies would have been identified as star-forming galaxies by the $BzK$ criterion. \\cite{pa06} finds that $\\sim$10\\% of 153 DRGs at $1.5\\le z_{\\rm phot}\\le 3$ show no signs of current star formation. Again, the selection criteria could play a role; none of our galaxies would have been classified as ``dead'' by \\cite{pa06}, as they apply the following criteria: age $>1 \\rm Gyr$ , $E(B-V)\\le0.1$, age$/\\tau>3$ and no X-ray or 24$\\mu$m detection. The sample selection could also be a factor, as the contribution of dusty star-forming galaxies to the DRG population is expected to be higher at $z<2$. Furthermore, the CDF-S field -- in which both the studies by \\cite{pa06} and \\cite{da04} were performed -- may be atypical \\citep{vd06}. We stress that our study is the first that is based on spectroscopic redshifts, and that this might also account for differences in the obtained fractions. This will be explored in a future paper.  We note that most current CDM galaxy formation models fail to produce the high fraction of red galaxies at the massive end as found by \\cite{vd06}. One easy way to solve this is by allowing more dust in the galaxies, but we have shown here that a large fraction of galaxies has low specific SFRs, and these are generally absent in these models at $z>2$ \\citep[e.g.,][]{so04,na05,ka06}. Our results may indicate that the suppression of the star formation in the most massive galaxies occurs at higher redshift than has been predicted by current models. In this context, it is interesting to note that not all of the line-emission in our sample may be due to star formation: from our low-resolution GNIRS spectra it appears that several of the emission-line galaxies with low specific SFRs exhibit high [N\\,{\\sc ii}]/H$\\alpha$ ratios, possibly indicating AGN. High resolution spectra of the line-emitting objects in our sample will be discussed in a future paper (M.~Kriek et al., in preparation)."
    },
    "0608/astro-ph0608670_arXiv.txt": {
        "abstract": "The Pierre Auger Observatory detects ultra-high energy cosmic rays by implementing two complementary air-shower techniques. The combination of a large ground array and fluorescence detectors, known as the \\textit{hybrid} concept, means that a rich variety of measurements can be made on a single shower, providing much improved information over what is possible with either detector alone. In this paper I describe the hybrid reconstruction approach and the latest hybrid measurements.  \\vspace{1pc} ",
        "introduction": "The energy spectrum of primary cosmic rays extends over more than $12$ decades in energy. It follows, over a large range, a simple power law indicating its non-thermal nature. The highest energy cosmic rays (which means particles with energies above $10$~EeV $\\equiv 10^{19}$~eV) must be produced by the most extreme non-thermal process in the Universe. Their energy spectrum, arrival directions, and composition can be inferred from air shower observations. However, the origin of these particles is still unknown and there is not a consensus whether there is a steepening in the energy spectrum at around $60$~EeV. Determining the shape of the spectrum in this energy range in sufficient detail requires an enormous detector exposure. The Pierre Auger Observatory was designed to study these ultra-high energy cosmic rays. Although under construction, the Auger Southern site is the largest cosmic ray detector in the world.  A fundamental characteristics of the Pierre Auger Observatory is its capability of \\textit{hybrid} reconstruction of cosmic ray showers~\\cite{Mostafa:2005kd}. Two independent detectors, the Surface Detector (SD), which samples the shower particles at ground, and the Fluorescence Detector (FD), which collects the fluorescence light emitted by the shower particles along their path in the atmosphere, are able to measure the energy and direction of the same cosmic ray shower. The enhanced capabilities of the Auger hybrid detector are examined in this paper.  The Pierre Auger Observatory was designed to observe, in coincidence, the shower particles at ground and the associated fluorescence light generated in the atmosphere. This is achieved with a large array of water Cherenkov detectors coupled with air-fluorescence detectors that overlook the surface array. It is not simply a dual experiment. Apart from important cross-checks and measurement redundancy, the two techniques see air showers in complementary ways. A single air shower is detected 3-dimensionally. The ground array measures the 2-dimensional lateral structure of the shower at ground level, with some ability to separate the electromagnetic and muon components. The fluorescence detector records the longitudinal pro{f}ile of the shower during its development through the atmosphere.  A \\textit{hybrid} event is an air shower that is simultaneously detected by the fluorescence detector and the ground array. The Observatory was originally designed and is currently being built with a \\textit{cross--triggering} capability. Data are recovered from both detectors whenever either system is triggered. If an air shower independently triggers both detectors the event is tagged accordingly. There are also cases where the fluorescence detector, having a lower energy threshold, promotes a sub--threshold array trigger. Surface stations are then matched by timing and location. This is an important capability because these sub--threshold hybrid events would not have triggered the array otherwise. The geometrical reconstruction of the air shower's axis is accomplished by minimizing a $\\chi^{2}$ function involving data from all triggered elements in the eye and on the ground. The reconstruction accuracy is far better than the ground array counters and the single eye could achieve independently~\\cite{Bonifazi:2005ns}.  The combination of the air fluorescence measurement and the particle detection on the ground provides an absolute energy calibration. The FD measurements determine the longitudinal development profile, whose integral is proportional to the total energy of the electromagnetic particle cascade. The SD independently estimates the shower energy by evaluating the particle density at $1000$~m from the core. This ground parameter, called $S(1000)$, is estimated from the surface stations by fitting the measured densities to a lateral distribution function~\\cite{Ghia:2005bn}. A conversion factor that relates $S(1000)$ to the shower energy (based on the FD information) is extracted from hybrid events. This reduces significantly the dependence on air shower models and on assumptions of the primary composition.  The Pierre Auger Collaboration started the construction of the Southern site in $2002$ in the city of Malarg\\\"{u}e, which is located at an elevation of $1400$~m a.s.l. in the province of Mendoza, Argentina. After a successful operation of a prototype experiment~\\cite{Abraham:2004dt}, the Southern Observatory started operation in \\textit{hybrid} production mode in January, 2004. Surface stations have a $100$\\% duty cycle, while fluorescence eyes can only operate on clear moonless nights. Both surface and fluorescence detectors have been running simultaneously $14\\%$ of the time. The number of hybrid events represents $10\\%$ the statistics of the surface array data. The Southern site will be completed in 2007. Another important objective is to obtain a uniform exposure over the full sky. This will be achieved by constructing a second instrument in the Northern hemisphere. Each site is conceived to cover an area of $3000$~km$^{2}$ in order to collect about $1$ event per week and site above $60$~EeV. The Northern Observatory is planned to be sited in the U.S. ",
        "conclusions": "The Pierre Auger Observatory is a hybrid detector with excellent capabilities for studying the highest energy cosmic rays. Much of its capability stems from the accurate geometric reconstruction it achieves. The shower geometry is reconstructed combining information from the eyes and the ground detectors. Arrival directions are determined to a small fraction of a degree and the shower core is located to an accuracy of about $60$~m~\\cite{Bellido:2005aw,Bonifazi:2005ns}.  The construction of the Southern Observatory is well under way. Eighteen FD telescopes and more than $60\\%$ of the surface array are in operation taking data routinely. At the present rate of deployment, construction will be finish in mid 2007. Detectors are performing very well and the first results are very encouraging.  Emphasis is placed on hybrid analysis that provide unprecedented quality in geometry, energy, and mass reconstruction. Of utmost importance for the near future will be the determination of the energy spectrum to study the GZK feature, and the search for anisotropies in arrival directions. It is important to note that both SD and FD techniques have different systematics, and results are preliminary at this stage while the Observatory is under construction. The possibility of studying the same set of air showers with two independent methods is valuable in understanding the strengths and limitations of each technique. The hybrid analysis bene{f}its from the calorimetry of the fluorescence technique and the uniformity of the surface detector aperture.  In parallel to the completion of the Southern Observatory and to the analysis of data towards the first scientific results, R\\&D has started for the development of the Northern site in the US."
    },
    "0608/hep-th0608071_arXiv.txt": {
        "abstract": "We consider global topological defects in symmetry breaking models with a non-canonical kinetic term. Apart from a mass parameter entering the potential, one additional dimensional parameter arises in such models --- a ``kinetic'' mass. The properties of defects in these models are quite different from ``standard'' global domain walls, vortices and monopoles, if their kinetic mass scale is smaller than their symmetry breaking scale. In particular, depending on the concrete form of the kinetic term, the typical size of such a defect can be either much larger or much smaller than the size of a standard defect with the same potential term. The characteristic mass of a non-standard defect, which might have been formed during a phase transition in the early universe, depends on both the temperature of a phase transition and the kinetic mass. ",
        "introduction": "% Topological defects are the class of solitons which may form as a result of symmetry-breaking phase transitions. Defects may be zero-, one- or two-dimensional depending on the topological properties of the minima of the effective potential. In the condensed matter physics topological defects arise rather commonly. The well-known examples are the flux tubes in superconductors \\cite{Abr} and vortices in superfluid helium-4.  In the cosmological context topological defects attract much interest because they might appear in a rather natural way during the phase transitions in the early universe. The discrete symmetry breaking leads to the appearance of domain walls, while the breaking of a global or a local $U(1)$-symmetry is associated with global \\cite{glstr} and local \\cite{lostr} cosmic strings correspondingly. The localized defects or monopoles arise in the gauge model possessing $SO(3)$ symmetry which is spontaneously broken to $U(1)$ \\cite{Pol,tHo}. Note that the gauge monopoles naturally arise in Grand Unified Theories (GUT); e.~g. in a symmetry breaking phase transition $G\\to H\\times U(1)$ where $G$ is a semi-simple non-Abelian group, monopoles inevitably appear. Similar to gauge monopoles, global monopoles can arise if a global symmetry is broken suitably.  Adding non-linear terms to the kinetic part of the Lagrangian makes it possible for defects to arise without a symmetry-breaking potential term \\cite{Sky}. Non-canonical kinetic terms are rather common for effective fields theories. During the last years Lagrangians with non-canonical fields were intensively studied in the cosmological context. So-called $k$-fields were first introduced in the context of inflation \\cite{k-inflation} and then $k$-essence models were suggested as solution to the cosmic coincidence problem \\cite{k-essence}. Tachyon matter \\cite{tachyon} and ghost condensate \\cite{ghost} are also examples of non-canonical fields in cosmology. A very interesting application of $k$-fields is the explanation of dark matter by the self-gravitating coherent state of $k$-field matter \\cite{Halo}. The production of large gravitation waves in the models of inflation with nontrivial kinetic term was considered in \\cite{MukVik}. The effects of scalar fields with non-canonical kinetic terms in the neighborhood of a black hole were investigated in \\cite{BH}.  In this paper we study the properties of topological defects arising in models with non-canonical kinetic terms. We dub such defects as $k$-defects to distinguish them from standard ones. The existence of non-trivial configurations is ensured by the symmetry-breaking potential term, as in the case of ``usual'' domain walls, vortices and monopoles. However, the kinetic term is non-trivial. We choose it to have asymptotically a standard form for small values of gradients, and to be different from the standard one at large gradients. The asymptotically canonical form at small gradients guarantees the standard behavior of a system for small perturbations and is needed to avoid problems with singularities and/or non-dynamical behavior of fields at zero gradients. By contrast, the deviation of a kinetic term from a standard one at large gradients makes the properties of topological defects quite different (depending on parameters of the Lagrangian) from the standard ones because of non-zero gradient of fields inside defects.  The paper is organized as follows. In Sec.~\\ref{SModel}, we describe our model and derive its equations of motion. The general properties of topological $k$-defects are studied in Sec.~\\ref{SGeneral}. In Sec.~\\ref{SConstraints} we find constraints on the parameters of the model. Solutions for domain wall, global vortex and global monopole are found in Sec.~V. We summarize and discuss results and cosmological applications in the concluding Sec.~\\ref{SDiscussion}. ",
        "conclusions": "\\label{SDiscussion}% In this work we considered the properties of topological defects which arise in models with a symmetry breaking potential and non-canonical kinetic term. The principal difference of the considered models from the ``standard'' ones allowing the existence of topological defects is the presence of non-standard kinetic part $K(X)$, where $X$ is a canonical kinetic term (\\ref{X}). The non-linearity of the kinetic term inevitably leads to the appearance of a new scale in the action [see Eq.~(\\ref{act})], in addition to the mass parameter in the potential term. Depending on the relation between the kinetic and potential masses and on the precise form of the function $K(X)$ the basic properties of topological defects can vary drastically.  We considered general properties of defects and restrictions on models having in mind the rather general case of a kinetic term with the asymptotic behavior (\\ref{model}). The constraints on the parameters of the model in this case are shown in Fig.~\\ref{constr}. It turns out, that the properties of $k$-defects are quite different from standard defects. In particular, the size of a defect depends both on the kinetic and a potential masses, cf.~Eq.~(\\ref{Lph}). Depending on the parameter $\\a$ in (\\ref{model}), a $k$-defect can be either smaller or greater than the standard defect.  As particular examples we studied in detail two concrete models having non-canonical kinetic terms (\\ref{m1}) and (\\ref{m2}). The field profiles of domain walls for different choices of $K(X)$ are shown in Fig.~\\ref{wall}: in the region of parameters where the deviation of a kinetic term from the standard one is small, the profiles are practically identical (left panel); if the non-linearities in $X$ are strong inside a defect then $k$-defect has quite different properties (right panel). The same result applies for global cosmic strings (see Fig.~\\ref{string}) and global monopoles.  In this paper we did not consider explicitly the problem of gauge models with non-standard kinetic term. Although this problem deserves a separate investigation, some basic properties of local $k$-defects can be inferred, basing on the present work. In particular, one can estimate the total the total mass of the $k$-monopole $M_k$ as the typical energy density (\\ref{ed}) multiplied by the volume of the core of the monopole. Using (\\ref{LS}) one obtains: \\begin{equation} \\label{MkM} M_k\\sim M_{st} \\left[\\l^{1/4}\\frac{\\eta}{M}\\right]^{6(\\a-1)/\\a}, \\end{equation} where $M_{st}$ is the mass of a standard monopole. In a similar way we find the tension of a cosmic $k$-string $\\mu_k$: \\begin{equation} \\label{Mks} \\mu_k\\sim \\mu_{st} \\left[\\l^{1/4}\\frac{\\eta}{M}\\right]^{4(\\a-1)/\\a}, \\end{equation} where $\\mu_{st}$ is the tension of a standard cosmic string. If we consider the case of different mass scales, $M\\ll\\eta$ [compare to our requirement (\\ref{Xcondg})], then for $\\a>1$ we have $M_k \\gg M_{st}$, while for $\\a<1$ we obtain $M_k \\ll M_{st}$. The same conclusion is true for tension of local cosmic $k$-strings. We can see that, similar to global $k$-defects, the properties of local $k$-defects may drastically differ from standard local defects.  This result may have important consequences for cosmological applications. In particular, standard defects which might have been formed during phase transitions in the early universe have a mass scale directly connected to the temperature of a phase transition $T_c$: $M_{st}\\sim \\eta \\sim T_c$ for monopoles and $\\mu_{st}\\sim \\eta^2 \\sim T_c^2$ for cosmic strings. In contrast, the mass scale of a resulting $k$-defect depends both on $T_c$ and the additional mass parameter, the kinetic mass $M$. This effect has an interesting application to cosmic strings. The new three-year WMAP data \\cite{WMAP} give an upper limit on the cosmic string tension: $G\\mu<2\\times 10^{-7}$ \\cite{csc1}, or even stronger: $G\\mu<3\\times 10^{-8}$ \\cite{csc2}; while the theoretical estimation of the tension for GUT cosmic strings gives $G\\mu\\sim 10^{-6}-10^{-7}$. This means that GUT cosmic strings are almost ruled out by the observations. However, if physics at the GUT scale involves non-standard kinetic terms, then the GUT phase transition may have lead to the formation of cosmic strings with smaller tension, $G\\mu_k \\ll 10^{-6}$, thereby evading conflicts with the present observations."
    },
    "0608/astro-ph0608514_arXiv.txt": {
        "abstract": "The NEMO Collaboration is conducting an R\\&D activity towards the construction of a Mediterranean km$^3$ neutrino telescope. In this work, we present the results of Monte Carlo simulation studies on the capability of the proposed NEMO telescope to detect and identify point-like sources of high energy muon neutrinos. ",
        "introduction": "\\label{intro}  The detection of high energy neutrinos is considered one of the most promising means to investigate non-thermal processes in the Universe. A first generation of small scale detectors has been realized (NT-200 \\cite{Baikal} in the Baikal lake and AMANDA \\cite{Amanda} at the South Pole), demonstrating the possibility to use the \\v{C}erenkov technique to track high energy neutrinos. Besides, these experiments have set limits on neutrino fluxes. Other small scale detectors are at different stage of realization (ANTARES \\cite{antares} and NESTOR \\cite{nestor}).  Actual expectations on neutrino fluxes, mainly based on the measured cosmic ray fluxes and the estimated fluxes for several high energy sources from theoretical models \\cite{flussinu}, require detectors of km$^3$ scale. Following the success of AMANDA, the largest operating detector, the realization of the IceCube km$^3$ detector \\cite{icecube} has started at the South Pole. On the other hand, many issues, as the full sky coverage, strongly support the construction of a km$^3$ scale detector in Mediterranean Sea. ",
        "conclusions": "The possibility to detect TeV muon neutrinos from point-like sources with the proposed NEMO-km$^3$ underwater \\v{C}eren\\-kov neutrino telescope has been investigated. In particular Monte Carlo simulations were carried out to determine the expected response of the km$^3$ telescope.  Our simulations show that it could be possible to observe the Moon shadow in about 100 days. This detection provides a measurement of the detector angular resolution and of the detector pointing accuracy.  We also computed the detector sensitivity to muon neutrinos from point-like sources, defined as the minimum flux detectable with respect to the atmospheric muon and neutrino background. The dependence of the sensitivity on the neutrino spectral index and energy range, on the source declination and on years of data taking has been studied.  Finally, we consider the case of two particular sources: microquasar LS 5039 and SNR RX J1713.7-3946; both observed in the TeV gamma-ray region. For each source, we compute the detector sensitivity and the expected number of source events compared with the background. Our results show that, assuming present predictions of TeV neutrino fluxes, the proposed NEMO telescope could identify both sources in a few years of data taking."
    },
    "0608/astro-ph0608387_arXiv.txt": {
        "abstract": "We present an analysis of thermal emission from comet 162P/Siding Spring (P/2004 TU$_{12}$) measured during its discovery apparition in December 2004.  The comet showed no dust coma at this time, so we have sampled emission from the comet's nucleus.  Observations using the ``Mid-Infrared Spectrometer and Imager\" (``MIRSI\") were performed at NASA's Infrared Telescope Facility, where the peak of the comet's spectral energy distribution was observed between 8 and 25 $\\mu$m.  In combination with the three near-infrared spectra presented by Campins {\\it et al.} (2006, {\\it AJ}, this issue) that show the Wien-law tail of the thermal emission, the data provide powerful constraints on surface properties of the nucleus. We find that the nucleus's effective radius is $6.0\\pm0.8$ km. This is one of the largest radii known among Jupiter-family comets, which is unusual considering the comet was discovered only recently.  Its geometric albedo is $0.059\\pm0.023$ in H-band, $0.037\\pm0.014$ in R-band, and $0.034\\pm0.013$ in V-band. We also find that the nucleus of 162P has little infrared beaming, and this implies the nucleus has low thermal inertia. Including all near-IR spectra yields a beaming parameter $\\eta$ of $1.01\\pm0.20$. This result is in agreement with others showing that cometary nuclei have low thermal inertia and little infrared beaming. If confirmed for many nuclei, the interpretation of radiometry may not be as problematic as feared. ",
        "introduction": "A narrative of the discovery and morphological evolution of comet 162P is given by \\citet{cam06} in a companion paper (hereafter Paper I).  The thin tail and lack of coma strongly suggested that observations of this comet would reveal flux from the nucleus. Since opportunities to directly study a cometary nucleus are infrequent, our group set up a multiwavelength observing campaign to take advantage of the comet's unusual behavior and probe the reflectance and thermal properties of the nucleus.  Paper I addressed the near-infrared spectroscopic properties; in this paper we describe the thermal emission and the albedo, which include observations from the Wien-law tail through the black-body peak in wavelength. ",
        "conclusions": "\\subsection{Absolute Magnitude and Phase Slope Parameter}  As mentioned, we have derived a value for the R-band absolute magnitude $H_R$  based on photometry that covers a range of phase angles. However the rotational context is not known, and for a proper analysis of the phase law, one needs to account for the variability of the magnitude due to the nucleus's rotation. Our lack of this information introduces an uncertainty into our derivation of $H_R$ and $G$. There is a possibility that if the nucleus's light curve amplitude is large then we are being fooled, and these two quantities may have quite different values.  In particular this is a concern because a slope parameter of $G=0.15$ is somewhat higher than has been found for several other cometary nuclei and outer Solar System objects \\citep{fer00,sr02, sj02,bau03,js04}, although it is comparable to that found by \\cite{bur04} for comet 19P/Borrelly.  The problem is compounded by the fact that the mid-IR and near-IR observations both took place at moderate phase angle ($46^\\circ$ and $49^\\circ$).  If 162P's slope parameter is indeed much smaller than $0.15$, this would mean that $H_R$ is brighter than we have calculated.  This, in turn, would mean that the nucleus's true albedo is higher and effective radius is lower than we have derived here.  For example, if the nucleus's $G=-0.05$, then the infrared observations would imply that $p_R\\approx0.06$ and $p_H\\approx0.09$.  Such an albedo would be somewhat higher than the other albedos known for Jupiter-family comets \\citep{lam05}, and be closer to those of Centaurs and transneptunian objects \\citep{stan05}. Furthermore the effective radius would be $R\\approx4.7$ km, admittedly still a fairly large nucleus.  \\subsection{December 11th Spectrum}  The fact that the December 11th spectrum is somewhat different from those of the other dates suggests that we should explore why this may be. First we can check the rotational context of the observations. The rotation period of the comet is about 32.78 hours (G. Masi, private communication), so we can calculate the relative rotational phase between the four epochs. If the December 3rd spectrum occurred at rotational phase $\\phi=0$, then at the other epochs, $\\phi=0.13$ for the December 10th spectrum, $\\phi=0.32$ for the December 11th spectrum, and $\\phi=0.59$ for the December 27th photometry.  Therefore there is no clear conclusion we can draw about what effect the rotation may have had on our interpretation of the modeling results. For example if the December 11th spectrum had been at a much different rotational phase compared to the other three datasets, then we would have reason to disregard it in the final analysis, but this is not the case. (We admit that if the rotation period is wildly off from 32.78 hours then we will have to readdress this argument.) In any case the lack of rotational context for the visible-wavelength photometry introduces extra uncertainty in the problem.  Another explanation may be secular changes in the comet's activity. For example, perhaps there was an outburst of activity sometime between December 10.25 and December 11.88.  This would have formed a dust coma that would affect the spectral measurements. If the dust grains had sizes that were approximately 1 to 10 $\\mu$m in scale, then they would have contributed extra reflected sunlight but little thermal emission in K-band. This is because the grains would have been isothermal and would not have reached the hotter temperatures that are achieved at and near the nucleus's subsolar point. Modeling such a spectrum on the assumption that we were only seeing light from the nucleus would yield higher albedos than normal -- which is exactly what is seen in the December 11th spectrum.  If this scenario is correct, then we should not include this spectrum in the derivation of the properties of the nucleus.  Finally, we note that the signal-to-noise ratio ($S/N$) of the spectrum on December 11th is somewhat lower than that from December 3rd and 10th. The $S/N$ of this spectrum is still good and cannot be rejected for that reason, but if we were to take just the two higher $S/N$ spectra we would find that the acceptable range of $\\eta$ is $0.98\\pm0.17$ (as can be found from Fig. 4). This is not too different from our solution in \\S 3.3.  Furthermore there would be no difference in our range of acceptable albedos $p_H$ (as can also be seen from Fig. 4). The acceptable range of radii $R$ would shrink by only a few tenths of a kilometer (as can be seen from Fig. 2).  In short,  there is no conclusive evidence letting us explain why the December 11th spectrum is different, and also no strong reason why it should be removed from consideration.  \\subsection{Thermal Properties}  Only recently have there been sufficient thermal measurements of cometary nuclei that can constrain the thermal inertia.  In particular, the thermal emission from the nucleus of 9P/Tempel 1 strongly indicates that the beaming parameter $\\eta$ is near unity, and that the thermal inertia is consistent with zero. This comes from both remote observations by \\cite{lis05} and from spacecraft observations reported by \\cite{ahe05} from the {\\sl Deep Impact} mission. \\cite{sod04} analyze spatially-resolved {\\sl Deep Space 1} near-IR spectra of comet 19P/Borrelly's nucleus and find that the nucleus's temperature is consistent with a similar thermal behavior.  Lastly, observations of comet 2P/Encke by Fern\\'andez et al. (2004; and in preparation) indicate that that nucleus has a near unity beaming parameter as well. This is based on mid-IR, near-IR, and visible observations, just as in our present analysis of 162P.  If cometary nuclei are confirmed to commonly have low thermal inertias and near-unity beaming parameters, this will greatly facilitate the interpretation of ground-based mid-IR radiometry. For faint comets often one can detect the comet at only one mid-IR wavelength. This is insufficient to constrain $\\eta$ and so in such a case one needs to assume a value for it. If we know {\\sl a priori} that there is a strong likelihood of $\\eta=1$, then this problem is removed.  Furthermore one would not necessarily need to obtain time-consuming photometry at many extra wavelengths. A survey of the thermal emission from cometary nuclei would be logistically easier and would yield significant physical information with relatively few model-dependent problems.  We note that our solution for $\\eta$ is fairly low given the phase angle of the observations.  \\cite{del03} found an approximately linear relationship between the beaming parameters and phase angles of several near-Earth asteroids (NEAs). Their best-fit line would suggest that $\\eta$ at $46^\\circ$ should be near 1.27. However there is much scatter in their trend and our solution here for 162P is not significantly off. In any case, cometary nuclei likely have different thermal properties anyway.  The fact that we do not see an elevated $\\eta$ at such phase angles, unlike what has been seen for many NEAs, is suggestive of this.  \\subsection{Radius and Albedo Context}  A large database of physical information on cometary nuclei was compiled by \\cite{lam05}.  By comparing with this work, our radius for 162P makes it  one of the largest Jupiter-family comets (JFCs). Only 28P/Neujmin 1 has an accepted radius that is larger, and 10P/Tempel 2 and 143P/Kowal-Mrkos have comparable radii. This is significant since it was only discovered in October 2004, despite its large size, so perhaps the sample of known JFCs is not as complete as previously thought. Studies of the JFC size distribution by \\cite{jf99}, \\cite{wl03}, \\cite{mhm04}, and \\cite{lam05} suggested that we knew about all the nuclei larger than a few kilometers, or at least those with perihelia within roughly 2 AU of the Sun.  162P's perihelion is at 1.2 AU, and its current orbit-intersection distance with Earth is 0.23 AU, so it is a fairly close yet large comet that avoided discovery.  While this is partly due to poor apparitions, it also suggests that the comet's weak level of cometary activity -- as shown in late 2004 and early 2005 -- kept it below the threshold brightness for discovery. The question of our completeness of the JFC sample perhaps needs to be revisited.  The albedo of the nucleus appears to be typical. While the error bar on the albedo is approximately 40\\%, it overlaps with a significant fraction of the known albedos among JFCs \\citep{lam05}. However, it is important to point out that there are only nine other albedos known for JFCs, and that we do not yet have a statistically robust sample with which to make a mathematical comparison."
    },
    "0608/astro-ph0608664_arXiv.txt": {
        "abstract": "{\\\\  {\\it Aims:} In this work the luminosity function of QSOs is measured in the redshift range $3.5<z<5.2$ for the absolute magnitude interval $-21<M_{145}<-28$. Determining the faint-end of the luminosity function at these redshifts provides important constraints on models of the joint evolution of galaxies and AGNs.  {\\it Methods:} We have defined suitable criteria to select faint QSOs in the GOODS fields, checking their effectiveness and completeness in detail. A spectroscopic follow-up of the resulting QSO candidates was carried out.  The confirmed sample of faint QSOs is compared with a brighter one derived from the SDSS.  We used a Monte-Carlo technique to estimate the properties of the luminosity function, checking various parameterizations for its shape and evolution.  {\\it Results:} Models based on pure density evolution show better agreement with observation than do models based on pure luminosity evolution. However, a different break magnitude with respect to $z\\sim 2.1$ is required at $3.5<z<5.2$.  Models with a steeper faint-end score a higher probability.  We do not find any evidence for a bright-end flattening at redshift $z>3.5$. ",
        "introduction": "Once considered peculiar and exotic objects, quasars (QSOs) have been recognized in recent years as an important and possibly a necessary phase of galactic evolution, at least for spheroids (Danese et al. 2003). In particular, models of a joint evolution of galaxies and QSOs \\citep[e.g.][]{granato01,granato04,hopkins05,hopkins06} reproduce the properties of present-day elliptical galaxies assuming that their vigorous star formation at high redshift is quenched by the feedback (i.e.  galactic winds). This process is triggered by the QSO activity fed by the accretion of matter onto a supermassive black hole (SMBH) at the center of the galaxy itself (see also Monaco \\& Fontanot 2005). In this way, the properties of high redshift QSOs and, in particular, their luminosity function (LF) are fundamental in understanding the phenomena driving galaxy formation. It has been hypothesized that the feedback from AGN inverts the dark matter halo (DMH) hierarchical sequence for the collapse of baryonic matter, a phenomenon also known by the term {\\it downsizing} \\citep[e.g.][]{granato04, croton06, bow05, paper2, menci06, scaoh}. The depth of the potential wells in different DMH is a key factor in determining the timescale for star formation, accretion, and eventually for the effects of feedback, setting the pace for AGN activity throughout cosmic history.  Great progress has recently been achieved thanks to the large amount of data coming from major observational programs such as the {\\it Two Degree Field QSO Redshift Survey} (2QZ, Croom et al. 2004, www.2dFquasar.org) and the {\\it Sloan Digital Sky Survey}, with the third edition of the Quasar Catalog (DR3QSO, Schneider et al. 2005). The 2dF-QRS contains more than 23000 QSOs typically at redshift $z \\le 2.1$, while the DR3QSO lists 46420 QSOs up to $z \\sim 5.4$.  However, at high redshift the SDSS is sensitive only to the most luminous QSOs ($M_{145} \\la -26.5$\\footnote{Absolute AB magnitude at a wavelength of $145$ nanometers}). The faint-end of the high-z QSO LF, which plays a key role in comparing different predictions of the formation and evolution of galaxies, is left almost unconstrained due to the challenging magnitude limit.  The recent {\\it Great Observatories Origins Deep Survey} (GOODS) provides an important database for studying QSOs with $M_{145} \\la -21$, thanks to the depth of its optical observations and to its multi-wavelength nature. GOODS is a deep survey, not a wide one, but it is much larger than most previous deep HST/WFPC2 surveys, covering $320 \\, arcmin^2$, 32 times the combined solid angles of the Hubble Deep Field-North and South, and four times larger than their combined flanking fields. The GOODS-ACS optical data were analyzed in a previous paper (Cristiani et al. 2004 hereafter paper I) in order to select reliable QSO candidates and make a first estimate of their space density in the redshift interval $3.5 < z < 5.2$. The selection was carried out by defining suitable optical selection criteria based on magnitude limits and color criteria, and then by matching the optical candidates with Chandra X-ray surveys (Alexander et al. 2003; Giacconi et al. 2002).  In this paper we take advantage of the SDSS $3^{rd}$ Data Release, of a new analysis and of a more complete spectroscopic follow up of the GOODS data in order to determine the QSO space density at $3.5 < z < 5.2$, down to the faint-end of the LF ($M_{145} < -21$).  As in paper I we adopt a definition of QSOs encompassing all objects with strong, high-ionization emission lines and $M_{145} \\leq -21$, including conventional, broad-lined (type-1) QSOs, and narrow-lined, obscured (type-2) QSOs. Throughout the paper we adopt a flat universe with cosmological parameters $\\Omega_m = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and Hubble constant $H_0 =70 Km \\, s^{-1} Mpc^{-1}$. Magnitudes are in the AB system.  In Sect.~2 we present the databases, the selection of the QSO candidates, and the spectroscopic follow up; in Sect.~3 we present our algorithm for the determination of the LF; in Sect.~4 we discuss the results in terms of implications for current models of galaxy formation. ",
        "conclusions": "In this paper we used the SDSS and GOODS databases to build an LF in the magnitude interval $-21.0 < M_{145} < -28.5$ and in the redshift interval $3.5 < z < 5.2$. In the first part of the paper we completed the selection of the QSOs candidates in the GOODS fields carried out by paper I.  We repeated the original analysis but using an improved version of the optical catalogues, and discuss the results of the spectroscopic follow-up of targets. The main conclusions of paper I on the relatively low surface density of faint high redshift QSOs are confirmed and strengthened.  We also explored the possible presence of ``X-ray faint'' AGNs, not detected in the A03 and G02 catalogues. This additional search was based on a morphological analysis carried out with the ACS optical images, selecting point-like candidates via a direct comparison with the expected QSO colors. We followed up the seven most promising candidates spectroscopically finding no additional QSO among them and no evidence of any significant population of faint X-ray high-z AGN with magnitude $M_{145} < -21$ in our fields.  We then combined the QSO sample in the GOODS fields with the DR3QSO to study the evolution of the QSO LF.  In order to understand the systematics in the different samples, we exploited a method for simulating the combined survey. This method is based on a library of QSO template spectra, built up starting from observed SDSS QSO spectra in the redshift interval $2<z<2.25$. Using this library we were able to predict with great reliability the color distributions of observed high-z QSO in the SDSS and GOODS photometric systems.  Assuming a parameterization of the LF and its redshift evolution, we were then able to simulate the expected object distribution in redshift and apparent magnitude. From the statistical comparison between real and simulated samples, we were able to quantitatively evaluate the agreement of the assumed LF with the data.  We tried different LF models and our results point out that PDE models show better agreement with observations than do PLE models, even if a luminosity evolution of the magnitude of the break, $M^{*}_{145}$, from $z=2.1$ to $z=3.5$ is required to obtain a good match.  The most interesting result in terms of physical models is related to the faint end slope of the LF. Parameterizations assuming a relatively steep faint-end slope similar to Ri5 score a higher probability with respect to those with a faint-end slope as flat as in Cr4.  If we also tried fitting also the faint-end slope, we obtain a best-fit value even higher than the Ri5 estimate.  A similar conclusion has been reached by Hao et al.  (2005). Finally, we do not find any evidence of a flattening at high-z of the bright-end of the LF with respect to the Cr4 value, at variance with the results of Fan et al. (2003).  We used our best estimate for the LF to compute the QSO contribution to the UV background at redshift $3.5<z<5.2$, and concluded that the QSO contribution is insufficient for ionizing the IGM at these redshifts.  When compared to physically motivated models, the present results indicate a suppression of the formation or feeding of low-mass SMBHs inside DMHs at high redshift.  Several processes can lead to this phenomenon: photoionization heating of the gas by UV background (Haiman, Madau \\& Loeb 1999), feedback from star formation (Granato et al. 2004), or the interaction between the QSO and the host galaxy. In a recent paper, Monaco \\& Fontanot (2006) studied this last possibility in detail, providing hints of a new mechanism for triggering galactic winds. In a forthcoming paper (Fontanot et al. 2006), this feedback mechanism will be inserted into a self-consistent model for galaxy and AGN formation (Monaco, Fontanot \\& Taffoni 2006): the present results for the high-z LF will be used to constrain the model and estimate the relevance of the joint feedback from supernovae and AGN in shaping the redshift evolution of AGNs."
    },
    "0608/astro-ph0608452_arXiv.txt": {
        "abstract": "{So far, red giant oscillations have been studied from radial velocity and/or light curve variations, which reveal frequencies of the oscillation modes. To characterise radial and non-radial oscillations, line profile variations are a valuable diagnostic. Here we present for the first time a line profile analysis of pulsating red giants, taking into account the small line profile variations and the predicted short damping and re-excitation times. We do so by modelling the time variations in the cross correlation profiles in terms of oscillation theory.} {The performance of existing diagnostics for mode identification is investigated for known oscillating giants which have very small line profile variations. We modify these diagnostics, perform simulations, and characterise the radial and non-radial modes detected in the cross correlation profiles. } {Moments and line bisectors are computed and analysed for four giants. The robustness of the discriminant of the moments against small oscillations with finite lifetimes is investigated. In addition, line profiles are simulated with short damping and re-excitation times and their line shapes are compared with the observations.} {For three stars, we find evidence for the presence of non-radial pulsation modes, while for \\object{$\\xi$ Hydrae} perhaps only radial modes are present. Furthermore the line bisectors are not able to distinguish between different pulsation modes and are an insufficient diagnostic to discriminate small line profile variations due to oscillations from exoplanet motion.} {} ",
        "introduction": "Techniques to perform very accurate radial velocity observations are refined during the last decade. Observations with an accuracy of a few m\\,s$^{-1}$ are obtained regularly (see e.g. \\citet{marbut}, \\citet{queloz2001a}) and detections of amplitudes of 1~m\\,s$^{-1}$ are  nowadays possible with e.g. HARPS \\citep{pepe2003}.  This refinement not only forced a breakthrough in the detection of extra solar planets but also in the observation of solar like oscillations in distant stars. Solar like oscillations are excited by turbulent convection near the surface of cool stars of spectral type F, G, K or M and show radial velocity variations with amplitudes of typically a few cm\\,s$^{-1}$ to a few m\\,s$^{-1}$, and with periods ranging from a few minutes for main sequence stars to about half an hour for subgiants and a couple of hours for giants.  For several stars on or close to the main sequence, solar like oscillations were detected for a decade (see e.g. \\citet{Kjeldsen1995}, \\citet{bouchy2001b}, \\citet{bouchy2003}, \\citet{bedding2003}). More recently, such type of oscillations has also been firmly established in several red (sub)giant stars.  \\citet{frandsen2002}, \\citet{deridder2006b}, \\citet{carrier2006a}, (see also \\citet{barban2004}) and \\citet{carrier2003} used the CORALIE and ELODIE spectrographs to obtain long term high resolution time series of the three red giants \\object{$\\xi$ Hydrae}, \\object{$\\epsilon$ Ophiuchi} and \\object{$\\eta$ Serpentis} and of the subgiant \\object{$\\delta$ Eridani}, respectively. They unravelled a large frequency separation in the radial velocity Fourier transform, with a typical value expected for solar like oscillations in the respective type of star (a few $\\mu$Hz for giants and a few tens $\\mu$Hz for subgiants), according to theoretical predictions (e.g. \\citet{dziembowski2001}). It was already predicted by \\citet{dziembowski1971} that non-radial oscillations are highly damped in evolved stars, and that, most likely, any detectable oscillations will be radial modes.  So far, red (sub)giant oscillations have only been studied from radial velocity or light variations. Line profile variations are a very valuable diagnostic to detect both radial and non-radial heat driven coherent oscillations (e.g. \\citet{aerts2003} and references therein), and to characterise the wavenumbers ($\\ell,m$) of such self-excited oscillations. It is therefore worthwhile to try and detect them for red (sub)giants with confirmed oscillations and, if successful, to use them for empirical mode identification. This would provide an independent test for the theoretical modelling of the frequency spectrum.  It is also interesting to compare the line profile diagnostics used for stellar oscillation analysis with those usually adopted to discriminate oscillations from exoplanet signatures, such as the line bisector and its derived quantities. Recently, \\citet{gray2005} pointed out that the wide range of bisector shapes he found must contain information about the velocity fields in the atmospheres of cool stars, but that the extraction of information about the velocity variations requires detailed modelling. Here we perform such modelling in terms of non-radial oscillation theory. We do so by considering different types of line characteristics derived from cross correlation profiles. \\citet{dall2006} already concluded that bisectors are not suitable to analyse solar like oscillations. We confirm this finding and propose much more suitable diagnostics. We nevertheless investigate how line bisector quantities behave for confirmed oscillators, in order to help future planet hunters in discriminating the cause of small line profile variations in their data.  The main problems in characterising the oscillation modes of red (sub)giants are, first, the low amplitudes of the velocity variations, which results in very small changes in the line profile. Secondly, the damping and re-excitation times are predicted to be very short. Indeed, \\citet{stello2004} derived an oscillation mode lifetime in \\object{$\\xi$ Hydrae} of only approximately two days. In general, mode lifetimes are difficult to compute for different stellar parameters from current available theory, cf. \\citet{houdek1999} versus \\citet{stello2004}. \\citet{stello2006} derived the observed mode lifetime of \\object{$\\xi$ Hydrae} from extensive simulations and find a large difference with theoretical predictions. The known mode identification methods from line profile variations all use an infinite mode lifetime so far. Here, we provide line profile variations simulated for stochastically excited modes, and investigate how the damping affects the behaviour of the diagnostics in this case. We apply our methodology to the four case studies of \\object{$\\xi$ Hydrae} (HD100407, G7III), \\object{$\\epsilon$ Ophiuchi} (HD146791, G9.5III), \\object{$\\eta$ Serpentis} (HD168723, K0III) and \\object{$\\delta$ Eridani} (HD23249, K0IV). Some basic properties of these stars are listed in Table~\\ref{propstar}.  \\begin{table*} \\begin{minipage}{17cm} \\caption{Basic stellar parameters of the four stars: Effective temperature ($T_{\\mathrm{eff}}$) in Kelvin, rotational velocity ($\\upsilon \\mathrm{sin}i$) in km\\,s$^{-1}$, parallax ($\\pi$) in mas, distance ($d$) in pc, the apparent magnitude ($m_{v}$) and absolute magnitude ($M_{V}$) in the V band.} \\label{propstar} \\centering \\renewcommand{\\footnoterule}{} \\begin{tabular}{lcccc} \\hline\\hline parameter & \\object{$\\epsilon$ Ophiuchi}\\footnote{\\citet{deridder2006b}, $^{b}$ \\citet{taylor1999}, $^{c}$ \\citet{carrier2003}, $^{d}$ \\citet{glebocki2000}, $^{e}$ \\citet{esa1997}} & \\object{$\\eta$ Serpentis} & \\object{$\\xi$ Hydrae} & \\object{$\\delta$ Eridani}\\\\ \\hline $T_{\\mathrm{eff}}$ [K] & $4887 \\pm 100$ & $4855 \\pm 19^{b}$ & $5010 \\pm 15^{b}$ & $5050\\pm100^{c}$\\\\ $\\upsilon \\mathrm{sin}i$ [km\\,s$^{-1}$] & $3.4 \\pm 0.5$ & $2.6\\pm 0.8^{d}$ & $2.4^{d}$ & $2.2\\pm0.9^{d}$\\\\ $\\pi$ [mas] & $30.34\\pm0.79$ & $52.81\\pm0.75^{e}$ & $25.23\\pm0.83^{e}$ & $110.58\\pm0.88^{e}$\\\\ $d$ [pc] & $33.0\\pm 0.9$ & $18.9 \\pm 0.3$ & $39.6\\pm1.5$ & $9.0\\pm0.1$\\\\ $m_{v}$ [mag] & $3.24\\pm0.02$ & $3.23\\pm0.02^{e}$ & $3.54\\pm0.06^{e}$ & $3.52\\pm0.02^{e}$\\\\ $M_{V}$ [mag] & $0.65\\pm0.06$ & $1.85\\pm0.05$ & $ 0.55\\pm0.04$ & $3.75\\pm0.2$\\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*}  The paper is organised as follows. In Sect.~2 the data set at our disposal and different observational diagnostics are described. In Sect.~3 the diagnostics for mode identification are described, while in Sect.~4 the damping and re-excitation effects in theoretically generated spectral lines are investigated. In Sect.~5 the results obtained from the simulations and observations are compared and in Sect.~6 some conclusions are drawn. ",
        "conclusions": "We have implemented a line profile generation code for damped and re-excited solar like oscillations. We found that the quantities $\\langle \\mathrm{v} \\rangle$, $\\langle \\mathrm{v}^{3} \\rangle$ and amplitude across the profile are good line diagnostics to characterise these oscillations. Our simulations were made for a single lineforming region while we compared them with cross correlation profiles based on a mask. For this reason, the phase variations across the profile are difficult to interpret, since they are an average over the depth in the atmosphere, and this average depends on the shape of the eigenfunctions. Moreover, we learned from our simulations that the phase behaviour can be quite different for realisations, as illustrated in Fig.~\\ref{amplphaseveleta2}. More extensive simulations across the depth would be needed to interpret the observed phase behaviour in detail. In order to do so, we would need to know the shape of the excited eigenmodes as a function of depth in the outer regions of the star.  From the frequency analysis performed on the bisector velocity span, it becomes clear that it is not a good diagnostic to unravel solar like oscillations.  Individual mode identification is not yet possible with the method used in the present work. This is not surprising because we did a first exploration of how well-working methods for modes with infinite lifetimes could be adapted for damped oscillations. The discriminant of the moments is not useful at the low amplitudes present in the sub(giants) investigated in the present work. Nevertheless, from the resemblance between the observed amplitude distribution with the simulated ones, and comparison between different modes of the same star, we presented clear observational evidence for modes with different m-values in the same star for three of our four examples, i.e. non-radial modes are present in red (sub)giants. The amplitudes for \\object{$\\xi$ Hydrae} are too low to make firm conclusions, and since we see only one frequency in the line profile variations, a comparison of the amplitude shapes for different modes is not possible. Therefore we can not exclude that this star has only radial modes as proposed by \\citet{frandsen2002}.  \\citet{dziembowski2001} and \\citet{houdek2002} show theoretically that stochastically excited radial pulsation modes are most likely to be observed in red giants. They predict that, for non-radial modes, damping effects will have influence in the p-mode cavity as well as in the g-mode cavity, while for radial modes only damping in the p-mode cavity is present. Therefore, according to this theory, radial modes are less severely influenced by damping than non-radial modes. Our result pointing towards the presence of non-radial modes, calls for a re-evaluation of the theoretical predictions.  \\begin{acknowledgement} SH wants to thank the staff at the Instituut voor Sterrenkunde at the Katholieke Universiteit Leuven for their hospitality during her three month visit. JDR is a postdoctoral fellow of the fund for Scientific Research, Flanders. FC was supported financially by the Swiss National Science Foundation. The authors are supported by the Fund for Scientific Research of Flanders (FWO) under grant G.0332.06 and by the Research Council of the University of Leuven under grant GOA/2003/04. \\end{acknowledgement}"
    },
    "0608/hep-ph0608092_arXiv.txt": {
        "abstract": "I discuss a theory of non-solar cosmic rays (CRs) based on a single type of CR source at all energies. All observed properties of CRs are predicted in terms of very simple and completely `standard' physics.  The source of CRs is extremely `economical': it has only one parameter to be fitted to the enormous ensemble of all of the data. All other inputs are `priors', that is theoretical or observational items of information independent of the properties of the source of CRs and chosen to lie in their pre-established ranges. ",
        "introduction": "A couple of years ago, I presented at a large cosmic-ray (CR) conference a preliminary version of the theory \\cite{Florence} I am to summarize here. The coffee break subsequent to my talk was attended by hundreds of participants. A prominent CR theorist, who had heard my talk, told me: {\\it `Your theory is unacceptable. It solves the whole problem'}. He paused to make a large hand gesture embracing the surrounding crowd, and added: {\\it `What are then all these people going to do?} Was he just joking? Not that I am convinced that we have {\\it the complete solution} to the problem. But, I shall argue, we have {\\it a} solution that Ockham would have favoured, for the reasons stated in the abstract.  The {\\it standard} theory posits that, up to the knee, CRs are produced by supernovae (SNe) \\cite{BZ}, via shocks in the interaction of their roughly spherical non-relativistic ejecta with the interstellar medium (ISM) \\cite{Shlovsky}. The NASA website lists almost 70000 refereed papers on `cosmic ray(s)', most of them theoretical and `standard'. As implied by a fraction of these papers, the standard theory has problems: it does not accelerate CRs up to the knee (e.g.~\\cite{Cesarsky}), the fraction of SN remnants perhaps compatible with the acceleration of nuclear CRs is insufficient to explain the CR luminosity of the Galaxy (e.g.~\\cite{HESS}), CRs produced by SN remnants (most of which are within the `solar circle') would diffuse outwards, generating a directional asymmetry which is not observed (e.g.~\\cite{asymmetry}). CRs above the knee are mysterious.  Our theory \\cite{DD2006} is {\\it non-standard}  in that CRs are accelerated {\\it at all energies} by the {\\it relativistic jets} ejected by  SNe. It is part of a `unified theory of high-energy astrophysics' \\cite{Florence}, based  on the `cannonball' model \\cite{GRB1} of the jets of accreting black holes and neutron stars, which also explains simply the properties of gamma-ray bursts (GRBs) \\cite{DD}, X-ray flashes  (XRFs) \\cite{DDDXRF}, their respective afterglows (AGs) \\cite{AGoptical,AGradio}, the gamma `background' radiation \\cite{DDGBR}, the CR luminosity of our Galaxy \\cite{DD2006,DDLum}, and the properties of galaxy clusters harbouring `cooling flows'  \\cite{CDD}. Many more and less self-referring citations are given in \\cite{DD2006,DD}. ",
        "conclusions": "Our theory is incomplete in several respects. The ejection of CBs in episodes of accretion onto a compact object is supported by observations, but not fundamentally understood (this is also the case for the ejection of SN shells). Because CBs deposit CRs along their kpc-long trajectories, CR diffusion does not play a crucial role, but we have not studied it in detail (in the standard theory diffusion results in directional asymmetries that are not observed, and calls for ad-hoc remedies, such as `CR reacceleration'). The temporary confinement of CRs ---in the Galaxy or within a CB--- is not fully understood. Neither is the `penetrability' of the Galaxy to extragalactic CRs of energy below the ankle. We have studied the dynamics of the expansion of CBs, but not modelled it in minute detail. We contend that a good fraction of the original energy of CRs ends up in the production of  `equipartition' MFs~\\cite{DDMF}, but we cannot predict this, nor determine the effect that it may have on the CR spectral shape.  In spite of the above limitations, we have demonstrated how our simple {\\it and single} accelerators ---cannonballs--- are effective at all observed energies. The mechanisms of CR acceleration, particularly the `elastic' one, could hardly be simpler: a fast massive object slows down by kicking out of its way the light particles it encounters.  Most of our results are `robust' in that ---within very large brackets--- they do not depend on the specific choices of parameters and priors:\\\\ $\\bullet$ An all-particle piecewise power-law spectrum with four features: two steepenings at the knee and the second knee, a softening at the ankle, and an end-point at the roughly-coincident GZK and proton-acceleration cutoffs.\\\\ $\\bullet$ An UHECR flux above the ankle, which is predicted ---to within a factor of a few--- and otherwise parameter-free.\\\\ $\\bullet$ A composition dependence at 1 TeV with the observed trend, so different from that of the ISM.\\\\ $\\bullet$ A very low-energy flux whose spectral shape is independent of any CB-model `prior' parameters.\\\\ $\\bullet$ Individual-element knees that scale like $A$ and occur at the predicted energies.\\\\ $\\bullet$ A non-trivial shape of the individual knees: an abrupt decrease in flux, followed by a spectrum steeper than that below the knee.\\\\ $\\bullet$ An ankle with the observed shape. The dominantly Galactic-Fe flux below it and the dominantly extragalactic-proton flux above it are comparable in magnitude at the estimated escape `ankle' energy \\cite{Cocconi} of Galactic protons.\\\\ $\\bullet$ A composition dependence that is almost energy-independent below the knee becomes `heavier' from the knee to the second knee, `lighter' again above it, and finally heavier at yet unmeasured ultra-high energies.\\\\ $\\bullet$ An `extended' distribution of CR sources along CB trajectories emerging from the central realms of the Galaxy, where most SN explosions occur, implying a CR flux at the Earth's location with a much smaller and less energy-dependent anisotropy than that of standard models of CRs.\\\\ $\\bullet$ Predictions for the values of a related set of observables: the CR luminosity, confinement time and volume of the Galaxy, the spectral indices of CR electrons and of the diffuse GBR. \\\\  Our results describe the observed properties of hadronic non-solar CRs very well from the lowest energies  to $\\sim 10^{10}$ GeV. Above that energy and up to the highest observed energies, $\\sim 10^{11}$ GeV, our theory opts for the data gathered with fluorescence detectors. Overall, the energy range for which the theory is successful covers ten decades and the flux extends over  three times as many."
    },
    "0608/astro-ph0608334_arXiv.txt": {
        "abstract": "We report long-slit spectroscopic observations of the Herbig Ae/Be star Elias~1 in the 3.2 - 3.6$\\,\\umu$m region covering the C-H stretch emission features of hydrogen-terminated diamonds and PAHs. The data were recorded at UKIRT using UIST and yield information on the profiles and intensities of the bands as a function of offset along the N-S and E-W axes centred on the close binary. The diamond and nearby IR continuum emission arises from a symmetrical inner core region ($\\leq$\\,0.34$^{\\prime\\prime}$ or 48 AU).  The 3.3$\\,\\umu$m PAH emission is extended along the E-W axis up to \\emph{c.} 100 AU each side of the star.  This result supports a suggestion of \\citet{haa97} of an E-W oriented bipolar nebula in Elias 1. ",
        "introduction": "The morphology, evolution and chemical composition of circumstellar disks and halos of pre-main-sequence intermediate-mass Herbig Ae/Be stars is of key importance in the study of star and planet formation \\citep{wat98}.  Emission features of silicates and PAHs are commonly though not always observed \\citep{ack04} and in a few cases emission from surface-hydrogen-terminated diamonds is also detected. This raises interesting and challenging questions as to how the diamonds form, why their emission is seen in only a few stars, and whether there is a link between circumstellar diamonds and the presence of nanodiamonds in meteorites.  The existence of organic PAHs in the nebular environment also provides astrobiological interest. In this paper we describe near-IR long-slit observations of one of the rare diamond C-H emitting objects, Elias~1. We focus on the spectroscopy and spatial distribution of diamond (C-H), PAH and continuum emission from the circumstellar environment, and make comparison with the related object HD~97048.  The pre-main-sequence star Elias~1 (V892 Tau, Elias~3-1) is a close binary with a separation of $\\sim$50 mas at PA $\\sim$50$^\\circ$ \\citep{smi05}. Its spectral type is uncertain ranging through A0 \\citep{eli78}, A6 with $A$$_V$~=~3.9 mag \\citep{zin94}, B9 with $A$$_V$~=~8.85 mag \\citep{str94}, with an $A_V$ value of 10.5 also reported \\citep{tei99}. Elias~1 has a very flat, possibly rising, SED in the near-IR to far-IR which is markedly different from that of HD~97048 \\citep{hil92}, the two stars being classified as group II and I, respectively.  The SED for Elias~1 is taken by \\citet{smi05} to suggest the possible existence of an envelope in addition to a disk, or a flared disk. Other Elias~1 characteristics include X-ray \\citep{zin94,gia04,ham05,ste06} and radio \\citep{ski93} emission.  An emission feature in the \\emph{c.} 3.5$\\,\\umu$m region for Elias~1 was first reported in 1982 by \\citet{all82} and was subsequently shown to have two components near 3.41 and 3.52$\\,\\umu$m with an additional PAH feature at 3.29$\\,\\umu$m \\citep{whi83}.  The only previously detected example was HD~97048 \\citep{bla80}.  Subsequent spectroscopic observations of Elias~1 have been made \\citep{whi84,tok91,geb97}, including ISO (SWS) data \\citep{vank02}. The features at 3.43 and 3.53$\\,\\umu$m remained unassigned until \\citet{gui99} found a convincing correspondence between the astrophysical bands of Elias~1 and HD~97048 and laboratory absorption spectra of the C-H stretching modes on hydrogen-terminated diamond nanocrystal films at $\\sim$~1000~K \\citep{cha95}. ISO spectra of Elias~1 covering the range 2.5~-~13.5 $\\,\\umu$m \\citep{vank02} contain absorption features due to water and CO$_2$ ice, and emission from PAHs, diamonds and silicates, the last of these being absent from the spectrum of HD~97048. In this paper we have chosen not to use the word `nanodiamond' in discussing the stellar emission features as their size is probably at least 50~nm \\citep{che02,she02,jon04} in contrast to meteoritic nanodiamonds which are typically only 2-3~nm across \\citep{lew87}. We also remark that the diamonds observed are exclusively those that are H-terminated at the surface; absence of hydrogen or its substitution by other chemical groups on the surface in general give spectra outside the region studied. The only other Herbig Ae/Be objects for which diamond emission is reported are BD$+$40$^\\circ$4124 \\citep{ack04} and MWC 297 \\citep{ter01,ack04}; emission features attributed to diamond emission are also seen in the post-AGB star HR~4049 \\citep{geb89} but are relatively weak. ",
        "conclusions": "We have conducted a long-slit study of Elias~1 in the 3.2 - 3.6$\\,\\umu$m region which has revealed a symmetrical distribution for the continuum (3.6$\\,\\umu$m) and diamond C-H emission features, and more extended PAH 3.3$\\,\\umu$m emission which is oriented along the E-W axis.  Given that for many objects the 3.3$\\,\\umu$m PAH emission is significantly less extended than for the UIR bands at longer wavelengths, it will be of interest to determine the mid-IR PAH emission distribution using long-slit spectroscopy in order to probe the size and geometry of the bipolar nebula."
    },
    "0608/astro-ph0608102_arXiv.txt": {
        "abstract": "We present an investigation of the properties and environments of bright extremely red objects (EROs) found in the fields of the quasars \\txs\\ and \\fourc, both at $z\\sim1.4$. There is marginal evidence from {\\it Chandra} ACIS imaging for hot cluster gas with a luminosity of a few $10^{44}$ ergs s$^{-1}$ in the field of \\fourc. The \\txs\\ field has an upper limit at a similar value, but it also clearly shows an overdensity of faint galaxies. None of the EROs are detected as X-ray sources.  For two of the EROs that have spectral-energy distributions and rest-frame near-UV spectra that show that they are strongly dominated by old stellar populations, we determine radial-surface-brightness profiles from adaptive-optics images.  Both of these galaxies are best fit by profiles close to exponentials, plus a compact nucleus comprising $\\sim30$\\% of the total light in one case and 8\\% in the other.  Neither is well fit by an $r^{1/4}$-law profile.  This apparent evidence for the formation of massive $\\sim2\\times10^{11}$ disks of old stars in the early universe indicates that at least some galaxies formed essentially monolithically, with high star-formation rates sustained over a few $10^8$ years, and without the aid of major mergers. ",
        "introduction": "Galaxies classified as {\\it extremely red objects} (EROs; \\eg\\ \\citealt{dad00,cha00,sto01,cim02,sar04,yan04, dad04}; earlier references can be found in these papers) are usually defined to have $R\\!-\\!K$ greater than a value somewhere in a range of 5--6. They have long been known to include at least two disparate classes.  Some are highly reddened high-redshift star-forming galaxies (\\eg\\ \\citealt{dey99,cim02,yan04b}).  The remainder are almost exclusively galaxies with $z\\gtrsim1$ whose stellar populations are already ancient and have little admixture from more recent star formation \\citep[\\eg][]{sto95,dun96,spi97,soi99,liu00,sto01,sto04,che04,cim04,fu05}. For convenience, we shall henceforth simply refer to this latter class as ``old galaxies'' (OGs; however, note that this term refers to the age of the stellar population, not necessarily the age of assembly of the galaxy itself). These OGs represent the earliest major episodes of star formation in the universe, and they seem to have formed the overwhelming majority of their stars over a very short period. Accordingly, they potentially offer us an avenue to investigate the formation and early evolution of the very first massive galaxies to appear in the universe.  There have been a number of discussions of the mix of morphologies in various samples of EROs.  It is commonly assumed that morphological types are surrogates for the two classes of EROs mentioned above; \\ie\\ that spheroidals as a group are more-or-less congruent with galaxies dominated by old stellar populations and that disk-like galaxies and those showing obvious merging activity are dusty star-forming galaxies \\citep{mor00,sti01,smi02,saw05}.  Nevertheless, both \\citet{gil03} and \\citet{yan04} mention the possibility of truly passive disk-dominated galaxies. \\citet{iye03} seem to have found such a galaxy at $z\\sim1.5$, and \\citet{sto04} and \\citet{fu05} have shown examples of galaxies that seem to be essentially pure disks of old stars at $z=2.48$ and $z=1.34$, respectively.  So far, the morphologies of only a few OGs at redshifts significantly greater than 1 have been studied in any detail (\\eg\\ \\citealt{bun03,sto04,yan04,cim04,fu05,dad05}). Here we present an investigation of the morphologies and environments of 3 OGs (and 1 dusty star-forming galaxy) found in the fields of the quasars TXS\\,0145+386 ($z=1.442$) and 4C\\,15.55 ($z=1.406$). We assume a flat cosmology with $H_0 = 70$ and $\\Omega_m=0.3$. ",
        "conclusions": "The clustering amplitude of luminous galaxies comprising old stellar populations at high redshifts \\citep{dad00,mcC01,bro05} is similar to that of early-type galaxies at the present epoch. It is also comparable to that of powerful radio sources at similarly high redshifts, so it would not be surprising if these two classes were to have a considerable overlap in their spatial distribution.  \\citet{hal01} and \\citet{wol03} have found an excess of EROs in radio-loud quasar fields that tends to corroborate such an overlap.  \\citet{wol03} suggested two possibilities for this excess: (1) a physical clustering of early-type galaxies around the quasars, and (2) massive clusters along the line of sight to the quasars, which bias the quasar catalogs via gravitational amplification.  Our spectroscopic confirmation that at least one of the OGs in each of our fields has a redshift agreeing with that of the quasar to within a few hundred \\kms\\ supports the first interpretation.  Our attempt to detect direct evidence for clusters in the two fields we discuss in this paper has produced ambiguous results.  While we clearly seem to find an enhanced surface density of faint galaxies in the field of TXS\\,0145+386, our Chandra ACIS observation of the same field gives only an upper limit to any diffuse X-ray emission from the cluster.  On the other hand, in the field of 4C\\,15.55, we appear to have a $\\sim2.4~\\sigma$ X-ray detection of diffuse emission, but we cannot be absolutely sure that this emission is not associated with the X-ray jet from the quasar, in spite of our effort to excise any such emission. Neither the upper limit in the first case, nor the possible detection in the second, are inconsistent with the presence of a fairly massive cluster surrounding the respective quasars.  Both of the OGs for which we have AO imaging appear to have compact nuclear components. After removal of these nuclear components, both of these galaxies are well fit by exponential profiles (and not by $r^{1/4}$-law profiles); however, the discrimination among profiles is much better for \\fourc\\ ER2 because of the much higher S/N of the Subaru image.  For this galaxy, there is no significant doubt that the stellar distribution has close to an exponential profile; even if one were to prefer the $n=2$ S\\'{e}rsic model (which eliminates the need to add a compact nucleus), it is only slightly different from an exponential beyond a radius of 0\\farcs2 ($=1.7$ kpc), and, to the extent it differs, it is a worse fit to the data.  It is true (as the referee has emphasized) that elliptical galaxies in the present-day universe can have a wide range of S\\'{e}rsic indices, including some as low as 2 \\citep{caon93,kho00}.  However, the same studies that have demonstrated this range have also shown that there is a strong positive correlation between the S\\'{e}rsic index and both the central surface brightness and the total luminosity of the galaxy. Ellipticals with $n\\sim2$ are, virtually without exception, galaxies with luminosities considerably below $L^*$ and with low central surface brightnesses (see, \\eg\\ Table 2 of \\citealt{caon93}). There are good reasons to believe that they have had different formation histories than have luminous ellipticals \\citep{kho00}, possibly involving secular evolution from disks. Such galaxies are in any case quite different from \\fourc\\ ER2, which is luminous ($\\sim2L^*$) and has a high surface brightness.  Some massive galaxies strongly dominated by old stellar populations at high redshifts clearly show closely $r^{1/4}$-law profiles \\citep{bun03,cim04,fu05,dad05}; others just as clearly show essentially pure exponential profiles \\citep{sto04,fu05,dad05}.  The only reasonably large structures in the present-day universe that show exponential profiles are the disk components of spiral and S0 galaxies and the ``pseudo-bulges'' \\citep[\\eg][]{kor04} found in late-type spirals.  The latter are not very massive, and they are produced over a significant fraction of a Hubble time via secular evolution from disks.  For these reasons, they are almost certainly not models for the exponential profiles we see in high-redshift OGs. This means that the exponential profiles in these massive high-redshift systems are due either to disks or to some configuration that is not found among massive galaxies at the present epoch. In at least one case \\citep{sto04}, the galaxy has a small axial ratio and looks very much like a highly inclined disk.  If these really are massive ($\\sim2\\times10^{11} M_{\\odot}$) disks of old stars, then they are significantly smaller than most massive disks at the present epoch. For example, if we take from the study of a complete sample of nearby disks by \\citet{vdK87} those with disk $M<-20.2$ (for our assumed $H_0=70$ km s$^{-1}$ Mpc$^{-1}$), the average exponential scale length is $5.5\\pm2.7$ kpc, which corresponds to an effective radius $r_e=9.2\\pm4.5$ kpc. In comparison, \\txs\\ ER1 and \\fourc\\ ER2 have $r_e = 1.69$ kpc and 3.0 kpc, respectively. Even the larger of these is barely more than half as large as the smallest of the 12 luminous galaxies from the \\citet{vdK87} sample. \\citet{dad05} have similarly found that early-type galaxies at high redshifts are significantly smaller than their counterparts in the present-day universe. Eight out of their 9 examples were found to have morphologies consistent with spheroids, and only 1 appeared to be a disk of old stars. But it is perhaps significant that this disk-like galaxy, at $z=2.47$, was by a fair margin the highest redshift galaxy in their sample.  \\citet{dad05} point out that this galaxy is quite similar to the galaxy discussed by \\citet{sto04}, which is also at a similar redshift.  Such massive disks would carry certain implications for the early formation of massive galaxies: \\begin{enumerate} \\item Such a disk must have formed quasi-monolithically and with strong dissipation. It is unlikely to have suffered any major mergers after a substantial fraction of the stars had been formed. \\item Depending on the age of the stellar population in the examples observed so far, quite large and sustained star formation rates are required to form most of the stellar mass in a short time. At a minimum, star formation rates of a few hundred $M_{\\odot}$ per year sustained for a few $10^8$ years are indicated.  Note that similar constraints are indicated by the $\\alpha$-element enhancement of the stellar content of massive ellipticals at the present epoch (\\eg\\ \\citealt{wor92, wor98}). \\item Since massive pure disks of old stars do not exist (or at least are extremely rare) at the present epoch, such galaxies must have a very high probability of being transformed into ellipticals via major mergers or into bulge-dominated S0s via less extreme encounters. This would not be surprising, given that our objects have been found in radio-source fields, which should be among the highest density regions in the early universe.  At high redshifts, most of these regions will not yet have become virialized, so velocity dispersions in sub-clumps may well be conducive to merging activity.  The fact that we suspect TXS\\,0145+386 ER2 to be an ongoing merger gives some observational credence to this line of reasoning.  This scenario might also explain why, among the admittedly small number of OGs that so far have published high-resolution morphological data, spheroids seem to be more common at $z\\sim1.5$ than at $z\\sim2.5$. \\end{enumerate}  All of this fits in rather well with what we know or suspect about massive cluster ellipticals at the present epoch.  Numerous lines of evidence \\citep[\\eg][]{pee02,tho02,tho05} point to early and rapid formation of the great majority of the stars in the most massive and luminous ellipticals, yet some simulations \\citep[\\eg][]{tho99,kau00} suggest late ($z$ between 1 and 2) assembly of the galaxies themselves.  The only way these two conditions can be reconciled is if the assembly is done from galaxies that formed most of their stars at much earlier times and that are largely free of gas.  Furthermore, the assembly process must be dominated by (at most) a few large such galaxies in order to preserve the well-known color-luminosity relation for ellipticals.  Thus, early forming massive disks in which star formation has been driven to completion or truncated by the expulsion of remaining gas could plausibly act as reservoirs for the stars that eventually end up in present-day elliptical or cD galaxies."
    },
    "0608/astro-ph0608428_arXiv.txt": {
        "abstract": "We construct a sample of extremely red objects (EROs) within the UKIDSS Ultra Deep Survey by combining the Early Data Release with optical data from the Subaru/\\textit{XMM-Newton\\/} Deep Field. We find a total of 3715 objects over 2013\\,arcmin$^2$ with $R-K>5.3$ and $K\\leq20.3$, which is a higher surface density than found by previous studies. This is partly due to our ability to use a small aperture in which to measure colours, but is also the result of a genuine overdensity of objects compared to other fields. We separate our sample into passively-evolving and dusty star-forming galaxies using their \\textit{RJK\\/} colours and investigate their radio properties using a deep radio map. The dusty population has a higher fraction of individually-detected radio sources and a higher mean radio flux density among the undetected objects, but the passive population has a higher fraction of bright radio sources, suggesting that AGNs are more prevalent among the passive ERO population. ",
        "introduction": "Extremely Red Objects (EROs) are variously defined on the basis of an optical--near-infrared colour, usually $R-K$ or $I-K$. These red colours can be due either to an evolved, passive stellar population at $z\\ga1$, or the presence of large quantities of dust associated with starburst activity (although edge-on discs may also contribute significantly; Yan \\& Thompson 2003). While the former population reveals the stellar mass already assembled in galaxies as a result of star formation at early cosmic times, the latter population reveals the star formation at $z\\sim1$ which, due to its dust-enshrouded nature, may be missed by other surveys. Since these two populations represent very different routes to the creation of present-day galaxies, there is a clear desire to reliably separate them and use them as tests of galaxy formation models.  Existing samples of EROs (for which we adopt the most common definition of $R-K>5.3$) have been relatively small (typically no more than about 100 objects), due to the difficulty in obtaining near-infrared imaging data with the necessary depth and areal coverage (e.g., Daddi et al.\\ 2000; Smail et al.\\ 2002; Moustakas et al.\\ 2004; Georgakakis et al.\\ 2006). The arrival of infrared cameras with large fields of view, such as WFCAM (Casali et al., in preparation), has largely eliminated this problem, allowing samples of EROs to be constructed which are large enough to permit further separation by magnitude, colour, and other properties. Furthermore, by locating infrared surveys in well-observed patches of sky, it becomes possible to study the multi-wavelength properties of EROs and learn more about their nature.  In this Letter, we combine near-infrared data from the UKIDSS Ultra Deep Survey Early Data Release (Dye et al.\\ 2006; Lawrence et al.\\ 2006) with optical data from the Subaru/\\textit{XMM-Newton\\/} Deep Field (SXDF; Furusawa et al., in preparation) to construct a sample of EROs to faint magnitudes ($K>20$) over an area larger than half a square degree. The Ultra Deep Survey will eventually cover 0.8\\,deg$^2$ to 5$\\sigma$ point source limits of $J=25.0$, $H=24.0$, and $K=23.0$ and has been sited in the SXDF because of the excellent deep multi-wavelength data in this region. We determine photometric redshifts for all objects and separate them into dusty star-forming and passively-evolving galaxies on the basis of their \\textit{RJK\\/} colours. We use deep radio data to determine the typical radio flux densities and star formation rates as a function of colour and magnitude. A more detailed analysis, including clustering and X-ray properties, will be undertaken after the UKIDSS DR1 release, which is expected to be a magnitude deeper than the present data, and will provide more reliable photometry. All magnitudes in this paper are on the Vega system. The SXDF catalogue is calibrated on the AB system, and we use the conversion $R_{\\rm Vega} = R_{\\rm AB} - 0.219$ (H.\\ Furusawa, private communication). We adopt a cosmology with $H_0 = 70\\rm\\,km\\,s^{-1}\\,Mpc^{-1}$ and $\\Omega_{\\rm m} \\equiv 1-\\Omega_\\Lambda = 0.3$. ",
        "conclusions": "We have described the construction of a sample of extremely red objects (EROs) within the Subaru/\\textit{XMM-Newton\\/} Deep Field (SXDF) and UKIDSS Ultra Deep Survey (UDS) which is much larger than previous samples. We have measured a higher surface density than found in other fields, but comparable to that measured in a small sub-region of our field, suggesting the SXDF samples a cosmic overdensity at $z\\sim1$. After using \\textit{RJK\\/} colours to separate the sample into passively-evolving (62\\% of EROs with $K\\leq19.5$) and dusty star-forming galaxies, we have used a deep radio map to study the radio properties of these two populations. Both populations are detected at high significance in a stacked radio image, with the dusty sample having a mean flux density (and derived mean star formation rate) equal to approximately twice that of the passive sample. The mean radio flux density of the passive sample is consistent with that of non-ERO galaxies with the same magnitudes."
    },
    "0608/astro-ph0608272_arXiv.txt": {
        "abstract": "Color confinement is only a supposition, which has not been proved in QCD yet. It is proposed here that macroscopic quark gluon plasma in astrophysics could hardly maintain colorless because of causality. The authors expected that the existence of chromatic strange quark stars as well as chromatic strangelets preserved from the QCD phase transition in the early universe could be unavoidable if their colorless correspondents do exist. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608044_arXiv.txt": {
        "abstract": "The recent discovery of two large trans-Neptunian objects (TNOs) 2003 UB$_{313}$ and 2005 FY$_9$, with surface properties similar to those of Pluto, provides an exciting new laboratory for the study of processes considered for Pluto and Triton: volatile mixing and transport; atmospheric freeze-out and escape, ice chemistry, and nitrogen phase transitions.\\\\ We studied the surface composition of TNO 2003 UB$_{313}$, the first known TNO larger than Pluto.\\\\ We report a visible spectrum covering the 0.35-0.95$\\mu$m  spectral range, obtained with the 4.2m William Herschel Telescope at ``El Roque de los Muchachos\" Observatory (La Palma, Spain). \\\\ The visible spectrum of this TNO presents very prominent absorptions bands formed in solid CH$_4$.  At wavelengths shorter than 0.6 $\\mu$m the spectrum is almost featureless and slightly red (S'=4\\%). The icy-CH$_4$ bands are significantly stronger than those of Pluto and slightly weaker than those observed in the spectrum of another giant TNO, 2005 FY$_9$, implying that methane is more abundant on its surface than in Pluto's and close to that of the surface of 2005 FY$_9$. A shift of 15 $\\pm$3 $\\AA$ relative to the position of the bands of the spectrum of laboratory CH$_4$ ice is observed in the bands at larger wavelengths (e.g. around 0.89 $\\mu$m), but not at shorter wavelengths (the band around 0.73 $\\mu$m is not shifted) this may be evidence for a vertical compositional gradient.  Purer methane could have condensed first while 2003 UB$_{313}$ moved towards aphelion during the last 200 years, and as the atmosphere gradually collapsed, the composition became more nitrogen-rich as the last, most volatile components condensed, and CH$_4$ diluted in N$_2$ is present in the  outer surface layers. ",
        "introduction": "The trans-Neptunian region is populated by icy bodies (TNOs), remnant planetesimals from the early solar system formation stages (Edgeworth \\cite{Edgeworth49}; Kuiper \\cite{Kuiper51}). TNOs are the source of the short period comets (Fern\\'andez \\cite{Fernandez80}) and are probably the most pristine objects in the Solar System. The low temperatures in this region ($\\sim$40K), implies that ices trapped at formation should be preserved and can provide key information on the composition and early conditions of the pre-solar nebula.  The recent discovery of three very bright TNOs, 2003 EL$_{61}$ (Ortiz et al.\\ \\cite{Ortiz2005}), 2003 UB$_{313}$ and 2005 FY$_9$ (Brown et al.\\ \\cite{Brownetal2005}, \\cite{Brownetal2005b}) provides an excellent opportunity to obtain spectra of TNOs with sufficient S/N to obtain reliable mineralogical information of their surfaces.   The spectra of three of the four largest members of the trans-neptunian belt, 2003 UB$_{313}$, Pluto, and  2005 FY$_9$  are dominated by strong methane ice absorption bands (Cruikshank et al.\\ \\cite{Cruik1993}, Brown et al.\\ \\cite{Brownetal2005}, Licandro et al.\\ \\cite{Licandroetal06}) Pluto also presents weak but unambiguous signatures of CO and N$_2$-ice (e.g. Cruikshank\\ \\cite{Cruik1998}). These bands are also detected in the spectrum of Neptune's satellite Triton (Cruikshank et al.\\ \\cite{Cruik1993}), a possibly captured ex-TNO. The presence of frozen methane on the surfaces of Pluto, Triton, 2005 UB${313}$ and 2005 FY$_9$ argues that the process suggested by Spencer et al.\\ (\\cite{Spencer1997}) in which surface methane is replenished from the interior, may be ubiquitous in large trans-neptunian objects.  2005 FY$_9$ and 2003 UB$_{313}$ provide an exciting new laboratory for the study of processes considered for Pluto and Triton: volatile mixing and transport; atmospheric freeze-out and escape, ice chemistry, and nitrogen phase transitions. In particular the abundance of volatiles like CO and N$_2$ is important to determine the possible presence of a bound atmosphere and constrain the formation conditions.    TNO 2005 UB$_{313}$ is the largest known object in the trans-neptunian belt, with a surface albedo higher than that of Pluto (2400+/-100 km or a size $\\sim$5\\% larger than Pluto, and p$_V$=86 $\\pm$7 \\%, Brown et al.\\ \\cite{Brownetal06}). Discovered near aphelion at 97.50 AU, it will take some 2 centuries to reach its perihelion at 38.2 AU. This huge variation in heliocentric distance causes large seasonal temperature variations that should affect the sublimation and recondensation of its surface volatiles.    In this paper we present visible spectroscopy of 2005 UB$_{313}$ and compare it with spectra of Pluto and 2005 FY$_9$ (Licandro et al.\\ \\cite{Licandroetal06}) in order to derive mineralogical information from its surface. ",
        "conclusions": "We present a new 0.35-9.4$\\mu$m spectrum of the TNO 2003 UB$_{313}$. The spectrum is very similar to that of Pluto, with prominent CH$_4$ ice absorptions bands. At wavelengths $<0.6 \\mu$m the spectrum is almost featureless and slightly red (S'=4$\\pm$1\\%/1000$\\AA$) supporting the existence of complex organics molecules (tholins) on its surface. The visible spectrum of 2003 UB$_{313}$ is not as red as spectra of Pluto and 2005 FY$_9$ (S'=8.8 and 8.9 \\%/1000$\\AA$ respectively), thus complex organics should be less abundant on the surface of 2003 UB$_{313}$ than on the surfaces of Pluto and 2005 FY$_9$.  The CH$_4$ ice bands in this new giant TNO are significantly stronger than those of Pluto, but weaker than those observed in the spectrum of 2005 FY$_9$ (Licandro et al.\\ \\cite{Licandroetal06}). Methane is more abundant and/or the methane ice grain particles (or the thickness of the surface ice layer) are larger on its  surface than on the surface of Pluto, and less abundant or composed of smaller grains than on the surface of 2005 FY$_9$.  A 15 $\\AA$ shift of the central wavelength of the 0.89$\\mu$m band relative to the pure methane band observed in the laboratory is observed. This shift is indicative of the presence of methane diluted in N$_2$. On the other hand, the 0.73 $\\mu$m band is not significantly shifted. This could be evidence for a vertical compositional gradient consistent with purer methane condensing first, with the composition becoming more nitrogen-rich as the last, most volatile components of the atmosphere condensed.  Such a compositional gradient could also arise via the solar gardening mechanism discussed by Grundy \\& Stansberry (\\cite{GruStan00}).    {\\bf Acknowledgements} W.M. Grundy gratefully acknowledges support from NASA Planetary Geology \\&\\ Geophysics grant NNG04G172G."
    },
    "0608/gr-qc0608003_arXiv.txt": {
        "abstract": "In this work, we explore the possibility of evolving $(2+1)$ and $(3+1)-$dimensional wormhole spacetimes, conformally related to the respective static geometries, within the context of nonlinear electrodynamics. For the $(3+1)-$dimensional spacetime, it is found that the Einstein field equation imposes a contracting wormhole solution and the obedience of the weak energy condition. Nevertheless, in the presence of an electric field, the latter presents a singularity at the throat, however, for a pure magnetic field the solution is regular. For the $(2+1)-$dimensional case, it is also found that the physical fields are singular at the throat. Thus, taking into account the principle of finiteness, which states that a satisfactory theory should avoid physical quantities becoming infinite, one may rule out evolving $(3+1)-$dimensional wormhole solutions, in the presence of an electric field, and the $(2+1)-$dimensional case coupled to nonlinear electrodynamics. ",
        "introduction": "Nonlinear electrodynamics has recently found many applications in several branches, namely, as effective theories at different levels of string/M-theory \\cite{Witten}, cosmological models \\cite{cosmo-models}, black holes \\cite{blackholes,Bronnikov1}, and in wormhole physics \\cite{Bronnikov1,wormhole1,wormhole2}, amongst others. Pioneering work on nonlinear electrodynamic theories may be traced back to Born and Infeld \\cite{BI}, where the latter outlined a model to remedy the fact that the standard picture of a point charged particle possesses an infinite self-energy. Thus, the Born-Infeld model was founded on a principle of finiteness, that a satisfactory theory should avoid physical quantities becoming infinite. Furthermore, Pleba\\'{n}ski extended the examples of nonlinear electrodynamic Lagrangians \\cite{Pleb}, and demonstrated that the Born-Infeld theory satisfy physically acceptable requirements.  In this context, in a recent paper, it was shown that $(2+1)$ and $(3+1)-$dimensional static, spherically symmetric and stationary, axisymmetric traversable wormholes cannot be supported by nonlinear electrodynamics \\cite{wormhole2}. This is mainly due to the presence of an event horizon and that the null energy condition is not violated. However, a particularly interesting situation arouse in the analysis of the $(2+1)-$dimensional static and spherically symmetric wormholes, namely, that in order to construct these geometries, they must necessarily be supported by physical fields that become singular at the throat. Thus, taking into account the principle of finiteness, and imposing a non-singular behaviour of the physical quantities, it was found that the wormhole possesses an event horizon, rendering the geometry non-traversable. We also point out that the non-existence of $(3+1)-$dimensional static and spherically symmetric traversable wormholes is consistent with previous results \\cite{Bronnikov1}.   Based on this analysis, it is also of interest to explore the specific case of evolving wormhole geometries, in the context of nonlinear electrodynamics. Time-dependent spherically symmetric wormholes have been extensively analysed in the literature, and a particularly interesting case of a dynamic wormhole immersed in an inflationary background, was considered by Roman~\\cite{Roman}. The primary goal in the Roman analysis was to use inflation to enlarge an initially small, possibly submicroscopic, wormhole, and test whether one could evade the violation of the energy conditions in the process. Further dynamic wormhole geometries were analysed, considering specific cases \\cite{Kar1,Kar2,Trobo}, and the evolution of a wormhole model was considered in a FRW background \\cite{Kim-evolvWH}. In the latter model, it was found that the total stress-energy tensor does not necessarily violate the energy conditions as in the static Morris-Thorne case \\cite{Morris,Visser}. Different scenarios for the weak energy condition violation were also explored, namely for a constant redshift function \\cite{Kar1,Kar2}, and an extension with a specific choice of a non-zero redshift function was further analysed \\cite{Trobo}.  Thus, in this work, we shall be interested in exploring the possibility that nonlinear electrodynamics may support time-dependent traversable wormhole geometries. This is of particular interest as the energy conditions are not necessarily violated in specific regions, as noted above. Therefore, analogously to Ref. \\cite{wormhole2}, we shall consider $(2+1)$ and $(3+1)-$dimensional spacetimes and find particularly interesting results. We find that for the $(3+1)-$dimensional case in the presence of an electrical field, the latter becomes singular at the throat, however, for a purely magnetic field, the solution is regular at the throat, which is extremely promising, and is in close relationship to the regular magnetic black holes coupled to nonlinear electrodynamics found in Ref. \\cite{Bronnikov1}. For the $(2+1)-$dimensional case we find that the physical fields become singular at the throat.     This paper is outlined in the following manner: In section \\ref{Sec:4dynWH} we analyze $(3+1)-$dimensional dynamic and spherically symmetric wormholes coupled with nonlinear electrodynamics, and in section \\ref{Sec:3dynWH}, $(2+1)-$dimensional evolving wormhole geometries in the context of nonlinear electrodynamics are studied. In section \\ref{Conclusion} we conclude. We shall use geometrised units, i.e., $G=c=1$, throughout this work . ",
        "conclusions": "In a recent paper, it was shown that $(2+1)$ and $(3+1)-$dimensional static, spherically symmetric and stationary, axisymmetric traversable wormholes cannot be supported by nonlinear electrodynamics. In this work, we explored the possibility of evolving time-dependent wormhole geometries coupled to nonlinear electrodynamics. For the $(3+1)-$dimensional spacetimes, it was found that the Einstein field equation imposes a contracting wormhole solution and that the weak energy condition be satisfied. It was also found that in the presence of an electric field, a problematic issue was verified, namely, that the latter become singular at the throat. However, regular solutions of traversable wormholes in the presence of a pure magnetic field were found. Time-dependent spherically symmetric $(2+1)-$dimensional wormhole spacetimes were also analyzed, and it was found that the Einstein field equation imposes that the electric field be zero. For this case, it was found that in order to construct wormhole geometries, these must necessarily be supported by physical fields that become singular at the throat. Thus, taking into account that the model construction of nonlinear electrodynamics, founded on the principle of finiteness, that a satisfactory theory should avoid physical quantities becoming infinite, one may rule out evolving $(3+1)-$dimensional electric wormhole solutions, and the $(2+1)-$dimensional case coupled to nonlinear electrodynamics.   It is also relevant to emphasize that the solutions obtained in this work and in Ref. \\cite{wormhole2}, can be obtained using an alternative form of nonlinear electrodynamics, denoted the $P$ framework \\cite{Bronnikov1}. The latter is obtained from the original form, the $F$ framework, by a Legendre transformation. The duality between the $F$ and $P$ frameworks connects solutions of different theories, but we emphasize that it is a dual description of the same physical system. Therefore, we have not made use of the $P$ formalism throughout this work, as we have only been interested in exploring the possible existence of evolving wormhole solutions coupled to nonlinear electrodynamics. Another point worth noting is that we have only considered that the gauge-invariant electromagnetic Lagrangian $L(F)$ be dependent on a single invariant $F$. As stressed in Ref. \\cite{wormhole2}, it would also be worthwhile to include another electromagnetic field invariant $G\\sim \\,^*F^{\\mu\\nu}\\,F_{\\mu\\nu}$, which would possibly add an interesting analysis to the solutions found in this work.        \\appendix"
    },
    "0608/astro-ph0608530.txt": {
        "abstract": "{Observations of present and future X-ray telescopes include a large number of serendipidious sources of unknown types. They are a rich source of knowledge about X-ray dominated astronomical objects, their distribution, and their evolution. The large number of these sources does not permit their individual spectroscopical follow-up and classification. Here we use Chandra Multi-Wavelength public data to investigate a number of statistical algorithms for classification of X-ray sources with optical imaging follow-up. We show that up to statistical uncertainties, each class of X-ray sources has specific photometric characteristics that can be used for its classification. We assess the relative and absolute performance of classification methods and measured features by comparing the behaviour of physical quantities for statistically classified objects with what is obtained from spectroscopy. We find that among methods we have studied, multi-dimensional probability distribution is the best for both classifying source type and redshift, but it needs a sufficiently large input (learning) data set. In absence of such data, a mixture of various methods can give a better final result. We discuss some of potential applications of the statistical classification and the enhancement of information obtained in this way. We also assess the effect of classification methods and input data set on the astronomical conclusions such as distribution and properties of X-ray selected sources.} ",
        "introduction": "Since the launch of Chandra and XMM-Newton X-ray observatories at the end of 90's our vision of the X-ray sky has been tremendously changed. The large collecting area of the XMM-Newton, and the sensitivity and high spatial resolution of Chandra (\\cite{cxcres}) have increased our knowledge about the details of the morphology and the energy spectrum of extended sources and the distribution and nature of point sources. They have also made it possible to study fainter and harder sources, and related physical processes. Each observed field, even with relatively modest exposure time of few kilo seconds, adds tens of new objects to the list of serendipitous sources. This gives the opportunity to increase our knowledge of individual sources as well as the general characteristics of classes of astronomical objects, and thereby the modeling of the underlying physical phenomena. In addition, in a large set of objects with various characteristics and behaviour, there is always the chance to find new classes/subclasses of sources or extreme and rare examples of known types (\\cite {faintx};\\cite {faintx1}).  Multi-wavelength spectroscopic follow-up of X-ray selected sources based on Chandra (\\cite {champ1}; \\cite {champ3}) and XMM-Newton (\\cite{axis}; \\cite {axis0}; \\cite {roberto}; \\cite {brightsam}) source catalogs have already investigated a number of important astronomical issues such as redshift evolution of different classes of Active Galactic Nuclei (AGN) (\\cite{agndust}; \\cite {agndust1}; \\cite {agndust2}; \\cite {agndust3}; \\cite {agndust4}; \\cite {agndust5}), their intrinsic absorption (\\cite{agnabs}; \\cite{agnabs1}; \\cite {agnabs2}; \\cite {agnabs3}) and their relationship with the characteristics of the host galaxy, its star formation history (\\cite{agnstar}; \\cite {agnstar1}; \\cite {agnstar2}; \\cite {agnstar3}; \\cite {agnstar4}; \\cite {agnstar5}) and the central super-massive black hole (\\cite{agnbh}; \\cite {agnbh1}; \\cite {agnbh2}). Observations have also revealed the relation between hard X-ray sources, hidden broad line AGN(BLAGN) or BL-Lac(\\cite{agnhard}; \\cite {agnhard1}; \\cite {agnhard2}) and X-ray background in hard (\\cite{agnxcb}; \\cite {agnxcb1}; \\cite {agnxcb2}; \\cite{agnxcb3}; \\cite {agnxcb4}) and soft bands (\\cite{softxcb}).  Although these observations have significantly improv-ed our knowledge about X-ray sources (\\cite{bepphard}; \\cite {bepphard1}; \\cite {bepphard2}), they only include relatively small number of spectroscopically identified X-ray objects. Despite the fact that they belong to randomly selected fields, it is not certain how much other criteria such as brightness of the optical counterparts and the tendency to study special classes of sources more than others, influence conclusions about populations and their physical conditions. Thus, it seems that much larger and unbiased samples representing different category of objects, with both typical characteristics of the population, and from extreme parts of the observable parameter space are necessary. It is however very difficult to obtain spectroscopic data for such large samples in a reasonable time.  An alternative to spectroscopic identification is classification based on multi-wavelength, photometric quantities such as fluxes, count-rates and hardness ratios. Abstractly speaking, the identification analysis exploits characteristics of the parameter space of spectroscopically identified objects to establish a series of rules to distinguish different classes. Unclassified objects are associated to a category according to these rules. This is similar to spectroscopic identification where observer must know the main spectral features discriminating one category of sources from others. The knowledge about characteristics of each class is either provided by the observer from prior information or can be found or learned in an automatic way using already classified objects. Some of the popular learning algorithms (\\cite{algo}) are mapping (\\cite{mapping}), neural network (\\cite{neuralnet}; \\cite {neurnetapp}), nearest neighbour (\\cite{algo}; \\cite {algo1}) and statistical algorithms (\\cite{stat})(Citations are examples and are not exclusive). These classification algorithms have an intrinsically statistical nature and as we will see later, depend on how much the learning sample is statistically representative of the population. For this reason we generally call them statistical identification methods.  In the present work we describe a number of statistical methods for classifying X-ray selected sources with candidate optical counterparts. Our aim is not just suggesting few algorithms but also assessing their reliability and their effect on astronomical conclusions. The input set of spectroscopically classified sources are taken from the publicly available Chandra Multi-wavelength project (ChaMP) (\\cite{champ2}) identifications. The unclassified sources to which we apply these methods are also taken from fields observed by ChaMP, and therefore are subject to the same selection conditions as the input data. No other prior assumption is made about sources or classes. As this paper concentrates on the technical aspects of classification, we don't investigate the completeness of the input sample, and therefore astronomical conclusions obtained from this data can not be extended to other data sets without caution.  In the following sections we first briefly review the X-ray and optical quantities used in the classification. Then, we explain methods we have investigated and compare their performance. Finally, we apply methods considered to be more efficient to a list of unclassified sources and show the effect of the input set as well as the classification method on the distribution of statistically classified sources and there-by, on the scientific conclusions.  \\section {X-ray and Optical Data} \\label {sec:data} ChaMP includes 137 high galactic latitude fields selected for having hydrogen column densities $N_h < 6 \\times 10^{20} cm^{-2}$ or optical extinction $E(B-V) < 0.1$mag (\\cite{champ1}). For the first public release of ChaMP spectroscopical identifications, 62 of these fields which have PI authorization have been used. With additional constraints explained in detail in \\cite{champ1}, it is expected that full ChaMP fields include $\\sim 6000$ background X-ray sources and indeed the public data available from ChaMP web site\\footnote{{\\scriptsize http://hea-www.cfa.harvard.edu/CHAMP/IMAGES\\_DATA/champ\\_xpc.tab}} includes 6512 X-ray sources with photometric information i.e. count rates, fluxes and hardness ratios.  The X-ray flux is calculated in 5 overlapping energy bands defined in Table \\ref{tab:xband}.  \\begin{table}[ht] \\caption{ChaMP X-ray energy bands \\label{tab:xband}} \\vspace{0.2cm} \\begin{tabular}{ll} \\hline Band & Energy \\\\ \\hline Broad (B) & 0.3-8.0 keV \\\\ Hard (H) & 2.5-8.0 keV \\\\ Soft (S) & 0.3-2.5 keV \\\\ Soft1 (S1) & 0.3-0.9 keV \\\\ Soft2 (S2) & 0.9-2.5 keV \\\\ \\hline \\end{tabular} \\end{table} In the classification methods described here we use three distinct bands S1, S2 and H. The broad band B is also used for normalisation when use, as parameter space, the ratio of fluxes to the X-ray flux. We also define two hardness ratios: \\bea H_{S1S2} &=& \\frac {S2-S1}{S2+S1}. \\label {hss}\\\\ H_{S2H} &=& \\frac {H-S2}{H+S2}. \\label {hhs} \\eea where $H$, $S1$ and $S2$ are total counts in the corresponding band. Energy conversion factors for 62 fields of the first release are given in \\cite{champ1} for three power-law spectrum models with $\\Gamma = 1.2,~1.4,~1.7$. For compatibility with what is used for the XMM-Newton data, through out this work we use fluxes determined with $\\Gamma = 1.7$.  Optical follow-up of the first 6 ChaMP fields for which observations and identifications are complete, was perform-ed by using NOAO telescope for both southern and northern fields (\\cite{champ2}). Optical filters for imaging correspond to Sloan filters $i$, $r$ and $g$. For northern fields if the counterparts are not too faint, it is also possible to use Sloan public data. The advantage is that it includes 2 more filters $u$ and $z$. We will show later the importance of having information in more filters, especially in $u$. For this reason, we have retrieved the 5- filter optical magnitudes from the Sloan 4$^{th}$ release for northern fields. Between unclassified sources in ChaMP fields, we use only objects with all Sloan 5 filters. It is not however possible to apply the same restriction to spectroscopically identified sources, because their number is quite limited. For identified sources out of Sloan coverage, we use ChaMP published magnitudes and estimate the missing $u$ and $z$ magnitudes by modeling color difference for each category of sources.  To model magnitudes, we define a color curve for each source with all the 5 magnitudes mentioned above. They are ordered in increasing central frequency and the difference of successive magnitudes - colors - are determined. For each category of sources we obtain the average color curve, considered as template for the category, and its standard deviation. We assume a Gaussian distribution around this template.  To estimate $u$ and $z$ magnitude for identified sources without these measurements, we choose a random value for $u-g$ and $i-z$ according to a Gaussian distribution with average and deviation taken from the template color curve and its deviation. Figs. \\ref{fig:colormodel}-a, b and c show the color and redshift distribution of sources with $u$ and $z$ magnitudes from Sloan archive and the distribution of sources for which their $u$ and $z$ magnitudes are estimated. These plots show that measured and modeled distributions are statistically similar. Evidently, at source by source level there is no guarantee that the estimated magnitudes correspond to real ones for a given source. But statistically speaking, there can be sources with $i$, $r$, and $g$ magnitudes as one of these sources and $u$ and $z$ magnitudes close to what we have estimated~\\footnote{Note that we have used only optical magnitudes $i$, $r$, and $g$ for extrapolating to $z$ and $u$. It would be better to use both optical and X-ray fluxes, but this makes the model too complicated.}. This level of precision is enough for us because our purpose here is studying the methodology of classification and not real properties of the sources we use. Therefore, a set of data with statistically the same properties as what can be encountered in reality is sufficient.  \\begin{figure*}[t] \\begin{center} \\begin{tabular}{cc} a) \\includegraphics[height=7.5cm]{fig1a.eps} & %{ugrizfit-gr-ug-new.eps} & b) \\includegraphics[height=7.5cm]{fig1b.eps} \\\\ %{ugrizfit-color-new.eps} \\\\ \\multicolumn{2}{c}{c) \\includegraphics[height=7.5cm]{fig1c.eps}} %{ugrizfit-z-ug-new.eps}} \\end {tabular} \\caption{(a) $u-g$ versus $g-r$ for data with $u$ magnitude (top) and modeled data without archival $u$ magnitude. Source type: BLAGN (violet), NLEG (magenta), ALG (cyan), star (gold). (b) Color distribution for sources in (a). Dash and full lines present respectively sources with redshift $\\leqslant 1.5$ if the source is a BLAGN or $\\leqslant 0.6$ otherwise. (c) Redshift distribution. To make the input sample to identification methods more uniform, 2 sources with largest deviation are removed before further application. \\label {fig:colormodel}} \\end{center} \\end{figure*}  \\subsection {Spectroscopic Identifications} The first ChaMP public data release includes 125 spectroscopically identified sources (\\cite{champ2}). Most of them are in the medium flux category i.e. have $10^{-15} \\lesssim f \\lesssim 10^{-14}$ \\enerunit in energy range $0.5$ keV $\\leqslant E \\leqslant 2$ keV. A small set of sources have fluxes as small as few times $\\sim 10^{-16}$ \\enerunit or as large as few times $\\sim 10^{-13}$ \\enerunit. Objects are classified in 4 categories: Broad Line Active Galaxy Nucleus (BLAGN), Narrow Emission Line Galaxy (NELG), Absorption Line Galaxy (ALG) and star. The details of criteria for associating a source to one of these categories can be found in \\cite{champ2}. For the purpose of completeness here we summarize them. BLAGN category is designated to sources with board band - FWHM $> 1000$ km~sec$^{-1}$ - emission lines. If all emission lines have FWHM $< 1000$ km~sec$^{-1}$, the source is classified as NELG. Objects with high $S/N > 20$ and no identifying feature are classified as BL-Lac. Sources with absorption lines are classified as ALG or stars. Their spectrum is compared to synthetic templates for classification and for determination of their redshift. A classification confidence flag is given to each source: 0 not classifiable, 1 insecure, 2 secure. Here we use only sources with secure identification in the input dataset.  The set of 125 serendipitiously identified sources is too small for statistical applications. Therefore, we have added also sources from two other published set of identified sour-ces by the ChaMP, i.e. hard AGNs (\\cite{cxcagnhardnew}) and normal galaxies (\\cite{cxcnormgal}). For this set of sourc-es as before we use the publicly available ChaMP data for x-ray, Sloan optical magnitude in 5 bands as explained above, and modeling when some of the optical magnitudes are mis-sing. After neglecting sources for which reliable data was not available and modeling couldn't give reasonable values, we put together a total of 268 sources with identifications and all the flux/magnitude information. This set includes 151 BLAGNs, 65 NELGs. 37 ALGs and 15 stars. It is clear that in this set the number of stars is under-estimated with respect to what one can expect from a serendipidious selection of X-ray sources in high galactic latitude fields. Although the addition of ChaMP hard AGN and normal galaxies have increased the size of the learning set for the statistical identification, due to the bias of these set toward a special type of objects, the final set does not anymore presents what one can find in a set of randomly selected sources subjected only to conditions on X-ray fluxes and optical magnitudes. This is not a major problem for present work where the main purpose is the investigation of performance of classification methods. However, we will show later that the learning (input) set to algorithms or classifying agents such as humans, inevitably leaves its trace on the classification and this is inherent to the concept of learning and extension of information. This issue should be considered when the results of statistical classification are used for obtaining scientific conclusions. ",
        "conclusions": "Irrespective of the method used for statistical classification, the input set leaves its fingerprint on the physical conclusions. Therefore, only if the input data is sufficiently - with but a few percent of fluctuations - representative of the wh-ole population, one can be confident on the reliability of the deduced conclusions about the physical behaviours from statistical classification. By giving some examples in this section, we show that the available data set is yet too small and a much larger input set is needed. At the same time, these examples clarify how one should investigate artifacts left in the output of the statistical classification from the input data or classification methods.  We have already seen in Fig. \\ref {fig:unidflux} the effect of under-representation or over representation of a category of objects on the classification. Fig. \\ref{fig:frfxfx} shows the distribution of corresponding statistically classified sources in $\\log (f_x/f_r)$-$\\log (f_x)$ plane. Suppose that from these plots we want to judge about X-ray to optical luminosity of different category of sources. According to the biased population of Fig. \\ref{fig:frfxfx}-b, significant number of stars have a very hard spectrum i.e. $f_x \\gg f_r$. But less biased classification of Fig. \\ref{fig:frfxfx}-a and spectroscopically classified sources Fig. \\ref{fig:frfxfx}-c don't show such a population. Similar effects are also visible at low X-ray luminosity, low X-ray to optical flux tail where in Fig. \\ref{fig:frfxfx}-b this area is dominated by stars but Figs. \\ref{fig:frfxfx}-a and \\ref{fig:frfxfx}-c show a significant contribution from NELG and ALG/gal. This has important implications for absorption column density in these types of extra-galactic objects. \\begin{figure*}[p] \\begin{center} \\begin{tabular}{cc} a)\\includegraphics[height=7cm,angle=-90]{fig16a.eps} & %{cxcunidclassflux-data-flux-frfx-fx.eps} & b)\\includegraphics[height=7cm,angle=-90]{fig16b.eps} \\\\ %{cxcunidclassopt-data-opt-frfx-fx.eps} \\\\ \\multicolumn{2}{c}{c)\\includegraphics[height=7cm,angle=-90]{fig16c.eps}} %{xxidedredfrfx-fx.eps}} \\end{tabular} \\caption{Distribution of statistically classified sources in $\\log (f_x/f_r) $-$\\log (f_x)$ plane. Plots a) and b) correspond to sets explained in Fig. \\ref {fig:unidflux}. The same distribution for spectroscopically identified is also shown in c). Definition of colors and symbols is the same as in Fig. \\ref{fig:hf}.} \\label{fig:frfxfx} \\end{center} \\end{figure*}  These are examples of misleading conclusions due to incorrect classification. Thus before any scientific conclusion be taken from statistical classifications, we must first assess their uncertainty and its effect. One way to take this into account could be to add an uncertainty to the distributions coming from a statistical classification. This is not however a simple task. The test of a classification algorithm is based on a small sample in which a classification bias can be difficult to detect. When such a method is applied to unknown objects, it is very difficult to know whether an observed behaviour of data is intrinsic or due to wrong classification. Moreover, quantifying the uncertainty of distributions is not simple either. We must also add the uncertainty about how representative is the set of spectroscopic identifications to classification uncertainty. There are a few more issues also to consider: to what extent is the selection of a source for spectroscopic follow-up a random process ? How much do observational conditions, source location on the sky, visibility to a special telescope, etc. bias the selection of sources for follow-up ? Or in general, does the set of followed-up fields present a random sampling of the whole sky ? One way to answer these questions is to divide the sample, plot separate parts and try to find any systematic difference. This is possible only if the size of the sample is enough large such that it can be dived to statistically significant subsets.  Before concluding this section we want to add few remarks about the general purpose of statistical identification, in X-ray or other domains of astronomy. As we mentioned in this and previous sections, whatever the classification me-thod, the input data set leaves its imprint on the statistical identified sources. This is quite natural because it is the base of information. Then, the question arises: what more we can learn about astronomical issues than what is in the input (learning) set ? Probably the most important and unbiased ones are issues related to quantities which are not used in the process of classification. For instance, large scale spatial correlation of different type of objects. More applications can be found with combination of observables used in the classification - here fluxes - and what is not used in classification such as spacial distribution. Note also that redshift is not an independent quantity because it is estimated from the knowledge coming from the input set."
    },
    "0608/gr-qc0608091_arXiv.txt": {
        "abstract": "We present numerical results from three-dimensional evolutions of scalar perturbations of Kerr black holes. Our simulations make use of a high-order accurate multi-block code which naturally allows for fixed adaptivity and smooth inner (excision) and outer boundaries.  We focus on the quasinormal ringing phase, presenting a systematic method for extraction of the quasinormal mode frequencies and amplitudes and comparing our results against perturbation theory.  The detection of a single mode in a ringdown waveform allows for a measurement of the mass and spin of a black hole; a multimode detection would allow a test of the Kerr nature of the source.  Since the possibility of a multimode detection depends on the relative mode amplitude, we study this topic in some detail. The amplitude of each mode depends exponentially on the starting time of the quasinormal regime, which is not defined unambiguously. We show that this {\\it time-shift problem} can be circumvented by looking at appropriately chosen {\\em relative} mode amplitudes. From our simulations we extract the quasinormal frequencies and the relative and absolute amplitudes of corotating and counterrotating modes (including overtones in the corotating case). We study the dependence of these amplitudes on the shape of the initial perturbation, the angular dependence of the mode and the black hole spin, comparing against results from perturbation theory in the so-called asymptotic approximation. We also compare the quasinormal frequencies from our numerical simulations with predictions from perturbation theory, finding excellent agreement. For rapidly rotating black holes (of spin $j=0.98$) we can extract the quasinormal frequencies of not only the fundamental mode, but also of the first two overtones. Finally we study under what conditions the relative amplitude between given pairs of modes gets maximally excited and present a quantitative analysis of rotational mode--mode coupling. The main conclusions and techniques of our analysis are quite general and, as such, should be of interest in the study of ringdown gravitational waves produced by astrophysical gravitational wave sources. ",
        "introduction": "One of the most useful methods to explore the response of black holes to external perturbations is based on wave scattering \\cite{Futterman88}. Early studies identified three main stages in the dynamics of a wave propagating on a black hole background, as observed at a fixed spatial point. In a first, transient phase the observed wave depends on the structure of the initial pulse. Vishveshwara and Press discovered that this initial ``burst'' is invariably followed by a second phase characterized by exponentially decaying oscillations: this phase is usually referred to as ``quasinormal ringing'' \\cite{Vishveshwara70b,Press71}.  In the third and last stage of the evolution, waves slowly die off as a power law tail \\cite{Price72}.  Astrophysical black holes should be well described by the Kerr solution, since charge is unlikely to play a major role in astrophysical scenarios (see e.g.\\ \\cite{Berti:2005eb} for a discussion).  As a consequence of the ``no hair theorem'', if general relativity is the correct theory of gravity, the quasinormal mode (QNM) frequencies of a Kerr black hole depend only on its mass and angular momentum. Earth-based and space-based gravitational wave detectors have the potential to measure the frequency and damping time of a QNM. From these two observables we can infer the black hole's mass and angular momentum \\cite{Echeverria89,Finn92}.  For the space-based {\\it Laser Interferometer Space Antenna} (LISA) , and possibly also for second-generation ground-based detectors, the signal-to-noise ratio can be large enough that we will be able to identify {\\it two} or more QNM frequencies in the signal \\cite{Berti06b}.  A multi-mode detection would provide a striking, direct test of the Kerr nature of the source (i.e., of the no-hair theorem). The basic idea is quite simple. Roughly speaking, the first mode in the pair is used to determine the black hole's mass and angular momentum, and the other mode(s) to verify that the QNM spectrum is indeed consistent with a general relativistic Kerr black hole \\cite{Dreyer04}.  The combined observation of supermassive black hole binary inspiral and ringdown with LISA can provide even more information \\cite{Flanagan:1998a}.  Parameter estimation during the inspiral phase can be very accurate, depending on the black holes' masses, spins and distance \\cite{Berti05a}. Combining information from the inspiral and ringdown phases we can estimate the energy radiated in the merger, and possibly improve parameter estimation from both phases (see e.g.\\ \\cite{Luna06} for a preliminary study of this effect in the context of earth-based detectors).  In the last thirty years the development of gravitational wave astronomy motivated a detailed investigation of the QNM frequency spectrum \\cite{Kokkotas99b,Nollert99,Berti04a}. In comparison, the problem of the {\\it relative excitation} of QNMs received very little attention (see e.g. \\cite{Berti:2006wq} and references therein).  Ideally, the relative QNM excitation should be determined by general relativistic simulations of binary black hole mergers. Despite recent progress, this information is not yet available \\cite{Berti06a}.  Given the recent progress of numerical relativity, by the time LISA flies we could have a good knowledge of the multipolar distribution of the energy and angular momentum radiated in a black hole merger under generic conditions. Knowing in advance which modes should be excited in a realistic merger will not only be useful to probe the Kerr nature of the source, but also to reduce the number of templates needed to perform matched filtering on ringdown waveforms.  In this paper we present a quantitative investigation of QNM excitation studying a simple model problem: the scattering of scalar waves on a Kerr background. We use our new infrastructure for multi-block simulations \\cite{Schnetter06a} for these studies. The infrastructure is based on the techniques described and applied in the context of numerical relativity in \\cite{Lehner2005a} and further extended in \\cite{Diener05b}.  Our infrastructure uses Carpet \\cite{Schnetter-etal-03b, carpetweb}, a driver for the Cactus computational toolkit \\cite{Goodale02a, cactusweb1}, originally designed to provide fixed and adaptive mesh refinement. The capabilities of Carpet have been recently extended \\cite{Schnetter06a} to include the type of boundary conditions that are needed for multi-block (also called multi-patch) simulations in Cactus.  Multi-block techniques yield increased efficiency and accuracy in our studies, for two main reasons. The first is that we can set up smooth excision and outer boundaries, and we can therefore apply boundary conditions in a clean and well understood way.  Furthermore, a suitable multi-block grid structure provides a natural and flexible way of implementing mesh refinement. We can keep a fixed angular resolution throughout the entire domain, avoiding the unnecessary high resolution at large distances from the central object that one would have using a cartesian grid. The resources thus saved can be used, for example, to set up a rather large number of grid points in the radial direction, so that outer boundaries are located at large radii and the noise produced at the boundaries does not affect the results. Even though we are working in three dimensions, our results are more accurate than previous studies using two-dimensional codes \\cite{Krivan96a,Krivan97a}. Krivan {\\it et al.}~\\cite{Krivan96a} studied the late time dynamics and the rotational coupling of massless scalar fields in a Kerr background, but not their quasinormal ringing.  Later they extended the analysis to gravitational perturbations, considering both the late time tail and the quasinormal ringing phase~\\cite{Krivan97a}.  For large rotation the damping times of corotating fundamental modes in \\cite{Krivan97a} are accurate within $\\sim 3\\%$ when compared to results from perturbation theory; our accuracy ($\\sim 0.3 \\%$) is roughly an order of magnitude better. In fact, we can extract the frequencies of some {\\em overtones} with an error of the order of a few percent or less.   Given the high accuracy of our multi-block infrastructure, a careful extraction of the QNM content of the waveforms becomes necessary. We discuss in detail the so-called {\\it time-shift problem} (exponential dependence of the quasinormal amplitudes on the time at which the quasinormal ringing regime starts), how it affects the determination of both absolute and relative QNM amplitudes, and how to choose pairs of modes so as to decrease the uncertainty on relative amplitudes. We also introduce a general criterion (based on minimizing a suitably defined residual) to determine the optimal fitting window to extract QNM frequencies and amplitudes. Using these tools we study the absolute and relative amplitudes of corotating and counterrotating modes for Gaussian initial data located in the far zone. We study the dependence of these amplitudes on the radial shape of the initial data, finding excellent agreement with results from perturbation theory \\cite{Berti:2006wq}.  We also discuss the problem of extracting overtones for modes with a given angular dependence, finding that the first overtones of corotating modes (e.g.\\ modes with $l=m=2$) contribute significantly to the waveform for rapidly rotating black holes.  The plan of the paper is as follows. In Sec.~\\ref{we} we briefly describe our multi-block code. After introducing the background metric, we discuss the numerical implementation of the scalar wave equation and our time evolution techniques. In Sec.~\\ref{QNM} we extract QNM frequencies and amplitudes from our evolutions, comparing with analytical predictions from perturbation theory. To start with, we point out some conceptual limitations in the extraction of QNM amplitudes due to the so-called {\\it time-shift problem}. Then we introduce a rather general method to determine the best fitting interval to extract QNM waveforms. We first check the accuracy of this method (and of our numerical code) by reproducing the QNM frequencies predicted by standard perturbation theory. Scalar QNM frequencies for Kerr black holes have been computed in \\cite{Berti:2005eb}, and they have never been systematically confirmed by numerical time evolutions\\footnote{See however \\cite{Saijo97a}, where the fundamental scalar mode with $l=0$ was observed to dominate the emission of scalar radiation by perturbed Kerr black holes in the Brans-Dicke theory of gravity.}.  Next we give a quantitative estimate of rotational mode mixing as a function of the black hole's spin and discuss the initial data dependence of the amplitudes of corotating and counterrotating modes. In this way we assess the validity of the amplitudes predicted by perturbation theory in the so-called asymptotic approximation (where both the observer and the initial data are located far away from the black hole). Finally we discuss the extraction of overtones from our waveforms. ",
        "conclusions": "The chances of a multi-mode detection by either Earth- or space-based gravitational wave detectors will depend on the relative amplitude of those modes. Knowing in advance which modes should be excited under a realistic binary merger would reduce the dimensionality of the template bank needed to perform matched filtering on ringdown waveforms. An answer that numerical relativity might provide is precisely which modes are likely to be dominant. This involves predicting the relative amplitudes of different pairs of modes under a variety of scenarios. In this paper we have taken a first step towards understanding the issues involved in such a prediction.  We first presented a systematic way of extracting QNMs from a given signal. Our procedure has a number of built in self-consistency checks, to make sure that when we keep adding modes to our fit we are fitting a true signal and not numerical noise. One of these self-consistency checks is to make sure that we extract the correct quasinormal frequency of each mode within a certain accuracy. If the data being analyzed comes from a numerical simulation, consistent frequencies can be used to monitor the accuracy of the code. If the data is experimental, consistency of the frequencies allows for a test of the no hair conjecture. In more detail: during our fitting procedure we first fit for the dominant mode(s), look at the residual (defined as the difference between our original signal and the fit), make sure that it has a consistent quasinormal ringing behavior and only then fit for the next set of modes, repeating the procedure as long as it makes sense to do so. By following this procedure we have hardly been able to go beyond the first few dominant modes, and this was only possible in very special cases. We expect this to happen with most numerical simulations.  We addressed in some detail the so-called {\\it time-shift problem}. In essence, this is the fact that the quasinormal amplitudes depend exponentially on the quasinormal ringing excitation time, which is not defined unambiguously (not even in the continuum). Furthermore, examining actual values of quasinormal frequencies we have seen that this exponential dependence is an important factor to take into account in practice.  To (partially) get rid of this exponential dependence we propose to look at relative amplitudes: choosing pairs of modes whose damping frequency is as close as possible, we can partially cancel each others exponential dependence. We analyzed in detail the exponential dependence of different pairs of modes as a function of the black hole spin.  In particular, we found that the {\\it time-shift problem} becomes more important as one increases the spin. For modes with the same value of $\\ell$, for example, the problem is not very relevant for spins $j \\lesssim 0.5$. On the other hand, an accurate extraction of the relative amplitude between the fundamental mode and the first overtone only seems feasible for very high spins and $m>0$.  Keeping this in mind, we first extracted the fundamental quasinormal frequencies for different values of spin, ranging from $j=0$ to a rapidly rotating black hole with $j=0.98$. Even using modest resolutions our frequencies agree with those obtained from perturbation theory within one part in $10^5$ to one part in $10^2$, depending on the black hole spin, location of the observer and angular dependence. To our knowledge this is the first time that quasinormal frequencies for scalar perturbations of Kerr, as predicted by perturbation theory, have been verified by numerical evolutions of the field equations.  Next we analyzed in detail the relative amplitude of corotating and counterrotating fundamental modes, as a function of the width of the initial perturbation and the black hole spin, being able to quantify (within the limitations imposed by the {\\it time-shift problem}) under what conditions the asymptotic approximation of Ref.~\\cite{Berti:2006wq} is valid. In particular, we were able to verify the widths of the initial perturbation corresponding to the maximal QNM excitation. Finally, we studied the excitation of overtones. We found that, according to expectations from perturbation theory \\cite{Berti:2006wq}, they get significantly excited for corotating modes and very high spins. In this particular case we were able to extract the complex QNM frequency for the fundamental mode and the first two overtones, with a difference with respect to the predicted values by perturbation theory of the order of a tenth of a percent to ten percent, depending on the mode and the black hole spin. We expect the techniques and results of this paper to be general enough to be useful for future work on ringdown waveforms."
    },
    "0608/astro-ph0608535_arXiv.txt": {
        "abstract": "We show that a star orbiting close enough to an adiabatically grown supermassive black hole can capture a large number of weakly interacting massive particles (WIMPs) during its lifetime.  WIMP annihilation energy release in low- to medium-mass stars is comparable with or even exceeds the luminosity of such stars due to thermonuclear burning.  The excessive energy release in the stellar core may result in an evolution scenario different from what is expected for a regular star.  The model thus predicts the existence of unusual stars within the central parsec of galactic nuclei.  If found, such stars would provide evidence for the existence of particle dark matter. White dwarfs seem to be the most promising candidates to look for. The signature of a white dwarf burning WIMPs would be a very hot star with mass and radius characteristic for a white dwarf, but with luminosity exceeding the typical luminosity of a white dwarf by orders of magnitude $\\la$$50L_\\odot$. A white dwarf with a highly eccentric orbit around the central black hole may exhibit variations in brightness correlated with the orbital phase. ",
        "introduction": "The nature of the non-baryonic dark matter, which dominates the visible matter by about 4:1, is perhaps the most interesting experimental challenge for contemporary particle astrophysics.  A hint for a solution has been found in particle physics where the WIMPs arise naturally in supersymmetric extensions of the Standard Model \\citep[e.g.,][]{haber-kane}, among other possibilities.  The WIMP is typically defined as a stable, electrically neutral, massive particle. Assuming that non-baryonic dark matter is dominated by WIMPs, the pair annihilation cross-section is restricted to provide the observed relic density \\citep{jkg96,bergstrom00}.  A pair of WIMPs can annihilate producing ordinary particles and \\grays.  WIMPs are expected to form high density clumps according to N-body simulations of test particles with only gravitational interactions \\citep{Navarro97,Moore99}.  The highest density ``free space'' dark matter regions occur for dark matter particles captured within the gravitational potential of adiabatically grown supermassive black holes \\citep{gs99,bm05}.  Higher dark matter densities are possible for dark matter particles captured inside of stars or planets.  Any star close enough to an adiabatically grown supermassive black hole can capture a large number of WIMPs during a short period of time.  Such an idea has been first proposed by \\citet{salati89} and further developed by \\citet{bouquet89} who applied it to main-sequence stars.  WIMP annihilation in stars may lead to considerable energy release in the stellar cores thus affecting the evolution and appearance of such stars.  The model led to the conclusion of suppression of stellar core convection, thus predicting a concentration of stars in the Galactic center masquerading as cold red giants.  We perform order-of-magnitude estimates of WIMP capture rates for stars of various masses and evolution stages located in high density dark matter regions. We use current limits on WIMP-nucleus interaction and WIMP annihilation cross sections, as well as recent estimates of WIMP energy density near an adiabatically grown supermassive black hole. We argue that white dwarfs, fully burned stars without their own energy supply, are the most promising candidates to look for. ",
        "conclusions": "Where does the energy released during the WIMP annihilation go?  Table \\ref{table1} shows that the typical radius of the thermal distribution of WIMPs in the star core $r_{th}\\ll R_*$, therefore, the products of their annihilation can not propagate to the star surface and are eventually converted into the thermal energy and neutrino emission. This may be not true for the white dwarfs whose radii are comparable to $r_{th}$, and where the effect of gravitational ``focusing'' is important (eqs.~[\\ref{C1}], [\\ref{C2}], second term in square brackets). While most of the energy is converted into heat, the surface of white dwarfs may emit products of WIMP annihilation, particles ($e^\\pm$, $p$, $\\bar p$, $\\nu$, $\\tilde{\\nu}$) and \\grays, and this might be visible with a \\gray\\ telescope.  The WIMP capture rate generally increases with the star mass. However, the relative importance of WIMP annihilation diminishes in massive stars since their own luminosity increases as $\\propto M_*^{3.5}$. Therefore, the interiors of massive stars will remain largely unaffected by WIMPs.  The WIMP burning affects mostly stars with $M\\la M_\\odot$ with low to moderate luminosities and a large rate of WIMP capture.  The typical luminosity of a sun-like star burning Hydrogen is $L_\\odot=3.8\\times10^{33}$ erg s$^{-1}$, while WIMP annihilation in its core can provide up to $L_\\chi\\sim4.6\\times10^{34}$ erg s$^{-1}$, i.e.\\ 10 times more. Such an additional energy source may exceed the star's own resources from nuclear burning.  The large concentration of WIMPs may also lead to more effective energy transport in the stellar core \\citep{bouquet89} thus decreasing the core temperature and limiting the replenishment of the burning core with fresh nuclear fuel.  Combined, these two effects could change the star's evolution scenario, life expectancy, and appearance.  A white dwarf, a star without its own energy supply consisting of Carbon and Oxygen, may emit $L_\\chi\\sim2\\times10^{33}-2\\times10^{35}$ erg s$^{-1}$, i.e.\\ 0.4--50 times luminosity of the sun, burning WIMPs only and this energy source will last forever!  To reach such a luminosity, the surface temperature of the white dwarf should be in the range of $T\\sim(80-150)\\times10^3$ K assuming $R=0.01R_\\odot$.  In comparison, the luminosity of a regular white dwarf does not exceed $\\sim$$0.01L_\\odot$. The interiors of white dwarfs are almost isothermal, and the energy transport is dominated by degenerate electrons \\citep[see][for a recent review]{hansen04}; therefore, the large number of captured WIMPs and their annihilation in the core would not change the structure of white dwarfs. Interestingly, since $L_\\chi\\propto\\rho_\\chi$, a population of such white dwarfs located at different distances from the central black hole will exhibit a luminosity correlated with the radial WIMP density distribution.  A signature of a ``dark matter burner'' would be a very hot star with a mass and radius characteristic for a white dwarf, but with a luminosity exceeding the typical luminosity for a white dwarf by orders of magnitude $\\la$$50L_\\odot$.  Note that an independent determination of the $M_*/R_*$ ratio is possible using the gravitational redshift which has to be equivalent to $v_R\\sim50$ km s$^{-1}$ \\citep{gt67}.  Furthermore, a low-mass star with a highly eccentric orbit around the central black hole may exhibit variations in brightness correlated with the orbital phase.  The brightness should increase as the star approaches the periastron and some time after that. To have this working, the orbital period should essentially exceed $\\tau_{eq}$.  Carbon burning stars have $\\tau_{eq}\\sim10$ yr. It is even shorter in case of a white dwarf $\\tau_{eq}\\sim0.5$ yr. If a white dwarf appears in a high-density WIMP region, the WIMP density in its material quickly reaches equilibrium. On the other hand, the small radii of white dwarfs $R_*\\sim 10r_{th}$ imply that the captured WIMP density close to the surface may be large enough thus the surrounding WIMP density change will result in variations in their brightness.  This makes white dwarfs ideal objects to test the WIMP density in the environment in which they are orbiting.  Advances in near-IR instrumentation have made possible the observation of stars in the inner parsec of the Galaxy \\citep{genzel00,ghez03,ghez05}.  The observed absorption line widths imply high temperatures and lead to a ``paradox of youth:'' apparently young stars in the region whose current conditions seem to be inhospitable to star formation.  One of the possibilities is that they are old stars masquerading as youths. If an independent determination of their mass reveals that some of them are low-mass stars, then they would become candidates for the dark matter burners.  The model predicts the existence of unusual stars and consequently unusual supernovae in the central parts of the galaxies. These are the stars burning WIMPs at a high rate whose luminosities essentially exceed the typical luminosity of a star of given mass. If found, such stars would be interesting probes of particle dark matter near supermassive black holes. Their luminosity, or rather their excess luminosity, attributed to the WIMP burning can be used to derive the WIMP matter density at their location. On the other hand, a lack of such unusual stars may provide constraints on WIMP density, WIMP-nucleus scattering and pair annihilation cross-sections."
    },
    "0608/astro-ph0608703_arXiv.txt": {
        "abstract": "{The  magnetorotational instability (MRI) in cylindrical Taylor-Couette flow with external helical magnetic field is simulated for infinite and finite aspect ratios.  We solve the MHD equations in their small Prandtl number limit and confirm with time-dependent nonlinear simulations that the additional toroidal component of the magnetic field reduces the critical Reynolds number from \\ord{10^6} (axial field only) to \\ord{10^3} for  liquid metals with their small magnetic Prandtl number. Computing the saturated state  we obtain velocity amplitudes which help designing  proper experimental setups. Experiments with liquid gallium require axial field $\\sim 50$ Gauss and axial current $\\sim 4$ kA for the toroidal field. It is sufficient that the vertical velocity $u_z$ of the flow can be measured with a precision of 0.1mm/s.\\\\ We also show that the endplates enclosing the cylinders do not destroy the traveling wave instability which can be observed as presented in earlier studies.  For TC containers without and with endplates the angular momentum transport of the MRI instability is shown as to be outwards.} ",
        "introduction": "The magnetorotational instability (MRI) has been formulated long time ago (Velikhov 1959) but recently it become of particular interest for astrophysics due to recognition of its importance as the source of turbulence in accretion disks with Keplerian rotation (Balbus \\& Hawley 1991). One of the most convenient models for MRI is magnetohydrodynamical cylindrical Taylor-Couette flow, i.e. motion of liquid metal between two concentric differentially rotating cylinders under imposed external magnetic field along the axis (R\\\"udiger \\& Zhang 2001; Ji, Goodman \\& Kageyama 2001; Willis \\& Barenghi 2002). Under the influence of a purely axial field MRI operates for magnetic Reynolds number $\\Rm=\\Omin \\Rin(\\Rout-\\Rin)/\\mdiff$ of order \\ord{10} and has never been observed in an experiment due to very small Prandtl number for laboratory liquids (see Table~\\ref{tab-physprop}) which results in need of high critical Reynolds numbers of order $10^6$  . Recently it has been shown by Hollerbach \\& R\\\"udiger (2005) and R\\\"udiger et al. (2005) that a current-free external toroidal magnetic field in addition to the  usual axial field can reduce the critical Reynolds number  to \\ord{10^3} and therefore it makes it much easier to design a MRI experiment due to slower rotation.  The aim of the presented paper is to compute the amplitudes of velocities and the frequencies in a nonlinear saturated state under the influence  of the endplates covering the cylinders.  \\begin{figure} \\center \\includegraphics[width=6.0cm]{container.eps} \\caption{The geometry of the problem -- two concentric cylinders with radii $\\Rin=4$ cm, $\\Rout=8$ cm rotating with $\\Omin$, $\\Omout$. $B_z$ and $B_\\phi$ are the external magnetic fields. $B_\\phi$ is due to an axial current inside the inner cylinder. In the PROMISE experiment cylinders have finite length of $H=40$ cm ($\\Gamma=10$) and are covered with endplates. The liquid is a mixture of gallium, indium and tin (Stefani et al. 2006).} \\label{fig-geom} \\end{figure}    The fluid confined between the cylinders is assumed to be incompressible and dissipative with the kinematic viscosity $\\nu$ and the magnetic diffusivity $\\mdiff$. Derived from  the conservation of angular momentum the rotation law $\\Omega(R)$ in the fluid is \\[ \\Omega(R)=a+\\frac{b}{R^2} \\] with \\[ a=\\frac{\\muh-\\etah^2}{1-\\etah^2} \\Omin, \\q b=\\frac{1-\\muh}{1-\\etah^2}\\Rin^2 \\Omin, \\] where $\\muh=\\Omout/\\Omin$. After the Rayleigh stability criterion, i.e. $\\partial_R(R^2\\Omega)^2 > 0$, the flow is hydrodynamically stable for $\\muh>\\etah^2$, i.e.  $\\muh>0.25$ for our geometry. In order to isolate MRI we are mainly interested in hydrodynamically stable regimes so that  $\\muh>\\etah^2$ should be  fulfilled. ",
        "conclusions": "We have shown that, as predicted, additional toroidal fields reduce the critical Reynolds number for the MHD Taylor-Couette flows for laboratory liquids with $\\Pm \\approx 10^{-5}-10^{-6}$. The nonlinear, axisymmetric simulations in the quasi-static approximation yield very similar characteristic values compared to those obtained with linear theory for infinite cylinders.  The velocity amplitudes which were not known until now, are small but still it is not very difficult to measure them experimentally. Similarly the velocity of the drift can easily be detected by looking at changes in the velocity field. We shall also mention that increasing Reynolds numbers result in linear (at least to some point) increase of velocity amplitudes, therefore it is easier to detect them for higher $\\Rey$, however for too high rotation rates the assumption of axisymmetry may not be valid.  For finite aspect ratio calculations, the velocities induced by Ekman pumping are of similar order as those resulting from the instability. Therefore one should not neglect the influence of the endplates in such experiments. Nevertheless the influence of the toroidal field and the traveling wave can easily be detected in all cases."
    },
    "0608/astro-ph0608229_arXiv.txt": {
        "abstract": "We present observations of the stripped Virgo Cluster spiral NGC~4522, a clear, nearby example of a galaxy currently undergoing ISM-ICM stripping. Utilizing SparsePak integral field spectroscopy on the WIYN 3.5m telescope and GALEX UV photometry, we present an analysis of the outer disk ($r > 3$ kpc) stellar population of this galaxy, beyond the HI and H$\\alpha$ truncation radius. We find that the star formation in the gas-stripped outer disk ceased very recently, $\\sim 100$ Myr ago, in agreement with previous claims that this galaxy is currently being stripped. At the time of this stripping, data and models suggest that the galaxy experienced a modest starburst. The stripping is occurring in a region of the cluster well outside the cluster core, likely because this galaxy is experiencing extreme conditions from a dynamic ICM due to an ongoing sub-cluster merger. The outer disk has a spectrum of a K+A galaxy, traditionally observed in high-redshift cluster galaxies.  In the case of NGC~4522, a K+A spectrum is formed by simple stripping of the interstellar gas by the hot intracluster medium.  These data show K+A spectra can be created by cluster processes and that these processes likely extend beyond the cluster core. ",
        "introduction": "K+A galaxies \\citep{dg83}, originally discovered in clusters at high redshift, differ from typical elliptical, spiral or irregular galaxies in an important way. These galaxies have spectra characterized by strong Balmer absorption lines and no significant emission from ongoing star formation. While the formation mechanisms of K+A galaxies are still unknown \\citep{tran03,poggianti04a,bekki05,goto05}, these spectra are generally interpreted as the signature of recent cessation of star formation (however, c.f.  \\citealp{burstein05}). \\citet{shioya02} show through star formation models that such cessation and passive evolution of the already-existing population can form a K+A spectrum.  While the interpretation of these spectra as the result of temporally truncated star formation is largely accepted, the cause of the termination of star formation is not settled.  Such catastrophic interruption of star formation seems likely to be caused by external processes. Events such as mergers, galaxy harassment, and ram pressure stripping all have the potential to terminate star formation. While K+A galaxies have been chiefly observed in clusters, observations of K+A spectra in lower density environments (i.e. \\citealp{zabludoff96,tran04,goto05}) show that they are not unique to the cluster environment. What role does environment play in the creation of these galaxies?  In the Virgo cluster, many spirals have truncated H$\\alpha$ disks \\citep{koopmann04}, suggesting that ram pressure stripping has cut off the star formation in the outer disk. These galaxies have normal star formation inside a truncation radius, with no detected star formation at larger radii. The truncation location varies from galaxy to galaxy, with some galaxies moderately truncated inside $0.8 R_{25}$ and others more severely truncated inside $0.4 R_{25}$. In these galaxies, it is possible to measure age-sensitive absorption lines from the stellar populations beyond the truncation radius without contamination of emission lines from star-forming regions.  The ages of the youngest stellar population in the outer disk tells us how much time has elapsed since the last epoch of star formation and, therefore, the time elapsed since the gas was removed. This, combined with the galaxy's location in the cluster, can tell us where in the cluster galaxies are stripped of their star-forming gas.  NGC~4522 is one of the clearest examples of ongoing ISM-ICM stripping observed in the nearby Virgo cluster. It is a highly inclined, 0.5L$_*$ Sc galaxy $3\\fdg3$ ($1.3 r_{200}$) south of M87, in a region of the cluster with modest ICM density (as measured from X-ray emission from \\citealp{bohringer94}).  There are several lines of evidence that NGC~4522 is experiencing ongoing ICM pressure. Despite an undisturbed stellar disk, the HI and H$\\alpha$ emission are spatially truncated in the disk at $r$=3 kpc=0.4 $R_{25}$ and significant amounts of extraplanar HI and H$\\alpha$ emission exist on one side of the disk close to the truncation radius \\citep{kk99,kenney04}.  A ridge of highly polarized radio continuum emission is observed in the disk opposite the extraplanar gas, suggesting compression of the magnetic field lines at the leading edge of the ISM-ICM interaction \\citep{vollmer04}. These data are consistent with a galaxy currently undergoing ram pressure; comparison of simulations and observations \\citep{vollmer06} suggest that we are observing this galaxy at a time close to peak pressure.  \\begin{figure*} \\includegraphics[width=7in]{f1.eps} \\caption{Left: $R$ band image of NGC~4522, with the position of the SparsePak fibers superimposed over the image. Middle: H$\\alpha$ image of NGC~4522, with HI contours \\citep{kenney04} and the SparsePak footprint overlayed. In both figures, the outlines of the fibers used to form the composite spectrum are bold. Right: Optical spectrum, resulting from the combination of several fibers just beyond the gas truncation radius.} \\label{images} \\end{figure*}  \\newpage ",
        "conclusions": "Our results clearly show that the cessation of star formation was recent and, therefore, that the outer disk of NGC~4522 was recently stripped of its star-forming gas. We would like to understand how rapidly star formation stops during an interaction with the ICM and how long dense molecular star-forming clouds might survive after the low-density atomic gas is stripped. Simulations \\citep{vollmer06} compared with the HI observations \\citep{kenney04} suggest that we are observing the galaxy $\\sim 50$ Myrs after peak pressure.  This timescale is consistent with that determined from the stellar population, suggesting that neutral gas stripping and termination of star formation nearly coincide.  \\subsection{Stripping Outside the Cluster Core}  From a simple timescale argument, the stellar populations of NGC~4522 tell us that this galaxy must be stripped locally, 0.9~Mpc ($1.3 r_{200}$) from M87, as opposed to closer to the cluster core. If we assume that the galaxy's motion in the plane of the sky is a factor of 2 higher than the average radial velocity of a Virgo Cluster galaxy (i.e.  we assume $v_{\\textrm{\\scriptsize t}} = 1400$ km/s), it would still take the galaxy $\\sim 500$ Myr to travel the $\\sim 0.75$ Mpc that separates its current location from the region where the ICM density is high enough to strip the outer disk.  \\emph{The observation that star formation ceased $\\sim 100$ Myr ago therefore rules out stripping in the core and argues that stripping must be occurring locally.} This result is in agreement with simulations of the HI morphology and kinematics \\citep{vollmer06} which also argue for local stripping.  Stripping so far out in the cluster shows either that gas stripping of galaxies in clusters is easier than the simple estimates of ram pressure imply \\citep{kenney04}, or that some galaxies experience extreme conditions due to a dynamic ICM \\citep{dupke06,heinz06}. In particular, NGC~4522 appears to be located between M87 and the M49 subcluster, which is likely in the process of merging with Virgo \\citep{schindler99,shibata01}.  Moreover, \\citet{shibata01} observe a shock front in the ICM near the location of NGC~4522, suggesting that the local ICM has a high velocity. This makes it plausible that gas stripping is important for transforming galaxies well outside the cluster core, including at the interface between merging subclusters.  \\subsection{Similarity to Higher-z K+A Galaxies}  The outer disk of NGC~4522 has strong Balmer lines and no emission, similar to K+A galaxies observed at higher redshift. While the outer disk of NGC~4522 shows a K+A spectrum, a global spectrum of this galaxy does \\emph{not} show the same signature, since it contains emission lines from the star-forming inner disk and absorption lines mostly from the stripped outer stellar disk \\citep{gavazzi04}. This makes it different from the K+A galaxies observed at higher redshift and more similar to the ``e(a)'' class \\citep{dressler99}. Spectral modeling \\citep{shioya02} has suggested that the e(a) class may evolve into K+A galaxies, but some observations (i.e. \\citealp{balogh05}) indicate that e(a) galaxies and K+A galaxies are morphologically distinct.  There may be more than one way to make a K+A and e(a) galaxies, but the present observations show that ram pressure stripping \\emph{can} create a K+A or e(a) spectrum, either from total stripping (K+A) or partial stripping (e(a)). This shows that ram pressure is effective in rapidly terminating star formation and is capable of creating the post star formation spectra observed at higher redshift. It may be that, in the case of Virgo, the ICM is not dense enough to completely strip the gas from a massive galaxy in a single passage so that a central star-forming gas disk will remain after this stripping event ends.  In the largest clusters, the ICM density is higher by a factor of $\\sim 10$ and galaxy velocities are higher by a factor of $\\sim 2$; this means that ram pressure is $\\sim 40$ times higher in these clusters.  In such clusters, it is likely that a galaxy like NGC~4522 would be completely stripped of its gas.  In fact, \\citet{poggianti04} have observed a striking correlation between K+A galaxies and the locations of X-Ray structure in the Coma Cluster, suggesting that we are observing such transformations there.  Both these observations and our observation of NGC~4522 indicate that cluster processes are important for the creation of some K+A spectra. While K+A/e(a) galaxies may not all be created in clusters, it appears that at least one route for their formation is the termination of star formation by ISM-ICM stripping.  Moreover, it appears that stripping can drive this transformation well outside the cluster core.  \\vspace{2mm}  We gratefully acknowledge the indispensable discussions and advice of Jim Rose, which were critical at the early stages of this paper. We also would like to thank Pieter van Dokkum, Jacqueline van Gorkom, and Bernd Vollmer for comments that served to clarify the discussion. Finally, we wish to thank the referee for insightful comments and suggestions. This research is supported by NSF grant AST 00-71251."
    },
    "0608/astro-ph0608190_arXiv.txt": {
        "abstract": "The discovery of luminous quasars at redshift $z \\sim 6$ indicates the presence of supermassive black holes (SMBHs) of mass $\\sim 10^9\\, \\Msun$ when the Universe was less than one billion years old. This finding presents several challenges for theoretical models, because whether such massive objects can form so early in the $\\Lambda$-cold dark matter ($\\Lambda$CDM) cosmology, the leading theory for cosmic structure formation, is an open question.  Furthermore, whether the formation process requires exotic physics such as super-Eddington accretion remains undecided. Here, we present the first multi-scale simulations that, together with a self-regulated model for the SMBH growth, produce a luminous quasar at $z \\sim 6.5$ in the $\\Lambda$CDM paradigm. We follow the hierarchical assembly history of the most massive halo in a $\\sim 3\\, \\Gpc^{3}$ volume, and find that this halo of $\\sim 8\\times 10^{12}\\, \\Msun$ forming at $z \\sim 6.5$ after several major mergers is able to reproduce a number of observed properties of SDSS J1148+5251, the most distant quasar detected at $z =6.42$ \\citep{Fan2003}. Moreover, the SMBHs grow through gas accretion below the Eddington limit in a self-regulated manner owing to feedback. We find that the progenitors experience vigorous star formation (up to $10^4\\, \\Msun\\, \\yr^{-1}$) preceding the major quasar phase such that the stellar mass of the quasar host reaches $10^{12} \\, \\Msun$ at $z \\sim 6.5$, consistent with observations of significant metal enrichment in SDSS J1148+5251.  The merger remnant thus obeys similar $\\mbulge$ scaling relation observed locally as a consequence of coeval growth and evolution of the SMBH and its host galaxy. Our results provide a viable formation mechanism for $z \\sim 6$ quasars in the standard $\\Lambda$CDM cosmology, and demonstrate a common, merger-driven origin for the rarest quasars and the fundamental $\\mbulge$ correlation in a hierarchical Universe. ",
        "introduction": "Quasars rank among the most luminous objects in the Universe and are believed to be powered by SMBHs (e.g., \\citealt{Salpeter1964, Lynden-Bell1969}). They constrain the formation and evolution of galaxies and SMBHs throughout cosmic time. The similarity between the cosmic star formation history (e.g., \\citealt{Madau1996, Bunker2004, Bouwens2004}) and the evolution of quasar abundances (e.g., \\citealt{Shaver1996}) suggests an intriguing link between galaxy formation and black hole growth. This is strengthened by tight correlations measured locally between the masses of the black holes and the global properties of the spheroid components of their hosts, such as their luminosities and masses (\\citealt{Magorrian1998, Marconi2003}), light concentration (\\citealt{Graham2001}), and velocity dispersions (\\citealt{Ferrarese2000, Gebhardt2000, Tremaine2002}).  Distant, highly luminous quasars are important cosmological probes for studying the first generation of galaxies, the star formation history and metal enrichment in the early Universe, the growth of the first supermassive black holes (SMBHs), the role of feedback from quasars and black holes in galaxy evolution, and the epoch of reionization. The Sloan Digital Sky Survey (SDSS, \\citealt{York2000}) has contributed significantly to the discovery of high redshift quasars.  Currently, there are over 1000 quasars known at $z \\gtrsim 4$ and 12 at $z \\gtrsim 6$ (\\citealt{Fan2001, Fan2003, Fan2004, Fan2006A}). As reviewed by \\cite{Fan2006B}, quasars at $z \\sim 6$ are characterized by: (1) a low space density ($\\sim 10^{-9}\\, \\Mpc^{-3}$ comoving); (2) high luminosities (absolute luminosity at rest-frame $M_{1450\\AA} < -26$), believed to be powered by SMBHs of $\\sim 10^9\\, \\Msun$; (3) Gunn-Peterson absorption troughs \\citep{Gunn1965} in their spectra, which place these quasars at the end of the epoch of reionization (e.g., \\citealt{Fan2001, Becker2001, Djorgovski2001, Lidz2002, Songaila2002, White2003, Sokasian2003}); and (4) a lack of evolution in the spectral energy distribution compared to lower-shift counterparts (e.g., \\citealt{Elvis1994, Glikman2006, Richards2006}), which demonstrates the existence of ``mature'' quasars at early times and comparable metal enrichment in quasars at all cosmic epochs.  The most distant quasar known, SDSS J1148+5251 (hereafter J1148+5251), was discovered by SDSS at z = 6.42 \\citep{Fan2003}. It is extremely bright optically with $M_{1450\\AA}$ = -27.8, and deep imaging surveys in both optical and radio (\\citealt{Carilli2004, White2005, Willott2005}) show no sign of gravitational lensing or other companions at the same redshift in the vicinity. Over the past few years, this quasar has been extensively studied at many wavelengths. Near-infrared observations by \\cite{Willott2003} and \\cite{Barth2003} imply a bolometric luminosity of $\\Lbol \\sim 10^{14}\\, \\Lsun$ powered by accretion onto a SMBH of mass 1--5 $\\times 10^9 \\, \\Msun$. Radio observations by \\cite{Bertoldi2003A} and \\cite{Carilli2004} suggest that the host is a hyper-luminous far-infrared (FIR) galaxy, with $\\Lfir \\sim 10^{13}\\, \\Lsun$, and these authors estimate a star formation rate of $\\sim 3\\times 10^3\\, \\Msun\\, \\yr^{-1}$ by assuming most of the FIR luminosity comes from young stars.  Emission from carbon monoxide (CO) has been detected (\\citealt{Walter2003, Bertoldi2003B, Walter2004}) corresponding to a mass of $\\sim 2\\times 10^{10} \\, \\Msun$.  Dust has been seen by several groups (e.g., \\citealt{Robson2004, Bertoldi2003A, Carilli2004, Beelen2006}) with an estimated mass of $\\sim 5\\times 10^8\\, \\Msun$.  In particular, {\\em Spitzer} observations by \\cite{Charmandaris2004} and \\cite{Hines2006} indicate that the dust is heated by an active galactic nucleus (AGN). Furthermore, the detection of iron by \\cite{Barth2003}, the carbon fine structure line [CII] by \\cite{Maiolino2005} and excess OI absorption by \\cite{Becker2006} indicate near-solar metallicity in the quasar host.  These observations raise many fundamental questions for models of quasar and galaxy formation: Where do such high-redshift, luminous quasars originate? How do they form? What are the mechanisms and physical conditions for SMBH growth? And, do these quasar hosts obey the same SMBH--host correlations as observed in the local Universe?  Interpretations of various observations of \\zquasar\\ have painted a rather conflicting picture for the formation site of the quasar halo and the SMBH--host relationship. The low abundance of these quasars leads to the view that they are hosted by massive halos ($\\gtrsim 10^{13}\\, \\Msun$) in the rarest density peaks of the dark matter distribution \\citep{Fan2003}. However, it has been argued, based on the lack of companion galaxies in the field, that this quasar may reside in a much lower mass halo in a less overdense region (\\citealt{Carilli2004, Willott2005}). Moreover, \\cite{Walter2004} suggest, based on the dynamical mass estimate from CO measurements, that \\zquasar\\ contains a small stellar spheroid, and that the SMBH may have largely formed before the host galaxy. However, the detection of metal lines (\\citealt{Walter2004, Barth2003, Maiolino2005}), along with dust (\\citealt{Bertoldi2003A, Carilli2004, Robson2004, Charmandaris2004, Hines2006, Beelen2006}), indicates that the interstellar medium (ISM) of \\zquasar\\ was significantly enriched by heavy elements produced through massive star formation at rates of $\\sim 10^3\\, \\Msun\\, \\yr^{-1}$ occurring as early as $z \\gtrsim 10$, and that large stellar bulges form before accreting SMBHs undergo luminous quasar activity.  In an expanding Universe that is dominated by cold dark matter and is accelerated by dark energy, the $\\Lambda$-cold dark matter ($\\Lambda$CDM) cosmology, the leading theoretical model for structure formation, assumes that structure grows from weak density fluctuations amplified by gravity, with small objects collapsing first and subsequently merging to form progressively more massive ones, a process known as ``hierarchical assembly'' (e.g., see \\citealt{Barkana2001} for a review). The formation of galaxies and quasars is therefore determined by the abundance of dark matter halos; i.e., the number density of halos as a function of mass and redshift. The most widely used analytic model for the halo mass function was first developed by \\cite{Press1974} (hereafter PS), which is based on Gaussian density perturbations, linear gravitational growth, and spherical collapse of dark matter. Using the PS mass functions, \\cite{Efstathiou1988} studied the abundance of rare objects, such as luminous quasars at high redshifts. These authors predicted a sharp ``cut-off'' of the quasar population at $z \\sim 5$. However, while the initial, linear growth of density perturbations can be calculated analytically, the gravitational collapse and subsequent hierarchical build-up of structure is a highly nonlinear process that can be followed only through numerical simulations. It has been shown by previous numerical studies (e.g., \\citealt{Jenkins2001, Sheth2002, Springel2003b}), and more recently by the state-of-the-art {\\em Millennium Simulation} by \\cite{Springel2005A}, that the PS formula underestimates the abundance of high-mass halos by up to an order of magnitude. Therefore, whether or not rare quasars such as \\zquasar\\ can form in the $\\Lambda$CDM cosmology remains an open question and an important test of the theory.  To date, only a limited number of analytical or semi-analytic models have addressed the early formation of a $10^9 \\, \\Msun$ SMBH at $z \\sim 6$ (\\citealt{Haiman2001, Haiman2004, Yoo2004, Volonteri2005}). These approaches use merger trees of dark matter halos generated using the PS theory, and assume a black hole accretion rate at or above the Eddington limit. However, as discussed above, the PS--based approach may be inaccurate. Moreover, it is not clear whether sufficient physical conditions for such large accretion rates exist in quasar systems as the hydrodynamics of the large-scale gas flow and feedback from black holes have not been incorporated in earlier modeling.  It is believed that the growth of SMBHs is closely linked to galaxy formation (e.g., \\citealt{Magorrian1998, Ferrarese2000, Gebhardt2000, Graham2001, Tremaine2002, Marconi2003, Haiman2004, Kazantzidis2005, Li2006A}), and that the growth is self-regulated by feedback (e.g., \\citealt{Silk1998, Haehnelt1998A, Fabian1999, King2003, Wyithe2003, DiMatteo2005, Springel2005B, Sazonov2005, Murray2005, Wyithe2005}). Remarkably, self-regulated models with SMBH feedback in the form of thermal energy coupled to the ambient gas have been demonstrated to successfully reproduce many observations of galaxies, including the $\\msigma$ relation (\\citealt{DiMatteo2005, Robertson2006A}), galaxy colors (\\citealt{Springel2005C, Hopkins2006F}), X-ray gas emission \\citep{Cox2006a}, elliptical kinematics \\citep{Cox2006b} and the fundamental plane \\citep{Robertson2006B}, quasar properties (\\citealt{Hopkins2005A, Hopkins2005B}), luminosity functions (\\citealt{Hopkins2005C, Hopkins2005D, Hopkins2006B}), and populations (\\citealt{Hopkins2006A, Hopkins2006C, Hopkins2006D}), and the luminosity function of low-level AGN \\citep{Hopkins2006E}. Furthermore, these simulations of binary mergers identify a plausible, merger-driven formation mechanism for massive black holes and luminous quasars (e.g., \\citealt{DiMatteo2005, Hopkins2006A, Robertson2006A}).  Here, we present a model that accounts for the SMBH growth, quasar activity and host galaxy properties of the most distant quasar known. In our scenario, the quasar and its host galaxy form in a  massive halo that originates from a rare density peak in the standard $\\Lambda$CDM paradigm, and they grow hierarchically through multiple gas-rich mergers, supporting an average star formation rate of $\\sim 10^3\\, \\Msun\\, \\yr^{-1}$ that peaks at $z\\sim 8.5$. Once the progenitors undergo sufficient dynamical friction to coalesce, the multiple SMBHs from the progenitor galaxies merge and exponentially increase their mass and feedback energy via accretion. At $z \\approx 6.5$ when the SMBH mass exceeds $10^{9}\\, \\Msun$, black hole accretion drives a sufficiently energetic wind to clear obscuring material from the central regions of the system and powers an optically luminous quasar similar to \\zquasar. We have devised a set of novel multi-scale simulations, which include cosmological N-body calculations on large scales and hydrodynamic simulations of galaxy mergers on galactic scales, coupled with the self-regulated growth of SMBHs, enabling us to follow galaxy assembly and quasar formation at $z \\sim 6$.  This paper is organized as follows. In \\S~2, we describe our computational methods and models, which includes a set of large scale cosmological N-body simulations, and hydrodynamical galaxy mergers along the merging history of the quasar halo. In \\S~3, we present the formation and evolution of the quasar and its host galaxy, including the assembly of the galaxy progenitors, star formation, and SMBH growth, as well as the SMBH--host correlations, and properties of the quasar such as luminosities and lifetimes. We discuss feedback from starburst-driven winds, quasar abundances for cosmological models with different parameters, the implication of black hole mergers, and galaxies in the epoch of reionization in \\S~4, and summarize in \\S~5. ",
        "conclusions": "\\subsection{Comparison with Previous Models and Robustness of Our Results} Our multi-scale simulations that include large-scale cosmological N-body calculations and hydrodynamic simulations of galaxy mergers, as well as a self-regulated model for black hole growth, have successfully produced a luminous quasar at $z \\sim 6.5$ with a black hole mass of $\\sim 2 \\times10^{9}\\, h^{-1}\\, \\Msun$ and a number of properties similar to those of \\zquasar, the most distant quasar detected at $z=6.42$ \\citep{Fan2003}. Our approach differs from previous semi-analytic studies by \\cite{Haiman2001, Haiman2004, Yoo2004, Volonteri2005} and \\cite{Volonteri2006} in the following ways:  \\begin{enumerate}  \\item We use a realistic merger tree derived directly from multi-scale, high-resolution cosmological simulations. The previous studies used merger trees of dark matter halos generated with the extended Press-Schechter theory \\citep{Press1974, Lacey1993}, which may underestimate the abundance of high-mass halos by up to one order of magnitude, as shown in \\S~\\ref{subsec_mf}. Also, the merger trees in those studies started from much higher redshifts than what we consider here. In our model, the quasar halo is the largest one in a volume of $1\\, h^{-3}\\, \\Gpc^{3}$. It has a mass of $\\sim 7.7 \\times 10^{12}\\, \\Msun$ at $z \\sim 6.5$ built up through seven major mergers from $z \\simeq 14.4$ to $z \\simeq 6.5$.  \\item We follow the evolution of the system and treat the gas dynamics, star formation, and black hole growth properly. This approach is critical to investigation of the properties of both black holes and host galaxies, and their evolution (e.g., \\citealt{DiMatteo2005, Springel2005B, Robertson2006A, Hopkins2006A}), but it was not included in those previous studies on formation of $z \\sim 6$ quasars.  \\item We employ a self-regulated model for the growth of supermassive black holes, in which the accretion is regulated by the black hole feedback, and the rate is under the Eddington limit. In the previous studies, the black hole growth was unregulated, but instead a constant- or super-Eddington accretion rate was used.  \\item In our simulations, we do not consider black hole ejection caused by gravitational recoil owing to insufficient resolution and lack of relativistic physics. However, the halo escape velocities in our simulations are in the range of $486$ -- $1284\\, \\rm{km\\, s}^{-1}$, much larger than the kick velocity $\\sim 100$ -- $475\\, \\rm{km\\, s}^{-1}$ (e.g., \\citealt{Herrmann2006, Baker2006, Gonzalez2006, Herrmann2007}, see \\S~\\ref{subsec_Vesc} and \\S~\\ref{subsec_bhgrowth} for more details). Therefore, black hole ejection may be negligible in our case. Previous studies had much smaller halo progenitors at higher redshifts than ours, so the black hole seeds would be more likely subject to ejection from their halos. This leads to the conclusion in these studies that constant or super-Eddington accretion is needed owing to significant black hole ejection. Our results are robust within the best estimates currently available for the recoil velocity of the black hole binary.  \\item The black hole seeds in the galaxy progenitors in our simulations are massive (e.g., $\\sim 10^5\\, \\Msun$ at $z \\sim 14$). The sub-resolution recipe in our model does not allow us to resolve the actual formation and accretion of such black holes below this mass scale. The formation of these seeds is an unsolved problem, but our results do not depend on the specific prescription of the formation process. We adopt a picture in which the seed holes come from the remnants of the first stars (which have a mass $200\\, \\Msun$ at $z = 30$) and grow under Eddington limit until they enter the merger tree we simulated. If the growth is delayed by radiation feedback from the PopIII stars (e.g., \\citealt{Johnson2006}), then super-Eddington accretion, or other proposed scenarios (e.g., \\citealt{Bromm2003, Begelman2006}) may be necessary to form massive seeds of $\\sim 10^{5}\\, \\Msun$ in the protogalaxies. \\end{enumerate}  Overall, we conclude that the results from our simulations, which are more realistic and more detailed than the models previously done, are robust. Suppermassive black holes of $\\sim 10^9\\, \\Msun$ can form rapidly through gas-rich hierarchical mergers under Eddington limit, even within a short period of time. We find that constant or super-Eddington accretion is not necessary unless the above assumptions in our modeling break, i.e., there are no massive black hole seeds of $10^5\\, \\Msun$ available at $z \\sim 14$, or the recoil velocity of the coalescing black hole binary is extremely high (e.g., $> 1000\\, \\rm{km\\, s}^{-1}$). Under these extreme circumstances, some ``exotic'' processes such as super-Eddington accretion may be necessary to grow a $\\sim 10^9\\, \\Msun$ SMBH within a few hundred million years. However, we should note, as pointed out by \\cite{Bogdanovic2007}, that most gas-rich galaxy mergers have a configuration such that the orbit and spins of both black holes are aligned with the large-scale gas flow owing to torques from accreting gas. Such a configuration has a maximum kick velocity $< 200\\, \\rm{km\\, s}^{-1}$, which is well below the escape velocity of a $10^{10}\\, \\Msun$ dwarf galaxy, as well as those of the halos in our modeling.  \\subsection{Merging History of Black Holes} \\label{subsec_bhbinary}  During the galaxy mergers, the black holes follow their host halos to the system center and can form binaries (or multiple systems). The coalescence of a black hole binary includes three distinct phases: inspiral, merger, and ringdown (e.g., \\citealt{Flanagan1998}). Whether black hole binaries can coalesce on short timescales is a matter of debate. In a stellar environment, it has been argued that a binary hardens very slowly owing to an eventual depletion of stars that cause the binary to lose angular momentum (e.g., \\citealt{Begelman1980, Milos2003}). In a gaseous environment, however, numerical simulations by \\cite{Escala2004} and \\cite{Li2007C} show that the binaries decay rapidly owing to strong dynamical friction with the gas, and they likely merge within $10^7$ years. Because our galaxies are very gas rich and have large central concentrations of gas during the mergers, we assume the black hole particles coalesce once their separation decreases below our spatial resolution ($30\\, h^{-1}\\, \\rm{pc}$) and their relative speed falls below the local gas sound speed \\citep{Springel2005B}.  In the simulations, we do not have sufficient resolution nor the relativistic physics to consider the ejection of black holes by gravitational recoil during the merger phase. However, as discussed in \\S~\\ref{subsec_Vesc} and \\S~\\ref{subsec_bhgrowth}, the halo escape velocities in our simulations are much larger than the maximum kick velocity for black hole binary estimated from the latest relativistic calculations. So black hole ejection is likely unimportant in our modeling. To accurately address gravitational recoil in the galaxy merger simulations, we need to include general relativity, resolve the dynamics of black hole binaries with extremely high resolution, and calculate the halo potential in a cosmological context (in which halo potential distribution may be different from that of a single object). However, such a comprehensive treatment is impossible at the moment. We therefore assume the black holes merge quickly once they reach the stage of gravitational radiation.  These coalescing supermassive black holes will be strong sources of gravitational radiation detectable by the Laser Interferometer Space Antenna ({\\em LISA}, \\citealt{Folkner1998}), as suggested by many authors (e.g., \\citealt{Thorne1976, Haehnelt1998B, Flanagan1998, Menou2001, Hughes2002, Sesana2005, Koushiappas2006}). By tracing the merging history of the SMBHs, {\\em LISA} could shed light on the distribution, structures and evolution of the associated dark matter halos. Because luminous, high-redshift quasars are likely sites of vigorous hierarchical mergers, they may be the best targets for {\\em LISA} to explore the early Universe.  \\subsection{Feedback from Starburst-driven Winds}  \\begin{figure} \\begin{center} \\includegraphics[width=3.5in]{f14.ps} \\vspace{0.5cm} \\caption{Evolution of star formation and black hole growth in merger simulations with a starburst-driven wind model. The simulation is run with lower resolution ($\\rm {N_{tot}} \\sim 5\\times 10^5$), and the specifications of the wind model is: wind efficiency $\\eta=0.5$, wind energy coefficient $\\chi=0.25$, wind free travel length $\\rm {L_{w} = 20\\, \\kpc}$, and a wind velocity $\\rm {V_{w} = 418\\, km\\, s^{-1}}$.} \\label{FigWind} \\end{center} \\end{figure}  Vigorous star formation would induce a galactic wind and mass outflow, a phenomenon that has been observed to prevail in both local star-forming galaxies as indicated by blueshifted optical absorption lines (e.g., \\citealt{Martin1999, Heckman2000, Martin2005, Rupke2002, Rupke2005}), and Lyman Break Galaxies at $z \\sim 3$ as indicated by blueshifted interstellar absorption lines and redshifted Ly$\\alpha$ emission lines (e.g., \\citealt{Pettini2002, Shapley2003}), as well as Ly$\\alpha$ emitters at $z \\sim 5.7$ \\citep{Ajiki2002}. These galactic winds are generally thought to play a significant role in galaxy evolution (e.g., see \\citealt{Veilleux2005} and \\citealt{Cox2006C} for recent reviews).  The strong starburst preceding the major quasar phase in our simulations may drive strong galactic winds and affect the black hole growth. To investigate the impact of the feedback from a starburst-driven wind on the growth of the black hole, we have done the same merger simulation with lower resolution ($\\rm {N_{tot}} \\sim 5\\times 10^5$) and with a canonical wind model from \\cite{Springel2003b}: the wind efficiency $\\eta=0.5$, which measures the coefficient of the star formation that determines the mass outflow; the energy fraction from supernovae injected into the wind $\\chi=0.25$, wind free travel length $\\rm {L_{w} = 20\\, \\kpc}$, and a wind velocity $\\rm {V_{w} = 418\\, km\\, s^{-1}}$. As demonstrated by \\cite{Cox2006C}, this wind model is able to reproduce the starbursts as observed in Lyman Break Galaxies, and therefore is suitable to our study.  We find that the impact of the starburst-driven wind on the quasar evolution is minor, as shown in Figure~\\ref{FigWind}. Both the histories of star formation and black hole growth remain roughly the same as in the simulation without a starburst wind, only the amplitude is lowered by a factor of $\\sim 1.5$. Similarly, the final masses of the black hole and the stars are reduced by roughly the same factor, but the quasar host is still on the $\\mbulge$ correlation. The peak quasar phase is delayed to $z \\sim 6$. Overall, the starburst wind affects the gas dynamics locally, but owing to the deep potential of the system, its impact on the process of quasar formation is minor. Our results support the finding by \\cite{Cox2006C} that feedback from starburst-driven winds alone is ineffective at regulating the growth of the central black hole, so feedback from the black hole plays the dominant role in the formation and evolution of quasars.  \\subsection{Abundance and Fate of Quasar Halos at $z \\sim 6$}  \\begin{figure} \\begin{center} \\includegraphics[width=3.5in]{f15.ps} \\caption{Comparison of halo abundances at $z \\sim 6$ from the zoom-in simulations with parameters from both WMAP1 (solid line) and WMAP3 (dashed line). The volume of the high-resolution zoom-in region is $\\sim 50^3 \\, h^{-3}\\, \\Mpc^3$.} \\label{FigNum} \\end{center} \\end{figure}  Because we have so far simulated only one quasar in a volume of $1 \\, h^{-3}\\, \\Gpc^{3}$, we are not yet able to constrain the expected abundance of quasars $z\\sim 6$. As mentioned in \\S~2.3, at a given redshift, cosmological simulations with parameters from WMAP1 produce more massive halos than runs with WMAP3 owing to a larger value of $\\sigma_8$. Figure~\\ref{FigNum} shows the number of halos at $z\\sim 6$ from the zoom-in runs with parameters from both WMAP1 and WMAP3. There are about three dozen halos with mass $M > 10^{12}\\, \\Msun$ in the WMAP1 run, while in the WMAP3 run there are only a handful of such halos. However, since in our picture the quasar activity depends not only on the halo mass, but also on the merging history, an accurate estimate of the quasar abundance and luminosity function would require hydrodynamical simulations of all the quasar candidates in a large box, which are currently unavailable. Nevertheless, all conditions being equal, the change from the WMAP3 parameters would produce fewer luminous quasars at $z \\sim 6$. This suggests that in a WMAP3 cosmology, the quasar observed with the largest redshift, \\zquasar, might have formed in a slightly higher overdensity peak than that we have presented here. In that event, if the WMAP3 determination of $\\sigma_8$ were correct, we would need to identify a rarer density fluctuation to match \\zquasar\\ at its observed redshift. However, this does not change our conclusion that the most distant and luminous quasars can form from hierarchical galaxy mergers in the $\\Lambda$CDM cosmology.  Imaging surveys of \\zquasar\\ show that there is no other luminous quasar from the same epoch in the field (\\citealt{Carilli2004, White2005, Willott2005}). In our simulations, around the peak of quasar activity at $z\\approx 6.5$, there are no other halos of mass $> 10^{12}\\, \\Msun$ within a few Mpc of this quasar. However, the numerous major mergers this halo experienced prior to the peak quasar activity demonstrates that this region was once highly clustered with massive halos, but they merged to become a bigger one by $z = 6.5$.  As seen from Figure~\\ref{FigTree}, this quasar halo will undergo a handful of major mergers at a later time from $z\\sim$~4--1, and eventually end up as a cD-like galaxy at the center of a rich cluster. Since we do not follow hydrodynamically the evolution of the quasar at $z<4$, the physical conditions of these mergers remain undetermined. It is not clear whether this halo would experience more episodes of starburst or quasar activity later on during these mergers. Therefore, the final black hole mass and other properties of this quasar at the present day are deferred to future simulations that follow its evolution to $z=0$.   \\subsection{Galaxies in the Epoch of Reionization}  The epoch of reionization (EoR) is an important landmark event in cosmic history that constrains the formation of the first luminous objects \\citep{Loeb2001}. The recent results of WMAP3 indicate that the Universe was 50\\% reionized at $z \\approx 9.3$ \\citep{Page2006, Spergel2006}, while studies of Gunn-Peterson absorption \\citep{Gunn1965} suggest that reionization began as early as $z \\sim 14$ and ended at $z \\sim 6$ \\citep{Fan2006A}. At present, it is believed that the reionization sources are star-forming galaxies since there are insufficient quasars at $z>6$ as indicated by the steep quasar luminosity function \\citep{Fan2006A}.  The galaxy progenitors of the quasar in our simulations underwent extreme and prolonged starbursts before $z \\sim 6.5$. Less extreme galaxies in this epoch may also have vigorous star formation histories. Detecting these galaxies and determining their contribution to reionization will be crucial to understanding the EoR (\\citealt{Hernquist2003, Barton2004, Dave2006}). As reviewed by \\cite{Hu2006}, recent observations using both broad-band colors (e.g., \\citealt{Dickinson2004, Yan2005, Bunker2004, Bouwens2004, Giavalisco2004, Egami2005, Eyles2005, Mobasher2005, Yan2006, Eyles2006}) and narrow-band Ly$\\alpha$ emission (e.g., \\citealt{Hu2002, Malhotra2004, Stern2005}) have detected $\\sim 500$ galaxies at $z \\sim 6$ and a handful at $z \\gtrsim 7$ \\citep{Bouwens2005}. The low luminosity density of galaxies currently detected at $z>7$ seems insufficient to reionize the Universe. However, ongoing surveys with the HST and Spitzer telescopes, and future missions such as the {\\em Dark Ages z Ly$\\alpha$ Explorer} (DAzLE, \\citealt{Horton2004}) and the {\\em James Webb Space Telescope} (JWST, \\citealt{Windhorst2006}) will search deeper and further for high-redshift objects, and may eventually unveil ionizing sources in the EoR."
    },
    "0608/astro-ph0608473_arXiv.txt": {
        "abstract": "We present the latest results from a multi-epoch timing and spectral study of the Transient Anomalous X-ray Pulsar XTE~J1810--197. We have acquired seven observations of this pulsar with the Newton X-ray Multi-mirror Mission (\\xmm) over the course of two and a half years, to follow the spectral evolution as the source fades from outburst. The spectrum is arguably best characterized by a two-temperature blackbody whose luminosities are decreasing exponentially with $\\tau_1 = 870$~d and $\\tau_2 = 280$~d, respectively. The temperatures of these components are currently cooling at a rate of 22\\% per year from a nearly constant value recorded at earlier epochs of $kT_1 = 0.25$~keV and $kT_2 = 0.67$~keV, respectively.  The new data show that the temperature $T_1$ and luminosity of that component have nearly returned to their historic quiescent levels and that its pulsed fraction, which has steadily decreased with time, is now consistent with the previous lack of detected pulsations in quiescence. We also summarize the detections of radio emission from XTE~J1810--197, the first confirmed for any AXP. We consider possible models for the emission geometry and mechanisms of XTE~J1810--197. ",
        "introduction": "Neutron star (NS) astronomy has been recently invigorated by the identification of a new class of magnetically dominated emitters. Known as anomalous X-ray pulsars (AXPs) and soft gamma-ray repeaters (SGRs), these objects are apparently young, isolated neutron stars (NSs), whose properties differ markedly from those of the Crab pulsar, previously considered prototypical of the young NSs (for a review see \\citealt{mer02}). All AXPs and at least 3 of the 4 SGRs are identified as relatively slow ($5-12$~s) pulsars. Many are located at the centers of recognized, young supernova remnants (SNRs), directly associating them with their supernova explosions. These objects emit predominantly at X-ray energies and are distinguished by their characteristic spectral signature. This radiation cannot be accounted for by rotational energy losses alone, as for the radio pulsars, but is most likely powered by the decay of an enormous magnetic field characterized by a dipole with $B_p \\simgt 4.4\\times 10^{13}$~G, at the pole.  Collectively, these isolated NSs are understood within the context of the magnetar theory (\\citealt{dun92,tho96}).  A unique pulsar has been discovered whose remarkable properties offer great promise for deciphering the emission mechanism(s) of magnetars. \\taxp\\ is a 5.54~s X-ray pulsar whose measured and derived physical parameters are fully characteristic of an AXP, but expresses behavior not previously associated with any such object \\citep[hereafter Paper~I]{got04}. \\taxp\\ is a transient AXP (TAXP) -- it was discovered during a bright impulsive outburst \\citep{ibr04} that is still fading steadily. Even more surprising is the discovery of highly variable radio emission, providing the first confirmed example of radio flux from an AXP \\citep{hal05b}, and the subsequent detection of radio pulsations at the X-ray period \\citep{cam06}.  The outburst that resulted in the detection of the transient AXP occurred sometime between 2002 November and 2003 January \\citep{ibr04}. Since then, over the course of a year, regular scans of the region with \\xte\\ recorded an exponential flux decay with a time-constant of $269\\pm25$~days from a maximum of $F (2-10\\ {\\rm keV}) \\approx 8 \\times 10^{-11}$ erg~cm$^{-2}$~s$^{-1}$. In comparison, the previous average quiescent flux, with its softer spectrum, gives $F(0.5-10\\ {\\rm keV}) \\approx 5.5 \\times 10^{-13}$ erg~cm$^{-2}$~s$^{-1}$ (Paper~I).  This contrast is unprecedented for an AXP.  Further \\xte\\ observations yielded SGR-like episodes of bursts \\citep{woo05}, similar to those seen from 1E~2259+586 \\citep{kas03}. A search for an optical/IR counterpart detected a fading IR source within the \\chandra\\ error circle, similar to ones associated with other AXPs, confirming its identification with \\taxp\\ \\citep{rea04a,rea04b}.  In this paper we present the latest results from TAXP \\taxp\\ including new \\xmm\\ observations extending to 2006 March. These observations allow us to characterize the spectral evolution of a TAXP in outburst as it returns to quiescence.  We show that the spectrum has now begun a marked transition back to the nominal quiescent state, with a distinct temperature and flux evolution. This long-term evolution provides the unique possibility of decomposing its fading spectral components in a manner unavailable for other magnetars (\\citealt[hereafter Paper~II]{hal05}). In the following, we used a revised distant to the NS, $d_{3.3}$, quoted in units of $3.3$~kpc and based on the radio pulsar DM measurement of \\cite{cam06}. ",
        "conclusions": ""
    },
    "0608/astro-ph0608159_arXiv.txt": {
        "abstract": "The disruption of binary stars by the tidal field of the black hole in the Galactic Center can produce the hypervelocity stars observed in the halo.  We use numerical models to simulate the full spectrum of observable velocities of stars ejected into the halo by this binary disruption process. Our model includes a range of parameters for binaries with 3-4 M$_\\odot$ primaries, consideration of radial orbits of the ejected stars through an approximate mass distribution for the Galaxy, and the impact of stellar lifetimes. We calculate the spectrum of ejection velocities and reproduce previous results for the mean ejection velocity at the Galactic center. The model predicts that the full population of ejected stars includes both the hypervelocity stars with velocities large enough to escape from the Galaxy and a comparable number of ejected, but bound, stars of the same stellar type. The predicted median speeds of the population of ejected stars as a function of distance in the halo are consistent with current observations. Combining the model with the data also shows that interesting constraints on the properties of binaries in the Galactic Center and on the mass distribution in the Galaxy can be obtained even with modest samples of ejected stars. ",
        "introduction": "Hypervelocity stars (HVSs) are objects moving with speeds sufficient to escape the gravitational influence of the Galaxy. Hills \\citeyearpar{hil88} predicted their existence, and Brown et al.\\ \\citeyearpar{broetal05} discovered the first HVS in a survey of blue stars within the Galactic halo. We now know of seven HVSs \\citep{ede05,hir05,broetal06a,broetal06b}.  One HVS is a subdwarf O star at a distance of $\\sim$ 20 kpc from the \\GC\\ \\citep{hir05}. The other HVSs are probably B-type main sequence stars with galactocentric distances of 50--100 kpc. The five late B stars in this group have masses $m \\sim$ 3-4 \\msun\\ \\citep{broetal06a,broetal06b}. The early B-type HVS is more massive, $m \\sim$ 8 \\msun\\ \\citep{ede05}.  Only two proposed models plausibly explain the origin of the observed population of Galactic HVSs. In the single black hole model, a binary system strays too close to \\Sgr, the 3.5 million solar-mass black hole in the \\GC\\ \\citep{hil88,hil92,yutre03,gouqui03}.  This encounter tears the binary apart, leading to the capture of one star and the high-speed ejection of its partner \\citep{hil88}.  For typical ejection velocities of $\\sim$ 2000 km s$^{-1}$, the HVS can reach galactocentric distances of $\\sim$ 50 kpc in $\\sim$ 25 Myr \\citep{hil88,hil92,yutre03,gouqui03,broetal06b,perhopale06}, roughly the lifetime of a 9 \\msun\\ star \\citep{schaetal92, schaetal93}.  Other HVS models involve three-body interactions between a single star and a binary black hole. Random encounters between stars near the \\GC\\ and a binary black hole can lead to an HVS with a typical ejection velocity of $\\sim$ 2000 km s$^{-1}$ \\citep{qui96,yu02,yutre03,gua05,ses06}. Alternatively, the in-spiral of an intermediate mass black hole with $m \\sim 10^3-10^4$ \\msun\\ can eject stars from the \\GC\\ \\citep[e.g.,][]{han03}.  Some of these encounters will produce HVSs with ejection velocities of $\\sim$ 2000 km s$^{-1}$ \\citep{lev05,BauGuaPor06,ses06}. Although binary black holes with circular orbits and in-spiraling black holes produce roughly isotropic distributions of HVSs on the sky, eccentric binary black holes can produce anisotropic space distributions of HVSs \\citep{hol05,BauGuaPor06,ses06}.  Here, we use numerical simulations to predict the distribution of observable radial velocities of HVSs ejected by a single black hole at the \\GC. Our results identify a new population of ejected stars of the same stellar type as the HVSs but more slowly moving.  Ongoing halo surveys \\citep{broetal06b} show evidence of this population.  The distribution of observable radial velocities among these stars may yield information about the nature of the progenitor binaries and can serve as a probe of the distribution of mass throughout the Galaxy. Even the small number of known HVSs can constrain the distribution of mass in the Milky Way \\citep[see also][]{gneetal05}.  We describe the basic model and the numerical simulations in \\S2, compare the numerical results with observations in \\S3, and conclude with a discussion and summary in \\S4. ",
        "conclusions": "We have developed a first calculation of the predicted distribution of radial velocities for HVSs at galactocentric distances of 10--120~kpc in the Galactic halo.  Our simulations quantify the ejection probabilities and velocities of binaries interacting with \\Sgr, generate the distribution of ejected stars from the vicinity of the \\GC, and produce the full phase-space radial velocity distribution of ejected stars in the halo from an integration of the orbits of ejected stars through the Galactic potential. In addition to providing a framework for interpreting the observed velocity distribution of HVSs and related stars, the results of these simulations (Figures \\ref{fig:vdisp}--\\ref{fig:fountainb}) demonstrate that this approach can provide interesting constraints on the Galactic potential and on the mass function of the population of binary stars in the \\GC.  Besides confirming Hills' (1988, 1992) results for the average velocity of ejected stars from the black hole, we derive a broad distribution of ejection velocities for each value of $\\abin$ and $\\rclose$.  A Gaussian with $\\sigma_v \\sim 0.2 v_{avg}$ provides a reasonable approximation to the derived velocity distribution.  We plan to consider the origin of this distribution in future studies.  The initial comparisons between the model and observations are encouraging.  The measured radial velocities of HVSs and stars with somewhat lower speeds agree well with the predictions. The current population of these stars provides a constraint on the Galactic mass distribution. As ongoing radial velocity surveys yield more HVSs with accurate stellar masses, the model will allow more robust tests for the origin of HVSs and lower velocity stars, will provide limits on the initial mass function of primary stars in the \\GC, and will yield better constraints on mass models for the Galaxy.  As surveys reveal more HVS, the observed velocity distribution of ejected stars in the Galactic halo may help to distinguish among alternative scenarios for producing HVSs. Compared to the single black hole models, binary black hole models produce relatively fewer ejections at the highest velocities \\citep{yutre03,BauGuaPor06}. Although binary black hole models may yield larger ejection rates overall \\citep[e.g.,][]{yutre03}, the two models produce similar velocity distributions at lower velocities.  Thus, better constraints on the HVS distribution at the highest velocities may provide a test of the two pictures for HVS formation.  Refinements to the mass models can include more detailed Galactic structure, in particular the Galactic disk.  Indeed, \\citet{gneetal05} have already proposed that HVSs can serve to probe halo triaxiality. Including the disk potential should lead to nonradial ejection velocities, and may cause some clustering of ejected stars in the halo. We may therefore use the breakdown of spherical symmetry in the distribution of ejected stars about the \\GC\\ to reveal the Galaxy's structure.  If ejected stars are produced in short-lived star forming regions near \\Sgr, or if they originate in individual star clusters which settle dynamically into the \\GC, ejected stars --- including HVSs --- may also cluster in time. Given a Galactic mass model, we can identify any temporal correlations by using the radial position and velocity of stars to determine the time since they were ejected.  The extraction of temporal and spatial clustering information may be difficult with rare HVSs. However, a clustering analysis might be feasible with the more numerous, slower population of ejected stars suggested by our model."
    },
    "0608/astro-ph0608645_arXiv.txt": {
        "abstract": "We describe observations of 9.7 {$\\mu$}m silicate features in 97 AGNs, exhibiting a  wide range of AGN  types and of  X-ray extinction toward the central nuclei.  We find that the strength of the silicate feature correlates with the HI column density estimated from fitting the X-ray data, such that  low HI columns correspond to  silicate emission while high  columns  correspond to  silicate  absorption.   The behavior  is generally consistent with unification models where the large diversity in AGN properties is  caused by viewing-angle-dependent obscuration of the nucleus.   Radio-loud AGNs and radio-quiet  quasars follow roughly the correlation  between HI columns  and the strength of  the silicate feature defined  by Seyfert galaxies.   The agreement among  AGN types suggests a high-level unification with similar characteristics for the structure of the obscuring  material.  We demonstrate the implications for  unification models  qualitatively with  a conceptual  disk model. The model includes  an inner accretion disk ($<$ 0.1  pc in radius), a middle disk (0.1-10  pc in radius) with a  dense diffuse component and with embedded  denser clouds, and an  outer clumpy disk  (10-300 pc in radius). ",
        "introduction": "The arrangement of material  around the central supermassive blackhole (SMBH)  in active  galactic nuclei  (AGNs) is  crucial  in unification schemes, where the  active nuclei classified as type  2 arise from the obscuration  of  broad  optical  emission  lines  in  type  1  objects \\citep{Antonucci93}.  The structure  of the circumnuclear material may be critical to  the growth of the SMBH and  may influence the feedback from the  AGN and its  effects on the  formation and evolution  of the host galaxies.   The infrared (IR) spectral  energy distribution (SED) arises from the heating of  surrounding material, and its shape should indicate  how the  material is  organized.  There  are at  least three different  possibilities:  compact  \\citep{Pier92,  Pier93},  extended \\citep{Granato94,     Granato97}    and     cloudy    \\citep{Krolik86, Rowan-Robinson95, Nenkova02}  structure models  all fit the  data well (assuming an extra component of  star formation in the compact model). The silicate emission and  absorption features now being observed with {\\it Spitzer} (e.g.,  Weedman et al. 2005) can be  used to probe these possibilities further.  In   this   paper,   we   present  $Spitzer$   Infrared   Spectrograph \\citep[IRS;][]{Houck04} observations  of 9.7 $\\mu$m  silicate features in    93    AGNs,    including    examples   from    the    literature \\citep{Siebenmorgen05,  Hao05, Sturm05,  Ogle06}.  We  complement them with  groundbased   measurements  for  4   more  \\citep{Roche91}.   By combining these observations with X-ray  data that explore the line of sight toward  the central  engine, we can  gain new insights  into the unification scheme. ",
        "conclusions": "We report  observations of 9.7{$\\mu$}m  silicate features in  97 AGNs. The  features  vary  from  emission  to  absorption  with  increasing HI  column density,  consistent with  unification models. Radio-loud AGN  (radio-loud QSOs and  FR II galaxies)  and radio-quiet QSOs  (PG  and 2MASS  QSOs)  lie  roughly  on the  Seyfert-correlation between HI column and  silicate feature strength.  The behaviors of LINERs and BALQs are not clear  due to the small number of objects for these two AGN types.  The scatter in the relation is large and several Compton-thick   AGNs   do   not   show   deep   silicate   absorption. Qualitatively, the correlation  requires a circumnuclear disk geometry with an  accretion disk outside  of which is  a middle disk  with high density and  with even denser  clouds embedded (0.1-10 pc  in radius), co-aligned  with an  outer clumpy  disk  (10-300 pc  in radius).   The similarity of the behavior of  various types of AGN suggests that this disk geometry may be typical for AGN in general. \\\\   We thank  the anonymous referee  for detailed comments. We  also thank Roberto Maiolino  for helpful suggestions.  Support for  this work was provided by  NASA through contract  1255094 issued by  JPL/ California Institute of Technology."
    },
    "0608/hep-th0608178_arXiv.txt": {
        "abstract": "The Gauss-Bonnet (GB) curvature invariant coupled to a scalar field $\\phi$ can lead to an exit from a scaling matter-dominated epoch to a late-time accelerated expansion, which is attractive to alleviate the coincident problem of dark energy. We derive the condition for the existence of cosmological scaling solutions in the presence of the GB coupling for a general scalar-field Lagrangian density $p(\\phi, X)$, where $X=-(1/2)(\\nabla \\phi)^2$ is a kinematic term of the scalar field. The GB coupling and the Lagrangian density are restricted to be in the form $f(\\phi) \\propto e^{\\lambda \\phi}$ and $p=Xg (Xe^{\\lambda \\phi})$, respectively, where $\\lambda$ is a constant and $g$ is an arbitrary function. We also derive fixed points for such a scaling Lagrangian with a GB coupling $f(\\phi) \\propto e^{\\mu \\phi}$ and clarify the conditions under which the scaling matter era is followed by a de-Sitter solution which can appear in the presence of the GB coupling. Among scaling models proposed in the current literature, we find that the models which allow such a cosmological evolution are an ordinary scalar field with an exponential potential and a tachyon field with an inverse square potential, although the latter requires a coupling between dark energy and dark matter. ",
        "introduction": "A great deal of efforts in modern cosmology, over the course of the past decade, have been made in trying to understand the nature and the origin of dark energy which is responsible for late-time acceleration of the universe \\cite{obser,review,CST}. Although cosmological constant is the simplest candidate of dark energy consistent with current observations, it suffers from a severe fine-tuning problem from a field theoretic point of view \\cite{Weinberg}. It is then tempting to find alternative models of dark energy which have dynamical nature unlike cosmological constant. An accelerated expansion can be realized by using a scalar field whose origin may be found in superstring or supergravity theories. Thus modern particle physics has been an attractive supplier for the construction of scalar-field dark energy models--such as quintessence \\cite{quin}, K-essence \\cite{Kes}, tachyon \\cite{tachyon}, dilatonic models \\cite{Gas}, and ghost condensate \\cite{ghostcon,PT04}.  Although the dynamically evolving field models have an edge over the cosmological constant scenario, they too, in general, involve fine-tuning of initial values of the field and the parameters of the model. Secondly, nucleosynthesis places stringent restriction on any additional degree of freedom (over and above the particle physics standard model degrees of freedom) which translates into a constraint on the ratio of the field energy density to the energy density of the background fluid \\cite{CST}. These problems can be addressed by employing scalar field models exhibiting {\\it scaling solutions} \\cite{Fer,CLW,scaling,TS04,Tsuji06,AQTW}. For instance, a standard scalar field $\\phi$ with an exponential potential  in General Relativity (GR) \\cite{Fer,CLW} can mimic the background (radiation/matter) remaining sub-dominant so as to respect the nucleosynthesis constraint \\cite{Nelson}. The scaling solutions as dynamical attractors can successfully address the second problem and can also considerably alleviate the fine-tuning problem of initial conditions.  For a viable cosmological evolution, the scalar field should remain unimportant for most of the cosmic evolution and should emerge from the hiding epoch only at late times to give rise to the accelerated expansion. Since scaling solutions are non-accelerating, the model should be supplemented by an additional mechanism allowing the field to exit from the scaling regime at late times. There are several ways of achieving this goal: \\begin{itemize} \\item (i) At late times the potential becomes shallow to lead to the accelerated expansion. For example this is realized by the double exponential potential given by $V(\\phi)=V_{0} (e^{-\\lambda \\kappa \\phi}+e^{-\\mu \\kappa \\phi})$ with $\\lambda^2>20$ and $\\mu^2<2$ \\cite{Nelson} (here $\\kappa^2=8\\pi G$ with $G$ being gravitational constant).  \\item (ii) A local minimum of the field potential appears at late times. The examples of this type of models are $V(\\phi)=V_{0} [\\cosh (\\lambda \\kappa \\phi)-1]^n$ \\cite{Sahni} and $V(\\phi)=V_{0} e^{-\\lambda \\kappa \\phi}[A+(\\kappa \\phi -B)^2]$ \\cite{Alb}. In the former case the acceleration occurs for $n<1/2$ on average, whereas in the latter case the system always approaches a de-Sitter universe after a damped oscillation of the field.  \\item (iii) Assisted  quintessence \\cite{KLT}. This is realized by considering many fields $\\phi_i$ with the sum of the exponential potential $V=\\sum A_i e^{-\\lambda_i \\kappa \\phi_i}$, since the presence of many scalar fields leads to the smaller effective coupling, i.e., $\\lambda_{\\rm eff}^{-2}= \\sum \\lambda_{i}^{-2}$ \\cite{Liddle}. \\end{itemize}  Although these are interesting attempts to provide a successful exit from the scaling regime, it is fair to say that the aforementioned mechanisms are still phenomenological. It is of great importance to find cosmological models from fundamental theories of particle physics like string theory which may incorporate the findings of the phenomenological studies. String theory necessarily includes scalar fields (dilaton $\\&$ modulus fields); the low energy string effective action also gives rise to a modification of standard GR in terms of Riemann invariants coupled to scalar field(s) \\cite{reviewstring}. The modifications are important to address the problems of past \\cite{past} and future \\cite{future} singularities. Out of these corrections, the Gauss-Bonnet (GB) term is of a special relevance. The GB term is a topological invariant quantity which contributes to the dynamics in 4 dimensions provided it is coupled to a dynamically evolving scalar field. It is the unique invariant for which the highest (second) derivative occurs linearly in the equations of motion thereby ensuring the uniqueness of solutions. Nojiri {\\it et al.} \\cite{NOS} studied scalar-Gauss-Bonnet cosmology to demonstrate the existence of a particular dark energy solution. Such a solution exists in the presence of an exponential potential $V(\\phi) \\propto e^{-\\lambda \\kappa \\phi}$ and the GB coupling $f(\\phi) \\propto e^{\\mu \\kappa \\phi}$. More recently Koivisto and Mota \\cite{KM06} showed that the GB coupling allows a possibility to lead to the exit from the scaling matter era to the dark energy dominated universe for the same model. See Refs.~\\cite{CTS,Neupane,NO05,Cog06,Cal,Sami06,Annalen,Sanyal,Andrew} for other aspects of the GB dark energy scenario.  In this paper we shall clarify the form of the GB coupling $f(\\phi)$ for the existence of scaling solutions by starting from a very general scalar-field Lagrangian density $p(\\phi, X)$ where $X$ is a kinematic term of the field. This includes a wide variety of scalar-field models such as quintessence, K-essence, tachyon, and ghost condensate. The demand for scaling solutions not only determines the field Lagrangian to be $p(\\phi,X)=Xg(Xe^{\\lambda \\phi})$ ($g$ is an arbitrary function) but also fixes the GB coupling in the form $f(\\phi) \\propto e^{\\lambda \\phi}$ (here we use the unit $\\kappa^2=1$). Thus our results are useful to construct viable dark energy models which possess scaling solutions in the matter-dominated epoch. We then investigate the fixed points for the scaling Lagrangian $p(\\phi,X)=Xg(Xe^{\\lambda \\phi})$ with a GB coupling $f(\\phi) \\propto e^{\\mu \\phi}$. We show the existence of a pure de-Sitter fixed point which exists only for specific models and study a possibility to obtain a scaling matter era that finally approaches the de-Sitter solution. A standard field with an exponential potential is a specific model to realize such a cosmological evolution provided that $\\mu>\\lambda$. A tachyon field also possesses the de-Sitter fixed point, but it does not have a scaling matter epoch. However the scaling matter era followed by the de-Sitter solution is possible if dark energy couples to dark matter. We shall discuss cosmological dynamics in such cases in details. ",
        "conclusions": "In this paper we have discussed the viability of Gauss-Bonnet (GB) dark energy models which possess cosmological scaling solutions. In order to alleviate the coincident problem of dark energy, it is of interest to find a cosmological trajectory which is in a scaling regime ($\\Omega_{\\phi}/\\Omega_{m}=$constant) during a matter era and finally exits to give rise to a late-time accelerated expansion. We have tried to understand the property of scalar-field dark energy models which allow such a possibility.  First we derived the form of the GB coupling together with the form of the scalar-field Lagrangian for the existence of scaling solutions by starting from a very general action given in Eq.~(\\ref{eqn:action}). The form of the GB coupling has been found to be $f(\\phi) \\propto e^{\\lambda \\phi}$ together with the field Lagrangian density $p=Xg(Xe^{\\lambda \\phi})$, where $\\lambda$ is defined in Eq.~(\\ref{lamdef}) and $g$ is an arbitrary function. The field Lagrangian density takes the same form as in the case where the GB coupling is absent.  We have also derived autonomous equations and fixed points for the scaling Lagrangian $p=Xg(Xe^{\\lambda \\phi})$ with the GB coupling $f(\\phi) \\propto e^{\\mu \\phi}$. When $\\lambda=\\mu$, we obtained the scaling solution characterised by $x_1=\\sqrt{6}/[2(Q+\\lambda)]$ and the scalar-field dominated solution characterised by $\\Omega_{\\phi}+\\Omega_{\\GB}=1$ without specifying any form of the Lagrangian. In this case, however, the solutions do not get away from the scaling matter regime to connect to a dark energy era. When $\\mu>\\lambda$ the presence of the GB term can lead to the exit from the scaling regime.  The GB term allows an interesting possibility to have a late-time de-Sitter fixed point. This point is present if the solutions for $gx_1^2=0$ and $x_1=0$ exist. The models which possess this de-Sitter solution are characterised by the form $g(Y)=(c_0-c/Y)f(Y)$, where $f(Y)$ is a function that approaches a constant as $x_1  \\to 0$. The ordinary field with an exponential potential (\\ref{mo1}) and the tachyon field with an inverse square potential (\\ref{mo3}) belong to this class, whereas the dilatonic ghost condensate model (\\ref{mo2}) does not.  We have studied cosmological evolution for the above two models in which the de-Sitter point is present. In the case of an ordinary field with an exponential potential one can in fact obtain a viable trajectory along which the scaling matter period is followed by the de-Sitter solution provided that $\\mu>\\lambda$ and $\\lambda^2>3$ for $Q=0$. There is also an interesting possibility in which the effective equation of state $w_{\\rm eff}$ temporarily becomes smaller than $-1$ because of the growth of the GB term. This is realized by choosing larger values of $\\mu$, but it also corresponds to a shorter matter-dominated period. The late-time acceleration is possible even for $\\lambda=\\mu$ by using a scalar-field \\& GB dominated point ($\\Omega_{\\phi}+ \\Omega_{\\GB}=1$), but in this case the contribution of the GB term is very small at the present epoch. Moreover we do not have a scaling matter era in this case, which is not a welcome feature in solving the coincidence problem.  The tachyon field has a de-Sitter fixed point, but in this case the scaling solution is absent for a non-relativistic background fluid ($w_{m}=0$) if $Q=0$. Hence we do not have a successful sequence of the scaling matter era and the de-Sitter expansion. If the coupling $Q$ is present, however, it is possible to realize this sequence when $\\mu$ is larger than $\\lambda$. The values of $Q$ are bounded from both above and below to get a scaling matter era with $w_{\\rm eff}$ close to 0. An alternative approach to the problem discussed here can be provided by Renormalization Group (RG) equations. The scaling solution presented here turns out to be the fixed point of RG flow equations; the form of the field Lagrangian and the GB coupling are automatically fixed in this method \\cite{sanjay}.  We also discussed ghost conditions for the GB dark energy models by considering gravitational perturbations. For the ordinary field with an exponential potential we found that the propagation speeds of scalar and tensor modes can be imaginary during the transition from the scaling matter era to the final de-Sitter attractor. This is associated with the appearance of ghosts which lead to negative instabilities of perturbations. This ghost stage can be avoided by tuning the model parameters at the expense of obtaining a minimum effective equation of state closer to $-1$.  In this paper we have not studied local gravity constraints on the GB coupling \\cite{ACD,ACD2}. If the contribution of the GB coupling is large at present epoch, this may contradict with local gravity experiments unless an accidental cancellation occurs to avoid the large variation of effective gravitational constant. We leave future work for the construction of scaling GB dark energy models consistent with such constraints. We hope that the results obtained in this paper will be useful to find such viable models."
    },
    "0608/astro-ph0608409_arXiv.txt": {
        "abstract": "We present time series data on the variable stars of the galactic globular cluster Messier~3 (M3). We give {\\it BVI$_{\\rm C}$\\/} light curves for 226 RR Lyrae, 2 SX Phe and 1 W Vir type variables, along with estimated fundamental photometric parameters such as intensity and magnitude-averaged brightness and pulsation periods. In some cases the periods we have found significantly differ from the previously published ones. This is the first published light curve and period determination for variable V266. The {\\it I\\/}-band light curve has not been observed previously for numerous (76) variables. Three new RR Lyrae variables have been discovered. Groups of RR Lyrae variables that belong to different evolutionary stages and have been separated previously on the basis of {\\it V\\/} data were found here for all colours and colour indices by cluster analysis. The {\\it I\\/}-band period \\--- luminosity relation is also discussed. From the 66 modulated (Blazhko type) RR Lyrae stars we investigated, six are newly identified and two of them are first overtone pulsators. In the case of 13 RR Lyrae, the period of Blazhko cycle has been estimated for the first time.  V252 is identified as a new RRd variable. Amplitude ratio of RRd stars have been investigated to search possible mode content changes. In contrast to previous publications no changes have been found. Problems with the sampling of the time series of typical cluster variability surveys is demonstrated. ",
        "introduction": "Multicolour photometry of RR Lyrae stars is very important for obtaining information on their physical parameters such as temperature, mass and luminosity. RR Lyrae stars in globular clusters deserve additional interest because they offer an excellent test for the stellar evolution theory and cosmological distance scale.  Regular monitoring of clusters' variables is also essential to investigate various subtle effects, such as period and modal content changes of multimode RR Lyrae stars, or the Blazhko phenomenon.  M3 (NGC~5272; $\\alpha_{2000}=13^{\\rm h}42^{\\rm m}11\\fs 2$, $\\delta_{2000}=28\\degr 22\\arcmin 32\\arcsec $) has been target of a large number of different studies. Although it is the most variable-rich galactic globular cluster (with 274 catalogued variables, with most of these belonging to the RR Lyrae type), the number of high quality time series observations are still limited. \\citet{CC01} (hereafter CC01) published a comparable amount of data in {\\it B\\/} and {\\it V\\/} colours to the dataset presented here. In addition, only three CCD-observations based publications exist with partial coverage of M3: \\citet{K98} (K98) in {\\it V}-band, \\citet{C98} (C98) and \\citet{hartman} (H05) in {\\it BVI}. \\citet*{Strader} (S02) derived {\\it V\\/} time series for stars in the core of the cluster, but their light curves represent differential fluxes on an arbitrary scale, instead of standard magnitudes.  The lack of data on M3 inspired several time series studies; such as \\citet*{Mallik}, \\citet{Ug} and \\citet{AFer}. The results of these investigations have not been published yet.  We started a multicolour {\\it BVI$_{\\rm C}$\\/} CCD survey of globular clusters at Konkoly Observatory in 1998.  The main goal of the project is to obtain high quality time series for as many globular cluster variables as possible. The instrumentation enables us to observe the clusters simultaneously with two telescopes (1-m~RCC and 90/60-cm Schmidt), both equipped with CCD cameras. Their different fields of view (FOVs) make it possible to obtain images with sufficient resolution for accurate photometry of both the crowded central regions and the dispersed halos, respectively.  In this work we present the observations of M3.  \\begin{figure} \\includegraphics[width=8.5cm]{frames} \\caption[]{The relative sizes and positions of the frames used in the published CCD time series papers on M3. The notations are: `K' \\--- \\citet{K98}, `C' \\--- \\citet{C98}, \\--- `H' \\citet{hartman}, \\---`RCC' + `SCHMIDT' \\--- present work. \\citet{CC01} have used a slightly larger FOV than `SCHMIDT'.} \\label{frames} \\end{figure} ",
        "conclusions": "We have conducted the most extensive CCD {\\it BVI\\/} survey of variables in the globular cluster M3 covering 238 stars. Image subtraction photometric data transformed into the Johnson-Cousins magnitude scale, light curves, periods and average magnitudes were presented  \\begin{itemize} \\item {\\it Varibility search\\/} In addition to our new variable stars  published in \\citet{Acta}, we have discovered three further RR Lyrae stars. We also found four non-variable stars in the instability strip.  Period, brightness were improved for most of the stars and several confused identifications were clarified.  \\item {\\it Color \\--- magnitude diagrams\\/} At least three significant subgroups have been separated by hierarchical clustering method in the CMDs of the RRab stars.  The identified three clusters agree well with previous studies of \\citet{letter} and with CC05 when we contract the appropriate clusters to 2 groups from our three ones or 3 from the four subgroups separated by \\citet{letter}.  \\item {\\it Period \\--- brightness relations\\/} A clear correlation have been found between period of RRab stars and their $\\langle I \\rangle$ magnitudes. The relation is in good agreement with synthetic horizontal branch calculation published by \\citet{Catelan}.  \\item {\\it Modulated RR Lyrae stars\\/} We identified 66 modulated (Blazhko) RR Lyrae stars in our sample where five of them are new discoveries. For six stars the ambiguous Blazhko classification have been verified. Two modulated RRc variable stars were found in the cluster. By combining the available CCD data modulation periods were estimated for 13 Blazhko star for the first time.  \\item {\\it Double mode RR Lyraes\\/} Investigating the Fourier spectra of V252 its double-mode pulsation has been confirmed.  Analyzing all available photometric data on RRd stars of the cluster no undoubted modal content change, neither on shorter nor longer time scales have been detected except for the case of V79. It was shown, that to detect a mode-switching it is practically impossible from available CCD photometric materials containing observing runs from 1-3 seasons. \\end{itemize}  At the end of this paper we would like to call the attention to the necessity of further long CCD time series observation of the variable stars of M3. To answer the question of modal changes of the double-mode variable stars or to find correct period for all modulated stars and many similar problems need long-term, homogeneous and high precision photometry."
    },
    "0608/astro-ph0608315_arXiv.txt": {
        "abstract": "In this paper we present a correlation between the peak spectral energy of gamma-ray bursts (GRBs) and the peak bolometric luminosity of the underlying supernovae (SNe), based on a sample of four pairs of GRBs-SNe with spectroscopically confirmed connection. Combining it with the well-known relation between the peak spectral energy and the isotropic equivalent energy of GRBs, we obtain an upper limit on the isotropic energy of GRBs, which is $\\approx 10^{52} {\\rm erg} \\left(L_{\\sn,\\p}/10^{43} {\\rm erg\\, s}^{-1} \\right)^{10}$, where $L_{\\sn,\\p}$ is the peak bolometric luminosity of the SNe. Our results suggest that the critical parameter determining the GRB-SN connection is the peak luminosity of SNe, rather than the feature of the SN spectra and/or the SN explosion energy as commonly hypothesized. Since it is generally believed that the peak luminosity of SNe powered by radioactive decays is related to the amount of \\nk\\, produced in the SN explosion, the mass of \\nk\\, may be a key physical factor for understanding the nature of GRBs and their connection with SNe. Application of our relation to Type Ibc SNe with normal peak luminosities indicates that if those normal SNe have GRBs accompanying them, the GRBs would be extremely soft and sub-energetic in gamma-rays, and hence easier to detect with X-ray or UV detectors than with gamma-ray detectors. ",
        "introduction": "\\label{intro}  The discovery of SN 1998bw within the error box of GRB 980425 \\citep{gal98} inspired a lot of consideration on the connection between gamma-ray bursts (GRBs) and supernovae (SNe). Since then, three more pairs of GRBs and SNe with spectroscopically confirmed connection have been found, which are GRB 030329/SN 2003dh \\citep{sta03,hjo03}, GRB 031203/SN 2003lw \\citep{mal04,saz04}, and the most recent one discovered by \\sw, GRB 060218/SN 2006aj \\citep{mas06,mod06,cam06,sol06,pia06,mir06,cob06}. Interestingly, all of the four SNe are among a special class of Type Ic, called the broad-lined SNe, which are characterized by smooth and featureless spectra indicating a very large expansion velocity \\citep[and references therein]{del06,woo06}. Modeling of the SN lightcurves reveals that the SNe with GRB-connection have very large explosion energy and mass production of \\nk\\, compared to normal Type Ibc SNe \\citep{iwa98,nak01,den05,maz06b}, except SN 2006aj which requires an explosion energy that is comparable to that of normal SNe Ibc \\citep{maz06}. These facts have motivated people to invent the term ``hypernovae'' for this special and much more powerful class of SNe (Iwamoto et al. 1998; see also Paczy\\'nski 1998a,b).  A less direct way for identifying the GRB-SN connection is observing the rebrightening and/or flattening (called ``red bumps'') in the late GRB afterglows, which can be interpreted as the emergence of the underlying SN lightcurves \\citep[and references therein]{blo99,zeh04,sod05,bers06}. Although alternative explanations with dust echoes \\citep{esi00} and dust sublimation \\citep{wax00} have been proposed, several groups have successfully fitted SN 1998bw templates to explain the late-time bumps \\citep{blo02,gar03,gre03,sta05,bers06}. A systematic study on the GRB afterglows with this approach by \\citet{zeh04} suggests that all long-duration GRBs are associated with SNe.  Despite the exciting developments in the past eight years in the detection and observation of GRB-SN connection, by now no any quantitative relation between the parameters of GRBs and that of SNe has been found (see, e.g., Zeh, Klose \\& Hartmann 2004; Ferrero et al. 2006), although it is commonly conceived that only very bright SNe can produce GRBs, based on the fact that all SNe with confirmed GRB-connection are much brighter than average and that the rate of GRBs and ``hypernovae'' are several orders of magnitude lower than the rate of core-collapse SNe \\citep{pod04}. The lack of a quantitative relation between GRBs and SNe has frustrated the advance in understanding the nature of GRBs, although many people believe that long-duration GRBs are produced by the core-collapse of massive stars \\citep{mac01,woo06a}.  In this paper, we present the discovery of a quantitative relation between the GRBs and the underlying SNe, based on the observational data of the four pairs of GRBs-SNe with spectroscopically confirmed connection (Table~\\ref{grbsn}). We show that, the peak spectral energy of the GRB is strongly correlated with the peak bolometric luminosity of the SN. Then, combining with the relation between the peak spectral energy and the isotropic equivalent energy of GRBs found by \\citet{ama02}, we explore the implications of the correlation that we have found for the GRB-SN connection and for the nature of GRBs.   \\begin{table*} \\centering \\begin{minipage}{160mm} \\caption{Gamma-ray bursts and supernovae with spectroscopically confirmed connection$^\\dagger$} \\label{grbsn} \\begin{tabular}{llllllll} \\hline GRB/SN\\hspace{1cm} & $z^{\\rm (a)}$\\hspace{0.48cm} & $E_{\\gamma,\\p}^{\\rm (b)}$\\hspace{0.33cm} & $E_{\\gamma,\\iso}^{\\rm (c)}$\\hspace{1.4cm} & $M_{\\sn,\\p}^{\\rm (d)}$\\hspace{0.76cm} & $E_K^{\\rm (e)}$\\hspace{0.84cm} & $M_\\ej^{\\rm (f)}$\\hspace{0.46cm}  & $M_{\\rm Nickel}^{\\rm (g)}$ \\\\ \\hline 980425/1998bw & $0.0085$ & $55\\pm 21$ & $0.00009\\pm 0.00002$ & $-18.65\\pm 0.20$ & $5.0\\pm 0.5$ & $10\\pm 1$ & $0.38$--$0.48$ \\\\ 030329/2003dh & $0.1687$ & $79\\pm 3$  & $1.7\\pm 0.2$ & $-18.79\\pm 0.23$ & $4.0\\pm 1.0$ & $8\\pm 2$ & $0.25$--$0.45$ \\\\ 031203/2003lw & $0.1055$ & $159\\pm 51$ & $0.009\\pm 0.004$ & $-18.92\\pm 0.20$ & $6.0\\pm 1.0$ & $13\\pm 2$ & $0.45$--$0.65$ \\\\ 060218/2006aj & $0.0335$ & $4.9\\pm 0.4$ & $0.0059\\pm 0.0003$ & $-18.16\\pm 0.20$ & $0.2\\pm 0.02$ & $2\\pm 0.2$ & $0.2\\pm 0.04$ \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^\\dagger$] Following \\citet{maz06b}, we assume $H_0 = 72$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m = 0.28$, and $\\Omega_\\Lambda = 0.72$. Quantities calculated in a different cosmology are converted to this cosmology. \\item[$^{\\rm (a)}$] Cosmic redshift. From \\citet{gal98} (GRB 980425/SN 1998bw), \\citet{sta03} (GRB 030329/SN 2003dh), \\citet{pro04} (GRB 031203/SN 2003lw), \\citet{pia06} and \\citet{mir06} (GRB 060218/SN 2006aj). \\item[$^{\\rm (b)}$] Peak energy of the integrated GRB spectrum in units of keV, measured in the GRB frame. From \\citet{yam03} (GRB 980425), \\citet{sak05} (GRB 030329), \\citet{ula05} (GRB 031203), and \\citet{cam06} (GRB 060218). \\item[$^{\\rm (c)}$] Isotropic equivalent energy of the GRB in units of $10^{52}$ erg, defined in the $1-10,000$ keV energy band in the GRB frame. From \\citet{cam06} (GRB 060218), and \\citet{ama06} (other bursts). \\item[$^{\\rm (d)}$] Peak bolometric magnitude of the supernova, defined in the $3,000-24,000$ \\AA\\, wavelength band in the SN frame. From \\citet{maz06b} (SN 1998bw, SN 2003lw), \\citet{pia06} (SN 2006aj), and \\citet{den05} (SN 2003dh, and Appendix~\\ref{sn03dh} of this paper). \\item[$^{\\rm (e)}$] Explosion kinetic energy of the supernova in units of $10^{52}$ erg. From \\citet{maz06} (SN 2006aj, error of 10 percent added), and \\citet{maz06b} (other SNe). \\item[$^{\\rm (f)}$] Ejected mass in the supernova explosion in units of $M_\\odot$. From \\citet{maz06} (SN 2006aj, error of 10 percent added), and \\citet{maz06b} (other SNe). \\item[$^{\\rm (g)}$] Mass of \\nk\\, produced by the SN explosion in units of $M_\\odot$. From \\citet{maz06} (SN 2006aj, the error corresponds to a 0.2 mag error in magnitude), and \\citet{maz06b} (other SNe). \\end{list} \\end{minipage} \\end{table*} ",
        "conclusions": "\\label{concl}  We have found a strong correlation between the peak spectral energy of GRBs and the peak bolometric magnitude (i.e., the peak luminosity) of their underlying SNe, based on the observational data of the four pairs of GRBs and SNe with spectroscopically confirmed connection (Fig.~\\ref{mag_epeak}, eqs.~\\ref{ep_mp} and \\ref{ep_lp}). The Pearson linear correlation coefficient between $\\log E_{\\gamma,\\p}$ (the peak spectral energy of GRBs) and $-M_{\\sn,\\p}$ (the peak bolometric magnitude of SNe) is $0.997$, corresponding to a probability $P = 0.003$ for zero correlation. Although the sample is limited by the small number of GRBs-SNe, we consider the result to be very suggestive because of the large correlation coefficient.  Combined with the relation between the peak spectral energy and the isotropic equivalent energy of GRBs \\citep{ama06}, the correlation that we have found leads to a relation between the isotropic energy of a GRB and the peak bolometric luminosity of the underlying supernova (eq.~\\ref{eiso_up}). If a GRB is among the normal cosmological class (i.e., it has a normal total gamma-ray energy), and is indeed associated with a supernova, then equation~(\\ref{eiso_up}) would take the equal sign. If a GRB is sub-energetic, like some of the SN-connected GRBs, equation~(\\ref{eiso_up}) gives an upper bound on the isotropic gamma-ray energy.  The slope of $\\log E_{\\gamma,\\iso} - \\log L_{\\sn,\\p}$ is extremely steep, which is $\\approx 10$ by equation~(\\ref{eiso_up}). This naturally describes the observational fact that GRBs have a very large diversity in properties (e.g., the isotropic equivalent energy) compared to SNe.  Applying the relations that we have obtained (eqs.~\\ref{ep_lp} and \\ref{eiso_up}) to normal Type Ibc SNe which are not as luminous as SN 1998bw, we found that the prompt emission from the potential GRBs associated with them peaks in the soft X-ray and UV band, and the total gamma-ray energy of the bursts is extremely small. Hence, the bursts associated with normal SNe Ibc would be more appropriately qualified as soft X-ray transients, which might be easier to detect with X-ray or UV detectors than with gamma-ray detectors.  Despite the fact that cosmological GRBs are typically more luminous than the sub-energetic GRB 980425 by five orders of magnitude or more, the potential SNe associated with them are expected to be brighter than SN 1998bw only by a factor $\\sim 2$. Although it is hard to predict in a general case up to what a distance a SN associated with a cosmological GRB can be observed, under favorable conditions a GRB-connected SN that is as luminous as a SN Ia should be observable up to redshift $\\sim 1.7$ with the upcoming {\\it SNAP} space observatory.  Our results suggest that the critical parameter characterizing the GRB-SN connection is the large peak luminosity of the SNe, rather than the broad-lined spectra (or, equivalently, the large expansion velocity) and/or the huge explosion energy as commonly hypothesized \\citep[and references therein]{nom04,del06,woo06}. Given the general Ansatz that the SN luminosity at peak equals the power generated by the decay of \\nk\\, \\citep{arn82,mae03,nom04,maz06,pia06}, our results may indicate that the mass of \\nk\\, produced in the SN explosion is a key physical factor for understanding the nature of the GRB-SN connection as well as the nature of GRBs (Fig.~\\ref{n56_epeak}). Although a physical relation between the peak spectral energy of GRBs and the mass of \\nk\\, of SNe cannot be established based only on the results in this paper, the following consideration may provide us a clue. Popular models of GRBs involve aspherical explosion of massive stars, where gamma-ray emission is produced along the axis of the explosion via a jet or a shock \\citep{woo93,pac98a,mac01,woo06a}. Investigation on nucleosynthesis in aspherical SN explosions indicates that \\nk\\, is distributed also preferentially in the direction along the jet axis where the ejecta carry more kinetic energy and the shock is stronger \\citep{mae02}.  Finally, we remark that a SN has been claimed to be detected in the afterglow of GRB 020903, an extremely soft burst at redshift $0.251$ \\citep{sod05,bers06}. The peak spectral energy of GRB 020903 is $3.37\\pm 1.79$ keV. The isotropic gamma-ray energy is $(2.8\\pm 0.7)\\times 10^{49}$ erg, which makes GRB 020903 consistent with the $E_{\\gamma,\\p}-E_{\\gamma, \\iso}$ relation \\citep{ama06}. Equation~(\\ref{ep_mp}) then predicts that the SN associated with GRB 020903 has a peak bolometric magnitude $\\approx -18.06$, fainter than SN 1998bw by $\\sim 0.6$ mag. Fitting the SN 1998bw template to the bump in the afterglow lightcurve of GRB 020903, \\citet{bers06} found that the SN is fainter than SN 1998bw by $0.8\\pm 0.1$ mag at peak in the $R$-band, consistent with the $0.6\\pm 0.5$ mag reported by \\citet{sod05} earlier. Considering the fact that the spectrum of the SN of GRB 020903 at 38.6 days after the burst was redder than SN 1998bw, the difference in the bolometric magnitudes is likely to be less than 0.8 mag (D. Bersier, private communication). It appears that GRB 020903 and its supernova are consistent with equations~(\\ref{ep_lp}) and (\\ref{eiso_up}), although a conclusion cannot be made without the availability of data in other filters."
    },
    "0608/gr-qc0608022_arXiv.txt": {
        "abstract": "Class of axially symmetric solutions of vacuum Einstein field equations including the Papapetrou solution as particular case has been found. It has been shown that the derived solution describes the external axial symmetric  gravitational field of the source with nonvanishing mass. The general solution is obtained for this class of functions. As an example of physical application, the spacetime metric outside a line gravitomagnetic monopole has been obtained from Papapetrou solution of vacuum equations of gravitational field. ",
        "introduction": "% \\label{intr}  A wide range of massive objects from the compact relativistic stars, as black holes and neutron stars,  up to modern cosmological models, including Early Universe and its evolution~\\cite{HawPen} where the solutions of Einstein equations are applicable and important is currently known. But the logic of scientific development and practical experience requires an introduction into the consideration of principally new objects, like cosmic strings, gravitomagnetic monopoles, quasiclosed worlds etc. The extension of objects under research leads to the additional analysis and extension of well known existing solutions of fundamental field equations. Therefore reconsideration of old and finding new exact exterior solutions of Einstein field equations describing space-time around these objects is one of the interesting and important branches of general relativity and theoretical physics. There exist a lot of interesting problems, especially in quantum gravity, because it is still not clear how to build quantum gravity (see, for example,~\\cite{Dadhich}) and in this level of our understanding it is very useful to investigate objects which are on the boundary between classical and quantum gravity. Of special importance also are external gravitational fields of isolated stars, as today is recognized they are at the core of most intriguing astrophysical phenomena. In spite of the fact that many exact solutions of Einstein equations can be found in classical monographs as~\\cite{Kramer} it is still interesting to study new possible solutions of this fundamental equation.  In our present work we concentrate our efforts in investigation and analysis of some exact solutions of Einstein equations~\\cite{Kramer}, namely, on the exterior axisymmetric fields in the vacuum, arising around rotating gravitational objects. In our consideration we look for some mathematical features and structure of Einstein equations, which could determine their physically acceptable solutions.  In fact, our motivation is, that the Papapetrou  class of axially-symmetric solutions~\\cite{Papa}$^{,}$\\cite{Djamal} of vacuum Einstein equations does not contain physically interesting one, because in the classical Newtonian limit there is no term being proportional to $r^{-1}$, what means that it describes rotation field for the bounded systems with zero mass~\\cite{Djamal}. However it should be noted, that one of the solutions of Einstein equations with the similar asymptotic behavior has been considered and interpreted as gravitational field of massless quadrupole~\\cite{Fomin1}. From geometrical point of view, the gravitational source is characterized  by a topological closure of internal three dimensional space with zero integral contribution according to its zero mass.  In the Papapetrou case the solutions are given in the class of pure harmonic functions $\\zeta$ and our assumption is to investigate this solution in the more extended functional space of functions $\\zeta$.  Here we follow to the logic of investigation, brilliantly reflected in~\\cite{Djamal}. As we will show in the Section~\\ref{ext}, small difference, introduced by us into the approach described in~\\cite{Djamal}, gives us a possibility to obtain solution of Einstein equations for more general class, which includes Papapetrou one as particular case.  Throughout, we use a space-like signature $(+,-,-,-)$ and a system of units in which $G = 1 = c$. Greek indices are taken to run from 0 to 3 and Latin indices from 1 to 3; partial derivatives are denoted with a comma. ",
        "conclusions": "We conclude, that obtained solution~(\\ref{16}) is an extension of the well known Papapetrou class of solutions, which can be applied to the physical systems, presenting rotating bounded nonzero masses. We have constructed the  solution for the particular function $\\eta(z)$, which in some sense, demonstrates the existence of $\\rho^{-1}$ term in asymptotics. In addition using properties of Poisson like equation for $\\zeta$ function, we have found the general solution (\\ref{28}) for any given function $\\eta (z)$, which generates various metric functions $f(\\rho, z)$. The connection of above obtained solution with metric of a line gravitomagnetic monopole was considered as a particular case. This is preliminary result and the physical interpretation of the obtained solution is outside the framework of this paper. More detailed investigation and application to gravitational sources is the scope of our future research. In this respect as the next step we want to investigate the regularity conditions at the axis $\\rho=0$ in dependence from the choice of function $\\eta (z)$."
    },
    "0608/astro-ph0608586_arXiv.txt": {
        "abstract": "A nonlinear kinetic theory of cosmic ray (CR) acceleration in supernova remnants is employed to investigate the properties of the remnant SN~1987A. It is shown that a large downstream magnetic field $B_d\\approx 10$~mG is required to fit the existing observational data. Such a strong field together with the strong shock modification due to CR backreaction provides the steep and concave radioemission spectrum and considerable synchrotron cooling of high energy electrons which diminish their X-ray synchrotron flux below the observed Chandra flux which has to be considered as an upper limit for nonthermal X-ray emission. The expected $\\gamma$-ray energy flux at TeV-energies at the current epoch is $2\\times 10^{-13}$~erg/(cm$^2$s). ",
        "introduction": "Supernova (SN) 1987A occurred in the Large Magellanic Cloud. It has been extensively studied in all wavelengths from radio to $\\gamma$-ray. The initial short outburst of radio emission \\citep{turtle87} is attributed to the synchrotron emission of electrons accelerated by the SN shock propagated in the free wind of presupernova star, which was the blue supergiant (BSG) \\citep{chf87}. After about 3 years radio emission was detected  again \\citep{ss92, gae97} as well as the monotonically increased X-ray emission \\citep{gro94, has96}. This second increase of emission is attributed to the entrance of the outer SN shock into the thermalized BSG wind and then in the H~II region occupied by much more dense matter, consists of the swept up wind of red supergiant (RSG) progenitor star \\citep{chd95}.  To describe the observed properties of nonthermal emission of SN~1987A we use here a nonlinear kinetic theory. It couples the particle acceleration process with the hydrodynamics of the thermal gas \\citep{byk96, bv00}. Therefore in spherically symmetric approach it is able to quite definitely predict the evolution of gas density, pressure, mass velocity, together with the energy spectrum and the spatial distribution of cosmic ray (CR) nuclei and electrons at any given evolutionary epoch $t$, and the properties of the nonthermal radiation produced in SNRs due to these accelerated CRs. The application of this theory to individual SNRs \\citep[see][for a review]{ber05} has provided the explanation of the observed SNR properties under the assumption of strong interior magnetic field. It was also shown that this strong field leads to the concentration of the highest-energy electrons in a very thin shell just behind the shock. Recent observations with the Chandra and XMM-Newton X-ray telescopes in space make it possible to resolve such filamentary structures, which are the result of strong synchrotron losses of the emitting multi-TeV electrons in strong magnetic fields downstream of the outer SN shock \\citep{ber05}.  Application of the nonlinear theory to SN~1987A \\citep{bk00} explained the properties of the observed radioemission. It predicted also the continuous increase of radioemission with roughly constant rate during the period from 1999 to 2006 yr, which was confirmed by further measurements \\citep{manch02}.  We analyze here radio data for more extended period of observation \\citep{mg01,manch02} together with the Chandra measurements of X-ray emission \\citep{park04} which provides a strong constraint on the amount of the nonthermal synchrotron emission at the same energies 0.5-10~keV. It is demonstrated that this constrain gives the evidence of strongly enhanced (amplified) magnetic field inside the SNR, that in turn provides a better fit of radio data than in our previous study. ",
        "conclusions": "Calculated shock radius $R_s$ and speed $V_s$ shown in Fig.\\ref{f1}a as a function of time are in satisfactory agreement with the values obtained on the basis of radio and X-ray measurements. Note that radio data compared with X-ray data gives larger shock size at any given time $t$ and calculated radius $R_s(t)$ goes between these two sets  To fit the spectral shape of the observed radio emission we assume a proton injection rate $\\eta=3\\times 10^{-3}$, which is a fraction of gas particles involved into the acceleration at the from of SN shock. This leads to a significant nonlinear modification of the shock: as it is seen in Fig.\\ref{f1}b  total shock compression ratio $\\sigma\\approx 5.3$ is essentially larger and a subshock compression ratio $\\sigma_s\\approx 2.8$ is lower than classical value 4.  Since the SN~1987A is very young CRs despite of very efficient acceleration accumulated only 7\\% of the explosion energy $E_{sn}$ (see Fig.\\ref{f1}c). Note that the injection is expected to be strongly suppressed at the quasiperpendicular part of the shock, therefore one should renormalize the results for the nucleonic spectrum, calculated within the spherically symmetric model. The lack of symmetry in the actual SNR can be approximately taken into account by a renormalization factor $f_{re}\\approx 0.2$, which diminishes the nucleonic CR production efficiency, calculated in the spherical model, and all effects associated with it \\citep[see][for details]{vbk03}. With  this renormalization the CRs inside SN~1987A already contain \\begin{equation} E_c=0.07f_{re}E_{sn}\\approx 2\\times 10^{49}~\\mbox{erg}. \\end{equation}  Since during the last 15~yrs the shock speed and the gas density are almost constant CR energy content grows roughly linearly with time. It gives the natural explanation of the linear increase of the radio emission detected during this evolutionary period.  Strongly modified SN shock generates CR spectrum $N\\propto p^{-\\gamma}$, which is very steep at momenta $p<m_pc$, with index $\\gamma =(\\sigma_s+2)/(\\sigma_s-1)\\approx 2.7$. CR electrons with such a spectrum produces synchrotron radioemission spectrum $S_{\\nu}\\propto \\nu^{-\\alpha}$ with spectral index $\\alpha =(\\gamma-1)/2\\approx 0.9$, that very well corresponds to the experiment, as it is seen in Fig.\\ref{f2}, where we present synchrotron energy spectra $\\nu S_{\\nu}$, calculated for five subsequent epoch together with the experimental data. Note that CR spectrum has a concave shape: it becomes flatter at higher momenta $p$. As a consequence synchrotron spectrum $S_{\\nu}(\\nu)$ is also concave as it is clearly seen in Fig.\\ref{f2} at $\\nu<10^{12}$~Hz. Radio data reveal this property in good consistency with theoretical prediction.  Strong downstream magnetic field $B_d\\approx 15$~mG, that corresponds to the upstream field $B_0\\approx 3$~mG (see Fig.\\ref{f1}), provides synchrotron cooling of electrons with momenta $p>10m_pc$ that in turn makes synchrotron spectrum at high frequencies $\\nu>10^{12}$~Hz very steep (see Fig.\\ref{f2}). Concave shape of electrons continuously produced at the shock front together with their synchrotron cooling lead to a formation of two peaks in synchrotron energy spectrum $\\nu S_{\\nu}$. The first one at $\\nu \\approx 10^{12}$~Hz corresponds to CR electron momentum $p\\approx 10m_pc$ above which synchrotron energy looses are relevant, whereas the second peak at $\\nu \\approx 10^{18}$~Hz corresponds to the maximum momentum $p\\approx 10^4m_pc$ of accelerated electrons. Under this condition calculated synchrotron flux at frequency $\\nu\\approx 10^{17}$~Hz, which corresponds to the photon energy $\\epsilon_{\\gamma}=0.5$~keV, is below the measured flux at the epochs $t>3000$~d. Since the contribution of the nonthermal radiation in the observed X-ray emission of SN~1987A is not very well known \\citep[e.g.][]{mich02}, the observed X-ray flux has to be considered as the upper limit for the expected nonthermal emission. At early epoch $t<2500$~d however the calculated flux exceeds the measured one (see the curve, corresponding $t=1970$~d in Fig.\\ref{f2}). This can be considered as indication, that the actual magnetic field $B_0$ is few times larger then given by the Eq.(3).  To illustrate the situation expected at considerably lower magnetic field we present in Fig.\\ref{f3} synchrotron energy spectra calculated at the same set of parameters as before except magnetic field, which was taken $B_0=200$~$\\mu$G independent of time. Since the interior magnetic field $B_d \\approx 2$~mG ($\\sigma \\approx 10$) is essentially lower in this case, synchrotron losses of high energy CR electrons are considerably smaller compared with the previous case. Due to this fact synchrotron spectra considerably exceeds at any given epoch the measured Chandra flux. Therefore we can conclude, that the actual interior magnetic field strength is not lower than $5$~mG.  Projected radial profile of the nonthermal X-ray emission at high interior magnetic field $B_d=10$~mG is characterized by an extremely sharp peak of thickness $L\\approx 10^{-3}R_s$ just behind the SN shock. Since the thermal X-rays have much wider radial distribution it provides the possibility for experimental determination of the nonthermal X-ray contribution in the observed emission. Experimental determination of the X-ray spectrum at $\\nu =10^{18}-10^{19}$~Hz ($\\epsilon_{\\gamma}=5-50$~keV) is the another possibility to discriminate the nonthermal emission. Such kind of measurements can be done with space instrument Suzaku.  Completely different possibility to have the synchrotron flux, which fit the radio data and goes below Chandra data, is low field scenario, when the magnetic field is as low as $B_0<2$~$\\mu$G. In such a case the cutoff frequency of synchrotron spectrum $\\nu_{max}$ is lower than $10^{17}$~Hz and calculated fluxes go down exponentially at $\\nu>\\nu_{max}$ below Chandra points. However the value $B_0<2$~$\\mu$G is unrealistically small for such a dense medium, which we have in SN~1987A. In addition the effective energy of electrons \\begin{equation} \\epsilon_e\\approx 5\\sqrt{(\\nu/\\mbox{1~GHz})/(B_d/10~\\mu\\mbox{G})}\\mbox{~~GeV}, \\end{equation} which radiate at $\\nu \\sim 1$~GHz in the interior field $B_d=10$~$\\mu$G is about 5~GeV. The electron spectrum at GeV-energies is characterized by a power low index $\\gamma\\approx 2$ that is considerably smaller than what is required for the observed radio spectrum. Note that as it is seen in Fig.\\ref{f3} even the field $B_d\\approx 2$~mG is too small to have a good fit of radio data as in Fig.\\ref{f2}. Therefore low field scenario should be rejected.  Calculated $\\gamma$-ray integral flux shown in Fig.\\ref{f4} at all energies is dominated by the $\\pi^0$-decay component. Since the SN shock is strongly modified $\\gamma$-ray spectrum at energies $\\epsilon_{\\gamma}>0.1$~TeV is very hard: $F_{\\gamma}\\propto \\epsilon_{\\gamma}^{-0.9}$. At the current epoch the expected $\\gamma$-ray energy flux at TeV-energies is about $\\epsilon_{\\gamma}F_{\\gamma}\\approx 2\\times 10^{-13}$~erg/(cm$^2$s) and during the next four years it expect to grow by a factor of two. This flux is by a factor of five lower compared with our previous prediction \\citep{bk00}, because we use here renormalization factor $f_{re} =0.2$, which reduces the number of accelerated CRs calculated within spherically symmetric approach. The existence of strongly asymmetric CSM structure, which is dense inner ring, makes our prediction of $\\gamma$-ray flux uncertain. According to the rough estimate this uncertainty is not very large, about a factor of two, due to the stronger SN shock deceleration in denser medium. At the moment there are only upper limits of TeV emission obtained by CANGAROO \\citep{eno03} and HESS \\citep{rowell04} instruments (see Fig.\\ref{f4})."
    },
    "0608/astro-ph0608065_arXiv.txt": {
        "abstract": "A reanalysis of the strongly metal-blanketed DZ white dwarf G165-7 is presented. An improved grid of model atmospheres and synthetic spectra is used to analyze $BVRI$, $JHK$, and $ugriz$ photometric observations as well as a high quality Sloan Digital Sky Survey spectrum covering the energy distribution from 3600 to 9000 \\AA.  The detection of splitting in several lines of Ca, Na, and Fe, suggesting a magnetic field of $B_s\\sim650$~kG, is confirmed by spectropolarimetric observations that reveal as much as $\\pm 7.5$\\% circular polarization in many of the absorption lines, most notably Na, Mg, and Fe.  Our combined photometric and spectroscopic fit yields $\\Te = 6440$~K, $\\logg=7.99$, $\\nh=-3.0$ and $\\log\\ ({\\rm Ca/He})=-8.1$. The other heavy elements have solar ratios with respect to calcium, with the exception of Na and Cr that had to be reduced by a factor of two and three, respectively.  A crude polarization model based upon the observed local spectral flux gradient yields a longitudinal field of 165~kG, consistent with the mean surface field inferred from the Zeeman splitting. The inclusion of this weak magnetic field in our synthetic spectrum calculations, even in an approximate fashion, is shown to improve our fit significantly. ",
        "introduction": "White dwarf stars of the DZ spectral type are cool, helium-atmosphere objects showing traces of heavy elements in their optical spectra. They are recognized mainly by the Ca~\\textsc{ii} H $\\&$ K doublet, while a few also show Ca~\\textsc{i} $\\lambda$4226, Mg~\\textsc{i} $\\lambda3835$ or Fe~\\textsc{i} $\\lambda3730$ (see \\citealt{SionAtlas}, \\citealt{WAtlas}, and \\citealt{harris03} for typical spectra). Since the heavy elements present in the atmospheric regions sink below the photosphere in a time scale much shorter than the white dwarf cooling time \\citep{paquette86}, the presence of these metals in DZ stars must be supplied by an external source, the most obvious mechanism being accretion from the interstellar medium, or, in some instances, from a circumstellar dust disk surrounding the white dwarf \\citep{zuckerman87,kilic05,becklin05}. Accretion of comets has also been proposed as a possible origin of metals in the atmospheres of DZ stars \\citep{Alcock86}. A precise determination of the heavy element abundances observed in the optical or UV spectra of DZ stars can thus provide important clues about the accretion rate and chemical composition of the interstellar or circumstellar medium.  Atmospheric abundance analyses of DZ stars are usually based on a restricted number of absorption lines detected in the optical portion of the spectrum \\citep[see, e.g.][]{LW87}, although a more important set of lines can be used when UV spectroscopy is available \\citep{Zeidler86,KW00,WKL02}. G165-7 (WD~1328+307, LHS 2745, $V=16.03$) is a particularly unusual DZ white dwarf that has practically no flux in the ultraviolet because of the severe blanketing by a large number of metallic lines, but it is also characterized by a very rich optical spectrum showing hundreds of absorption lines from various heavy elements such as Ca, Cr, Fe, Mg and Na. The first model atmosphere analysis of this star was carried out by \\citet{WL1980} who found $\\Te=7500\\pm 300$ K, $\\logg=8.0\\pm0.3$, $\\nh<-4.0$, and abundances for the heavy elements of about $1/30$ of the solar value under the assumption of scaled solar abundances.  Their analysis was based on low and moderate (3$-$5 \\AA) resolution spectroscopy in several segments covering the 3350$-$6700 \\AA\\ region, as well as on the multichannel spectrophotometry of \\citet{greenstein76}.  \\citet[][hereafter BLR]{blr01} also analyzed the broad-band energy distribution (optical $BVRI$ \\& near-infrared $JHK$ photometry) with a grid of pure helium model atmospheres and found $\\Te=7320$~K and $\\logg=8.21$ (the surface gravity in this case is constrained by the trigonometric parallax). However, analyses of the energy distribution of DQ white dwarfs based on pure helium models tend to overestimate the effective temperature by 300-800~K and the surface gravity by 0.10-0.24 dex with respect to the atmospheric parameters obtained from models that include carbon. This effect is understood in terms of an increase of the He$^-$ free-free opacity resulting from the additional free electrons \\citep{provencal02,dufour05}. A similar effect is thus expected in the case of DZ stars.  We present here a reappraisal of the analysis of G165-7 based on our latest metal blanketed atmospheric models and a new high quality spectrum. The observations are described in \\S~\\ref{observation} while in \\S~\\ref{theoretical} we describe our theoretical framework including our model atmosphere and synthetic spectrum calculations. Our detailed analysis is presented in \\S~\\ref{analysis} and our conclusions follow in \\S~\\ref{conclusion}. ",
        "conclusions": "Our study of G165-7 represents the first detailed model atmosphere analysis of a magnetic DZ white dwarf. High quality spectroscopic observations revealed that several metallic features were split by a surface magnetic field of approximately 650 kG. We achieved a best fit with a model atmosphere and synthetic spectrum at $\\Te=6440$~K, $\\logg=7.99$, $\\nh=-3$, and $\\log\\ ({\\rm Ca/He})=-8.1$, with other heavy elements in solar abundances relative to the calcium abundance (except for Na and Cr whose abundances had to be reduced by factors of 2 and 3, respectively). The predicted emergent spectrum provided an excellent match to the whole energy distribution of G165-7, as well as to most spectral features. By using the approximation of the linear Zeeman effect in our synthetic spectrum calculations, we were able to reproduce the splitting observed in several metallic lines. We also found that, contrary to the original analysis of G165-7 by \\citet{WL1980}, hydrogen is present with an abundance of $ \\nh\\sim-3$. With an effective temperature $\\sim1000$~K cooler than previously found, the molecular formation of MgH can no longer be completely ruled out. We have shown that the so-called ``asymmetric Mg I profile'' might possibly be explained by a MgH molecular feature instead of a quasi-static profile, although we cannot be entirely certain of this interpretation until a coherent theoretical framework becomes available. Other spectra of DZ stars with strong metallic absorption features have been found in the SDSS \\citep{harris03}. Analysis of them may lead to more insight into the properties of G165-7 and whether they are repeated in other DZ stars."
    },
    "0608/astro-ph0608253_arXiv.txt": {
        "abstract": "We perform Monte Carlo simulations of cosmic ray-induced hard X-ray radiation from the Earth's atmosphere. We find that the shape of the spectrum emergent from the atmosphere in the energy range 25--300~keV is mainly determined by Compton scatterings and photoabsorption, and is almost insensitive to the incident cosmic-ray spectrum. We provide a fitting formula for the hard X-ray surface brightness of the atmosphere as would be measured by a satellite-born instrument, as a function of energy, solar modulation level, geomagnetic cutoff rigidity and zenith angle. A recent measurement by the INTEGRAL observatory of the atmospheric hard X-ray flux during the occultation of the cosmic X-ray background by the Earth  agrees with our prediction within 10\\%. This suggests that Earth observations could be used for in-orbit calibration of future hard X-ray telescopes. We also demonstrate that the hard X-ray spectra generated by cosmic rays in the crusts of the Moon, Mars and Mercury should be significantly different from that emitted by the Earth's atmosphere. ",
        "introduction": "\\label{s:intro}  Starting in the 1960's, observations from balloons \\citep{petetal73,danste74,schetal80} and spacecraft \\citep{goletal74,schpet74,imhetal76,wilmah92,shamur01,shaetal03} have shown that the Earth's atmosphere is bright in hard X-rays and gamma-rays owing to its bombardment by cosmic rays of Galactic origin. This is the case despite the fact that the energy deposited by the cosmic rays into the atmosphere is some nine orders of magnitude smaller than the total energy budget coming to us from the Sun and that only $\\sim 1$\\% of the energy deposited by the cosmic rays is re-radiated into space as gamma-rays and much less as hard X-rays.  Atmospheric hard X-ray emission is of interest not only because it provides insight into the interaction of energetic cosmic particles with air, but also in connection with the possibility of measuring the spectrum of the cosmic X-ray background (CXB) using the Earth as a shield. In January--February 2006 a series of observations following this approach were performed by the INTEGRAL orbital observatory \\citep{chuetal07} and drew our attention to the problem of atmospheric X-ray radiation. In observations of this kind, cosmic ray-induced emission on the one hand causes significant contamination of the CXB decrement (due to occultation by the Earth) at energies $\\sim 30$--100~keV, but on the other, could in principle be used for calibrating a detector's response in the high-energy region (substantially above 100~keV) where atmospheric emission dominates over the CXB signal. During the INTEGRAL observations the Earth covered $\\sim 100$~sq. deg of the sky and was as bright as $\\sim 200$~mCrab in the spectral band 100--200~keV \\citep{chuetal07}.  Despite the long history of experimental (see references above) and theoretical \\citep{puskin70,danste74,ling75,grasch77,deaetal89} studying, there remains significant uncertainty regarding the spectrum of atmospheric X-ray and gamma-ray radiation. However, the nuclear and electromagnetic processes responsible for the generation and propagation of X- and gamma-rays in the atmosphere, although complex, are well understood and already incorporated in computer codes. Also the incident energy spectrum of Galactic cosmic rays has been measured with high precision (see e.g. \\citealt{gaietal01} for a review) and there is a good understanding of how it is modified below $\\sim 10$~GeV by the magnetic fields of the Solar system and Earth. All this suggests that the energy and angular distributions of atmospheric X-ray and gamma-ray emission should be predictable by simulations with fairly good ($\\sim 10$\\%) accuracy.  We have performed such a numerical modelling using the popular toolkit {\\sl GEANT} \\citep{geant03}, focusing on making predictions for observations of the Earth from space in the hard X-ray band 25--300~keV. We show below that the hard X-ray emission of the Earth is characterized by a well-defined spectral shape (Sect.~\\ref{s:simul}) and that its intensity depends in a predictable way on the phase of the solar cycle and on the geomagnetic rigidity map corresponding to a given observation (Sect.~\\ref{s:predict}). Therefore the Earth could become a good absolute calibrator for future orbital hard X-ray telescopes, in particular when used in combination with accurate monitoring of the cosmic-ray spectrum by large-area, high-resolution particle physics detectors such as the Alpha Magnetic Spectrometer to operate on the International Space Station \\citep{baretal04}.  In addition we demonstrate (Sect.~\\ref{s:planets}) that the cosmic ray-irradiated solid crusts of the Moon, Mars and Mercury should generate hard X-ray spectra different from the one produced in the Earth's atmosphere. ",
        "conclusions": "We have revisited the problem of the X-ray and gamma-ray emissivity of the Earth's atmosphere taking advantage of the recent high-quality measurements of the Galactic cosmic-ray spectra and the availability of efficient Monte Carlo codes incorporating state-of-the-art descriptions of hadronic physics and electromagentic cascades. We found that emission generated in the energy band 25--300~keV has a practically invariant spectral shape and a flux that is expected to show significant variations both with time (over the solar cycle) and with position on the Earth (due to geomagnetic effects). We provided two simple expressions -- equations~(\\ref{eq:spec_fit}) and (\\ref{eq:norm_fit}) -- which allow one to predict the spectral intensity of atmospheric hard X-ray emission for a particular observation from a spacecraft. It should be noted that our results are not directly applicable to observations performed from balloons floating at altitudes below $\\sim$~50~km, since then in addition to the upward flux there will also be a significant amount of downward radiation generated in the several overlying g~cm$^{-2}$ of air.  The statistical accuracy of our predictions is better than 5\\%, while the systematic uncertainty is probably dominated by the current uncertainty in our knowledge of the cosmic-ray flux, which is $\\sim 10$\\%. We can try to assess the overall accuracy of our model through direct comparison with real observations of the Earth from spacecraft. Among the few published measurements of atmospheric emission in the hard X-ray domain are those done in the 1970's from the satellites Kosmos~461 {\\citep{goletal74} and 1972-076B \\citep{imhetal76}. The spectra obtained in the energy range $\\sim$~50--1000~keV could be well described by a power law with a photon index ($dN_\\gamma/dE\\propto E^{-\\Gamma}$) $\\Gamma=1.55\\pm 0.10$ in the former experiment and $\\Gamma\\sim$~1.31--1.47 in the latter. These values are in good agreement with our calculated value of 1.4 for the spectral slope above $\\sim 50$~keV. Furthermore, \\cite{goletal74} (see also \\citealt{mazetal75}) measured the dependence of the hard X-ray flux on geomagnetic cutoff rigidity over the range $3<\\rcut~({\\rm GV})<17.5$, and that result too appears to be consistent with our computations.  Recently during a scanning observation of the Earth, the INTEGRAL observatory measured the spectrum of the cosmic X-ray background together with that of atmospheric hard X-ray emission. Our analysis of the data \\citep{chuetal07} showed that the atmospheric spectrum at energies from $\\sim 40$ to $\\sim 200$~keV is consistent with the prediction of the present study, although the spectral shape was not well constrained by the data. When we adopted the spectral shape given by equation~(\\ref{eq:spec_fit}), we found the best-fit amplitude of the atmospheric component to lie within 10\\% of (actually only $\\sim 4$\\% higher than) the flux predicted by equation~(\\ref{eq:norm_fit}) for the relevant value of the solar modulation potential ($\\phi=0.45$~GV) and taking into account the geomagnetic cutoff ridigity map corresponding to the INTEGRAL observations \\citep{chuetal07}.  This good agreement of our model with observations suggests that the Earth could be used as an absolute calibrator of future satellite-born sensitive hard X-ray detectors, complementing the Crab nebula. This possibility as well as the predictive power of our model could be tested soon via new INTEGRAL observations of the Earth performed at a different phase of the solar cycle.  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{planets_albedo.ps} \\caption{{\\sl Top:} Total X-ray spectral surface brightness of the Earth (soild lines) -- sum of cosmic ray-induced atmospheric emission (dotted lines) and the reflected CXB radiation (dashed lines). The two sets of curves correspond to observations of a geomagnetic polar region ($\\rcut=0$) during solar minimum ($\\phi=0.5$~GV) and of the geomagnetic equator ($\\rcut=15$~GV) during solar maximum ($\\phi=1.5$~GV).  {\\sl Bottom:} The same but for the Moon. The two sets of curves correspond to observations performed during solar minimum and solar maximum. } \\label{planets_albedo} \\end{figure}  We also demonstrated that under the impact of cosmic rays the solid surfaces of the Moon, Mars and Mercury should be about as bright in hard X-rays as (in fact somewhat brighter than) the Earth's atmosphere, making these objects potentially interesting for observations. It is important to keep in mind though that in a real experiment reflected CXB radiation will significantly contribute to the observed signal at energies below $\\sim 50$~keV. Based on the computations of \\citet{chuetal07b}, we show in Fig.~\\ref{planets_albedo} several examples of hard X-ray spectra that could be observed from the Earth and Moon. This figure indicates that the cosmic ray-induced hard X-ray continuum could complement the CXB-induced fluorescent X-ray lines in probing the chemical composition of the crusts of planets and the Moon, i.e. during the Bepi Colombo mission to Mercury\\footnote{http://sci.esa.int/home/bepicolombo}."
    },
    "0608/astro-ph0608247_arXiv.txt": {
        "abstract": "{ We have measured the gamma-ray fluxes of the blazars \\object{Mrk~421} and \\object{Mrk~501} in the energy range between 50 and 350~GeV ($1.2$ to $8.3 \\times 10^{25}$~Hz). The detector, called CELESTE, used first 40, then 53~heliostats of the former solar facility ``Th\\'emis'' in the French Pyrenees to collect Cherenkov light generated in atmospheric particle cascades. The signal from Mrk~421 is often strong. We compare its flux with previously published multi-wavelength studies and infer that we are straddling the high energy peak of the spectral energy distribution. The signal from Mrk~501 in 2000 was weak ($3.4~\\sigma$). We obtain an upper limit on the flux from 1ES~1426+428 of less than half that of the Crab flux near 100~GeV. The data ana\\-ly\\-sis and understanding of systematic biases have improved compared to previous work, increasing the detector's sensitivity. ",
        "introduction": "Measurements of high energy emission from active galactic nuclei (AGN) give insights into a variety of open problems, such as the nature of the AGNs themselves, and the density and evolution of the extragalactic diffuse infrared background. After GLAST is launched in 2007 (\\cite{glast}), the large number of high galactic latitude GeV gamma-ray sources to be seen will make AGNs become a background for searches for new classes of emitters~-- insufficient understanding of high energy AGNs might just limit potential glimpses of a hidden Universe. Several AGNs of the ``blazar'' class emit above 250~GeV and have been detected by atmospheric Cherenkov imaging telescopes. They are called HBLs, for high energy BL Lac objects. Of these, Mrk~421 (\\cite{blaze,cat421,hegra421b}) and Mrk~501 (\\cite{ir,cat501}) are the brightest .  The Synchrotron Self Compton model (``SSC'') is consistent with much of the multiwavelength blazar data. For discussion of these models as applied to the blazars discussed in this paper, see for example (\\cite{TMG}) or (\\cite{KSK}). In its simplest variant, relativistic electrons generate synchrotron photons with a $\\nu F_\\nu$ peak in the optical to X-ray range, and then up-scatter these same photons to produce the peak seen in the GeV range. The inverse Compton peak position can constrain the physical parameters of the model. Many, if not most, observed characteristics of blazars remain unexplained. Examples are the relation between variations at different wavelengths, as highlighted by the ``orphan'' gamma ray flares observed for 1ES1959+650 (\\cite{kwaz1959}), and the role that proton acceleration might play. Progress requires simultaneous measurements across the spectrum.  The CELESTE telescope was one of the first instruments sensitive near the Compton peak for these blazars, beyond the energy range of the rela\\-ti\\-ve\\-ly large sample of EGRET blazars (\\cite{egretblazars,3rdcatalog}). Mrk~421 was studied with EGRET, but Mrk~501 was visible only during a flare (\\cite{egret501}), while 1ES 1426+428 was undetected by EGRET, although well-measured at TeV energies (e.g. \\cite{cat1426,whipple1426}). A key motivation for our experiment was to provide blazar measurements in an unexplored energy window.   This paper presents the first observations of three blazars below 100~GeV, by a solar heliostat array, and is structured as follows: we describe the detector, the analysis method, and the data samples. We emphasize improvements in our knowledge of the systematic uncertainties. We present light curves for the three blazars as well as the integrated fluxes, and place the results in the context of SSC models. ",
        "conclusions": "We have provided some of the first astrophysical flux measurements around 100~GeV, beyond the reach of EGRET on the Compton CGRO and below the range of previous ground-based telescopes, by recording atmospheric Cherenkov light with the reconverted solar tower facility at Th\\'emis. We have increased our detector sensitivity as compared to our earlier work, to a level near that predicted in our original proposal. The gain came mainly from a better analysis procedure but also from refinements in our detector simulation and electronics improvements. 80~\\% of the gamma-rays we detect have energies between 50 and 350~GeV. Bad weather at our site lead to a low duty-cycle, resulting in a modest number of detections in spite of our detector's good performance.  Our results for the blazar Mrk~421 help constrain the shape of the high energy peak of the spectral energy distribution. We have made the first detection of Mrk~501 in this range, although the detection is weak, and we have placed an upper limit on the flux from the distant blazar 1ES~1426+428 during a period of X-ray inactivity.   \\begin{figure*} \\hspace*{-3.cm} \\centering\\includegraphics[width=0.7\\textwidth, angle=270]{./1ES1426_LC.eps} \\caption{Light curves for 1ES~1426+428~--  CELESTE detected no signal. The CELESTE daily averages run from 2002 March 22 to 2004 April 25. The ASM points are 5 day running averages. \\label{1426lc}} \\end{figure*}"
    },
    "0608/gr-qc0608036_arXiv.txt": {
        "abstract": "{Measurement of gravitomagnetic field is of fundamental importance as a test of general relativity. Here we present a new theoretical project for performing such a measurement based on detection of the electric field arising from the interplay between the gravitomagnetic and magnetic fields in the stationary axial-symmetric gravitational field of a slowly rotating massive body.  Finally it is shown that precise magnetometers based on superconducting quantum interferometers could not be designed for measurement of the gravitomagnetically induced magnetic field in the cavity of a charged capacitor since they measure the circulation of a vector potential of electromagnetic field, i.e., an invariant  quantity including the sum of electric and magnetic fields, and the general-relativistic magnetic part will be totally cancelled by the electric one which is in good agreement with the experimental results. } ",
        "introduction": "General relativity predicts that a spinning central massive body creates a gravitomagnetic field in addition to the Newtonian-like monopolar gravitoelectric one, that is static gravitational mass generates solely gravitoelectric field and a moving body, an additional gravitomagnetic one as in electrodynamics (see, for review, ~\\cite{ignaz}). Almost all the classical tests of general relativity (gravitational redshift, perihelion precession of Mercury, gravitational deflection of light) confirm the validity of the general-relativistic corrections to the Newtonian gravitoelectric field with high accuracy.  But due to the weakness of the gravitomagnetic field in the Solar system it is very hard to detect it although many theoretical projects on measuring gravitomagnetism have been proposed ~\\cite{ignaz}.  In our recent papers~\\cite{99pla,00adp} the gravitomagnetic effect on the electric current and magnetic field has been investigated on the theoretical level. In fact, the influence of the angular momentum of the rotating gravitational source may appear as a galvanogravitomagnetic effect in the current carrying conductors~\\cite{99pla} and as a general-relativistic effect of charge redistribution inside conductors in an applied magnetic field ~\\cite{00adp}. It is natural to ask whether the gravitomagnetic interaction with electric field can lead to the applicable general-relativistic effects. Our aim here is to investigate an answer to this question and present a new measurement scheme which could lead to a detection of the earth's gravitomagnetic field in the electromagnetic systems.  The paper is organized as follows: in section~\\ref{em-capacitor} we deal with electromagnetic fields in both the cavity of charged capacitor and solenoid embedded in a space-time of slow rotating gravitational object.  The next section~\\ref{feasibility} is devoted to a proposal for measuring gravitomagnetic field in the electromagnetic systems.  In section~\\ref{vasil} we discuss the impossibility of detection of gravitomagnetic field through direct measurement of a gravitomagnetically induced magnetic field in the cavity of charged capacitor. Finally section~\\ref{conclusion} contains concluding remarks. ",
        "conclusions": "\\label{conclusion}  As we have shown here, on one side, due to the gravitomagnetic effect, a magnetic field~ (\\ref{field}) will be induced inside a cavity of condenser in addition to the electric one. On other side, the second term arising on the right hand side of the flux~(\\ref{zero}) totally compensates the gravitomagnetically induced magnetic field. In this sense the experiment~\\cite{vas78} is analogous to the null result experiment~\\cite{jain} on test of the principle of equivalence for electromagnetic systems, where the gravitational redshift of the microwave frequency $\\Delta\\nu/\\nu=\\lambda$ was totally compensated by the general-relativistic effect due to the fact that the gravitoelectrochemical potential ${\\tilde{\\mu_e}}=\\mu_e(1+\\lambda )$ (rather than the electrochemical one $\\mu_e$) is constant during thermodynamical equilibrium (parameter $\\lambda =Lg/c^2$ depends on the height difference $L$ and the acceleration due to gravity $g$). However both compensating general-relativistic effects were separately detected in a number of experiments, for example, the gravitational redshift for the frequency by Pound and Rebka~\\cite{pound} and the electric field inside conductors produced by the inhomogeneity of the chemical potential in the gravitational field by Witteborn and Fairbank~\\cite{wf,wf71,lwf}.  The advantage of the time-of-motion experiment over the Vasil'ev's one ~\\cite{vas78} is that a non-null result is predicted by general relativity as opposed to the null result for the latter experiment. Also in the proposed experiment, with the time-of-motion technique, Newtonian physics would predict a null result, whereas in the proposed experiment it would be necessary to assume the existence of the earth's gravitomagnetic field to distinguish between the results predicted by general relativity and Newtonian gravity.  It should be mentioned that on earth, the angular velocity of the capacitor with respect to a local inertial frame ${\\mathbf\\Omega}_{L}$ is given by~\\cite{ignaz}:  $${\\mathbf\\Omega}_{L}={\\mathbf\\Omega}_{VLBI}-{\\mathbf\\Omega}_{Th}- {\\mathbf\\Omega}_{S}-{\\mathbf\\Omega}_{LT},$$  where ${\\mathbf\\Omega}_{VLBI}$ is the angular velocity of the laboratory with respect to an asymptotic inertial frame, ${\\mathbf\\Omega}_{Th}$ and ${\\mathbf\\Omega}_{S}$ are, respectively, the contributions of the Thomas precession arising from non-gravitational forces and of the de Sitter or geodetic precession. As a result, in order to detect ${\\mathbf\\Omega}_{LT}$ (which is $\\omega$ in metric~(\\ref{eq:corot})) one should measure ${\\mathbf\\Omega}_{L}$ and then substract from it the independently measured value of ${\\mathbf\\Omega}_{VLBI}$ with VLBI (Very Long Baseline Interferometry, see, for example, ~\\cite{koval}) and the contributions due to the Thomas and de Sitter precessions.  The typical parameters for earth which can be considered as uniformly rotating body are $a_\\oplus=({2}/{5}) \\Omega_\\oplus R_\\oplus^2$, $2\\times G{M_\\oplus}/{c^2}=2\\times 0.44cm$, $R_\\oplus\\approx 6.37\\times 10^8cm$, $\\Omega_\\oplus=7.27\\times 10^{-5}rad/s$ and consequently $\\Omega_{LT}\\approx 10^{-9}\\Omega_\\oplus$ is extremely weak. The experiment is quite difficult since the measured effect is small compared to potential disturbances and careful attention must be paid to possible systematic experimental errors coming from seismic accelerations, local gravitational noise, earth's electromagnetic field, atmospheric disturbances etc. We mention only how  few major difficulties can be reduced: (i) shielding the earth's electric field with help of Faraday cage with high accuracy, (ii) shielding the earth's magnetic field by means of superconducting shells, and (iii) vacuum measurements to prevent the effects arising from the air ionization.  We do not provide here any numerical estimates for the experimental technique since we would like only to underline the existence of the possibility of detecting the gravitomagnetically generated electric field. The detailed design and study of the method under real experimental conditions is under our consideration. However for some technical details and study of unwanted effects as electrostatic effects, magnetostatic effects, the patch effect, the electron and the lattice sag effects, thermal and vacuum effects one can look to~\\cite{wf71}.  From our research, we can conclude that (i) the effect of the gravitomagnetic force on the elecrostatic (magnetic) field is to induce a magnetic (electric) field, (ii) gravitomagnetically generated magnetic field can not be directly measured by the SQUID which is in agreement with the results of the null experiment~\\cite{vas78}, (iii) the electric field arising from the interplay between the general relativistic gravitomagnetic and magnetic fields seems to be experimentally verifiable with the measurement of the trajectory-of-motion of charged particles (see equations~(\\ref{sol_period}))."
    },
    "0608/astro-ph0608071_arXiv.txt": {
        "abstract": "The Rossiter-McLaughlin (RM) effect is the distortion of stellar spectral lines that occurs during eclipses or transits, due to stellar rotation. We assess the future prospects for using the RM effect to measure the alignment of planetary orbits with the spin axes of their parent stars, and to confirm exoplanetary transits. We compute the achievable accuracy for the parameters of interest, in general and for the 5 known cases of transiting exoplanets with bright host stars. We determine the requirements for detecting the effects of differential rotation. For transiting planets with small masses or long periods (as will be detected by forthcoming satellite missions), the velocity anomaly produced by the RM effect can be much larger than the orbital velocity of the star. For a terrestrial planet in the habitable zone of a Sun-like star found by the {\\it Kepler} mission, it will be difficult to use the RM effect to confirm transits with current instruments, but it still may be easier than measuring the spectroscopic orbit. ",
        "introduction": "}  When an exoplanet transits the disk of its parent star, there is both a photometric signal and a spectroscopic signal. The photometric signal is a small reduction in the received flux due to the partial obscuration of the stellar disk, as first detected by Charbonneau et al.\\ (2000) and Henry et al.\\ (2000). The transit light curve depends chiefly on the radius of the planet, the radius of the star, the orbital inclination, and the stellar limb darkening function. The spectroscopic signal is subtler and less familiar. Changes in the stellar absorption lines can be produced by absorption features in the planetary atmosphere (see, e.g., Charbonneau et al.\\ 2002, Vidal-Madjar et al.\\ 2003). Or, if the emergent stellar spectrum varies with position across the stellar disk, then the planetary obstruction will cause changes in the disk-integrated spectrum.  In particular, the Rossiter-McLaughlin (RM) effect refers to the spectral distortion caused by the spatial variation in the emergent spectrum due to stellar rotation. The exposed portion of the photosphere has a net rotational Doppler shift. When the planet covers part of the blueshifted half of the stellar disk, the integrated starlight appears slightly redshifted, and vice versa. Thus, the spectral distortion of the RM effect is often manifested as an ``anomalous'' radial velocity, i.e., a Doppler shift that is greater or smaller than the shift expected from only the star's orbital motion. Although the RM effect has a long history in the context of eclipsing binary stars (Forbes 1911, Schlesinger 1911, Rossiter 1924, McLaughlin 1924), its importance in the context of exoplanets is only starting to be appreciated.  Suppose that the photometric signal has already been observed, and that the parameters governing the photometric signal have been accurately determined. The two most important additional parameters that govern the RM effect are $V_S \\sin I_S$, the projected rotation speed of the stellar surface, and $\\lambda$, the angle between the sky projections of the stellar spin axis and the orbit normal (i.e., the angle between the transit chord and lines of stellar latitude, for $I_S=90\\arcdeg$). Observations of the RM effect allow these two parameters to be measured (see, e.g., Queloz et al.\\ 2000, Ohta et al.\\ 2005, Winn et al.\\ 2005, Wolf et al.~2006). The latter parameter, $\\lambda$, is especially interesting, as it provides information about exoplanetary spin-orbit alignment. Solar system planetary orbits are generally within $\\sim$5\\arcdeg\\ of the solar equatorial plane (Beck \\& Giles 2005), but this may or may not be true for the full range of exoplanetary systems.  Of particular interest in this regard are hot Jupiters---planets with masses of $\\sim$$M_{\\rm Jup}$ and orbital periods of $\\sim$3~days---because there are some theoretical reasons why one might expect large misalignments for such planets.   Hot Jupiters are thought to have formed at large orbital distances, and then migrated inward to their current positions. The migration mechanism is still unknown, and some of the proposed mechanisms differ in the degree to which they would affect spin-orbit alignment. Thus, measurement of spin-orbit alignment offers a possible means of discriminating among migration theories. The most widely discussed category of migration theories involves disk-planet interactions (Lin et al.\\ 1996). There is a large literature on this subject (as reviewed by Papaloizou \\& Terquem 2006), but broadly speaking, there is no apparent reason why these mechanisms would perturb spin-orbit alignment; in fact they may even drive the system toward closer alignment. In contrast, other migration theories involve disruptive events such as planet-planet interactions (Rasio \\& Ford 1996, Weidenschilling \\& Marzari 1996) or planetesimal collisions (Murray et al.\\ 1998), which could act to randomize spin-orbit alignment. Another proposal involves the Kozai mechanism, in which a companion star causes oscillations in the planetary orbit's eccentricity and inclination (Innanen et al.\\ 1997, Holman et al.\\ 1997). By the time tides circularize the orbit and halt ``Kozai migration,'' the orbital inclination can change substantially (Wu \\& Murray 2003; Eggenberger, Udry, \\& Mayor 2004; D.\\ Fabrycky \\& S.\\ Tremaine, priv.\\ comm.).  For planets at larger orbital distances, there is no particular reason to expect large misalignments, but the field of exoplanets has rewarded observers with surprises in the past. There are also a few empirical hints of multiple orbital planes in some systems, such as the double debris disk recently reported around $\\beta$~Pic (Golimowski et al.\\ 2006), and the apparently counter-rotating disks around the protostar IRAS~16293--2422 (Remijan \\& Hollis 2006).  The exoplanetary RM effect was first observed by Queloz et al.\\ (2000) and Bundy \\& Marcy (2000) during transits of HD~209458b. The former authors were able to place an upper bound on $\\lambda$ of about 20\\arcdeg. Snellen (2004) described the interesting idea of using the wavelength-dependence of the RM effect to search for planetary absorption features, but with only null results for HD~209458b. Winn et al.\\ (2005) used improved photometric and spectroscopic data to show that $\\lambda =-4^\\circ\\hskip-4.5pt.4\\pm 1^\\circ\\hskip-4.5pt.4$, a small but significant misalignment reminiscent of Solar system planets.  \\citet{wolf06} performed a similar study of HD~149026b, finding $\\lambda = 11\\pm 14\\arcdeg$. Apparently, in these two cases, the migration mechanism was fairly quiescent. We can soon expect similar studies of the other transiting planets with host stars bright enough for high-precision spectroscopy, namely, HD~189733b \\citep{bouchy05c}, TrES-1 \\citep{alonso04}, and XO-1b \\citep{mccullough06}. The first goal of this paper is to provide useful guidance for future observations of these and other systems.  The second goal of this paper is to investigate another important and timely application of the RM effect: transit detection and confirmation. The spectroscopic detection of transits offers certain advantages, in some cases, over photometric detection.  Indeed, \\citet{bouchy05c} discovered the transits of HD~189733b through observations of the RM effect, and followed up this discovery with photometric observations. The reason for this ordering may simply have been that the observers had more convenient access to spectroscopic resources than to photometric resources. However, as pointed out by Welsh et al.\\ (2004) and Ohta et al.\\ (2005), the importance of the RM effect in transit confirmation is likely to increase soon. The satellite missions {\\it Corot} \\citep{baglin03} and {\\it Kepler} \\citep{borucki03} aim to find smaller and longer-period transiting planets than those currently known, for which it will be very difficult to confirm the existence of transits photometrically from the ground. It has been envisioned that the confirmation process will require measuring the orbital velocity of the parent star, and there is hope that this can be achieved, given the recent excellent progress in improving the accuracy of Doppler measurements \\citep{mayor03, marcy05}. As we will show in this paper, the RM effect offers an alternative path to confirmation that is suitable for at least a subset of stars.  Previous theoretical work on this topic has concentrated on analytic descriptions of the spectral distortion (Ohta et al.\\ 2005, Gimenez 2006), or on numerical simulations of the spectral distortion involving a discretized model of the stellar surface (Queloz et al.\\ 2000, Welsh et al.\\ 2004). In this work, we are not concerned with high accuracy in describing the spectroscopic distortion; instead we are concerned with the measurement and estimation problem. In \\S~\\ref{sec:rmeffect}, we explain our notation, review some of the previously derived results and provide some useful approximate scaling relations. In \\S~\\ref{sec:character}, we estimate the achievable accuracy in measuring the key RM parameters, $V_S \\sin I_S$ and $\\lambda$, as a function of the orbital geometry of the system and of the characteristics of the data. We do this both for the general case, and for the specific cases of the 5 transiting exoplanets with bright host stars. A secondary parameter that affects the RM signal is the degree of differential rotation across the stellar disk, and in \\S~\\ref{sec:limb} we investigate whether or not this effect is important for near-term observations. In \\S~\\ref{sec:confirm}, we derive an analytic expression for the signal-to-noise ratio ($S/N$) in the detection of the RM effect, and we apply this formula to assess the prospects for state-of-the-art Doppler measurements to confirm the existence of transits of planets detected by {\\it Kepler} and {\\it Corot}. The final section summarizes all of these results and suggests some avenues for future work. ",
        "conclusions": "\\label{sec:discussion}  The transit of an exoplanet is a rare and information-rich event that should be exploited in every possible way. In this paper we have examined one aspect of transit physics, the Rossiter-McLaughlin effect, whose origin is the rotation of the parent star. Our goal has been to assess the ability of near-term observations to exploit the information in the RM signal.  We have calculated the achievable accuracy with which one can measure the key parameter $\\lambda$, with which one can assess the degree of alignment between the stellar spin axis and the planetary orbit, and have assessed the prospects for the currently known sample of transiting exoplanets. One might ask further: how accurately must $\\lambda$ be known in order to enable interesting theoretical advances, such as the discrimination among different planetary migration theories? In other words, how good is good enough? This question does not seem to have been addressed in the literature on planet formation theory, and is an appealing topic for future research. As suggested in \\S~3, for a system with a very hot Jupiter, even a crude measurement with $\\sigma_\\lambda=30\\arcdeg$ would be of interest. As a more general benchmark, one would like to achieve at least enough accuracy to tell whether a given system is similar to the Solar system, or not. To date, the only system that has been observed with enough precision for this task is HD~209458, and the result is that it is Solar-like with a small but significant misalignment of $\\lambda\\approx -4\\arcdeg$ (Winn et al.~2005). A system could be different either by having a larger misalignment ($\\lambda \\gsim 10\\arcdeg$), for which an accuracy of a few degrees or more would suffice, or by having a more perfect alignment, for which an accuracy of $\\sim 0.1\\arcdeg$ would be needed.  We have also explained and quantified the potentially important role of the RM effect in the confirmation of transiting planet candidates. It is a difficult task, but in many cases it will be easier than the task of detecting the orbital velocity of the parent star. The comparative advantages of RM confirmation are that the velocity amplitude is generally larger, and the {\\it acceleration} is certainly larger. The full range of RM variations occur during a single transit (typically a single night or two) whereas the orbital velocity variations occur over the full orbital period (a year, for the interesting case of an Earth analog). A disadvantage of RM confirmation is that one does not learn the planetary mass; it is not a dynamical measurement, {\\it per se}, but rather an alternate method of verifying that a portion of the stellar surface is periodically occulted. Furthermore, RM observations must take place during transits, whereas the spectroscopic orbital data can be obtained on any other nights. In addition, the ease or difficulty of the RM method depends on the rotation rate of the star, which will vary from system to system.  Another interesting question for future research is whether or not the RM effect could ever profitably become the basis of a planet {\\it search} technique, rather than simply a follow-up technique. Ohta et al.~(2005) suggested that transiting planets might announce themselves through large outliers in data bases of radial velocity measurements. The relevant efficiency factors, selection effects, and detection algorithms for such an undertaking have yet to be worked out.  Even less certain are the prospects for an RM-based transit survey, in which stars are spectroscopically monitored specifically for transits. The obvious drawback is that typical high-resolution spectrographs examine only one star at a time, whereas photometric transit surveys observe thousands or even millions of stars at a time. There are plans for multiplexed Doppler spectrographs in the near future (Ge et al.~2004), but the multiplexing factor is only $\\sim$50.  On the other hand, the RM effect is perhaps uniquely sensitive to small planets around hot and rapidly rotating stars. The amplitude of the RM effect depends on the stellar radius and projected rotation speed, according to Eq.~(\\ref{eqn:krm}), but these factors are not completely independent. Among main-sequence stars, both the stellar radius and the typical rotation speed increase with stellar mass (or temperature), with an especially dramatic rise near the F5 boundary. The typical rotation speed of G0 stars is $\\sim$10~km~s$^{-1}$, and for F0 stars it is $\\sim$100~km~s$^{-1}$ (see, e.g., Tassoul~2000). This is related to the strong variation in the depth of the convective zone across this same boundary.  We have attempted to illustrate the net effect of these astrophysical correlations on the strength of the RM signal, in the following manner. We considered main-sequence stars ranging in spectral type from O5 to G0, and imagined that a Jupiter-sized planet makes an equatorial transit of each star.\\footnote{Later-type stars are not amenable to even this crude analysis, as there is a very wide dispersion in velocity (and achievable velocity accuracy) that is linked to age and stellar activity through the relations of Skumanich~(1972) and Noyes et al.~(1984).} We calculated $K_R$ using estimates of the stellar radius from Cox et al.~(2000), and estimates of typical rotation speeds from Tassoul~(2000). The results are shown in Fig.~\\ref{fig:kr-vs-spt}. The maximum signal is achieved for late A stars.  Of course, the achievable velocity accuracy is also a function of stellar type, as it depends on the number and width of absorption lines available for velocity determination. We have not attempted to quantify the velocity accuracy because in this regime it would certainly be advantageous to abandon the description of the RM effect as a net Doppler shift and model the distortion directly, as is done in Doppler-imaging algorithms.  \\begin{figure}[ht*] \\epsscale{1.0} \\plotone{kr-vs-spt.eps} \\caption{ Rough estimate of the amplitude of the RM effect of a transiting Jovian planet as a function of the spectral type of the parent star. The effect reaches a maximum for late A stars, which is a compromise between the faster rotation rates and the larger radii of hot stars. \\label{fig:kr-vs-spt} } \\end{figure}  As with so many other areas of stellar and planetary physics, interest in the Rossiter-McLaughlin effect has been re-invigorated by the discovery of exoplanets. It has been nearly 100 years since the first observations of this effect, and yet we believe that there are still important applications waiting to be developed."
    },
    "0608/astro-ph0608592_arXiv.txt": {
        "abstract": "{ % The absolute calibration of space-borne instruments in the soft X-ray regime rests strongly on model spectra of hot white dwarfs.} { % We analyze the \\chandra\\ LETG+HRC observations of the white dwarfs \\ohz\\ and \\osi\\ and of the neutron star \\orx\\ with the aim of resolving current uncertainties in the soft X-ray spectral fluxes and photospheric parameters of the three stars. We apply the derived photon spectra to a cross-calibration of the LETG+HRC-S with the short-wavelength \\euve\\ spectrometer and the \\rosat\\ PSPC.} { % We tie \\ohz\\ to the flux of \\rxj\\ in the 44--48\\AA\\ range and perform a simultaneous least squares fit to the LETG+HRC spectra of the three stars.  This allows us to determine an internally consistent set of spectral energy distributions and an empirically derived wavelength-dependent correction to the LETG+HRC-S effective area. We employ NLTE model atmospheres calculated with \\textit{TMAP} for the white dwarfs and a two-blackbody model for \\rxj, tied to the respective optical fluxes.} { % The two-blackbody model for \\rxj\\ features a hot spot on a cooler star and yields k$T_\\mathrm{spot}=62.8\\pm0.4$\\,eV and k$T_\\mathrm{star}=32.3\\pm0.7$\\,eV with a stellar angular radius as seen from infinity of $0.1371\\pm0.0010$\\,km\\,pc$^{-1}$. For \\ohz, our fit yields \\teff=51126$\\pm 660$\\,K and \\logg=7.90$\\pm0.080$ (cgs) with anti-correlated errors (1-$\\sigma$) which include not only the statistical but also the systematic uncertainties of the fit. HZ43AB displays a previously detected bremsstrahlung component with a temperature k$T\\simeq 0.6$\\,keV.  For \\osi, we find \\teff=24923$\\pm 115$\\,K for fixed \\logg=8.6.  The calibration of the short-wavelength \\euve\\ spectrometer differs from that of the LETG+HRC-S by $15\\pm 7$\\%. The \\rosat\\ PSPC is found to be correctly calibrated within a few percent and reports of a major miscalibration are unfounded. } { % We have obtained improved parameters for \\orx, \\ohz, and \\osi\\ which fit the observations from the optical to the soft X-ray regime. Our approach allows us to quote their absolute spectral fluxes at selected wavelengths which may aid the calibration of other space-borne instruments.}  \\keywords {Methods: data analysis -- Stars: white dwarfs -- Stars: neutron -- Stars: fundamental parameters -- Stars: individual (\\ohz, \\osi, \\orx) -- X-rays: individual (\\ohz, \\osi, \\orx)}  \\titlerunning{Soft X-ray standard stars \\ohz, \\osi, and \\orx} \\authorrunning{K.~Beuermann et al.} ",
        "introduction": "Hot white dwarfs with pure hydrogen atmospheres are well established standard stars in the optical and ultraviolet \\citep[e.g. ][]{bohlinetal01} and have played a decisive role in the calibration of satellite instruments in all spectral regimes from the optical to the extreme ultraviolet and soft X-rays. The FOS and STIS on the \\textit{Hubble} Space Telescope \\citep{bohlin00}, the Extreme Ultraviolet Explorer (\\textit{EUVE}) \\citep{boydetal94}, and the Low Energy Transmission Grating Spectrometer on board \\textit{Chandra} (LETG+HRC-S) with the High Resolution Camera as detector \\citep{brinkmanetal00,peaseetal00} are prominent examples. The latter two instruments were calibrated on the ground, but their calibration was partly relinquished in favor of an in-flight calibration using hot white dwarfs. The Position Sensitive Proportional Counter (PSPC) on board the R\\\"ontgen Satellite (\\rosat), on the other hand, was accurately calibrated on the ground \\citep{truemper82,pfeffermannetal87} and not recalibrated in flight. Based on observed PSPC spectra of white dwarfs, however, \\citet{napiwotzkietal93,jordanetal94, wolffetal95,wolffetal96} claimed the PCPC calibration at soft X-ray energies to be incorrect by a factor of about two. Given this important role of white dwarfs it is deplorable that their soft X-ray fluxes are much less well documented than their optical and ultraviolet fluxes. Given, furthermore, the large body of scientific data of all types of objects potentially affected, the resolution of such open calibration issues is of general interest.  Our initial aim was to investigate the PSPC calibration which requires, however, to take recourse to spectrally higher resolved data as provided by the \\euve\\ or the \\chandra\\ LETG+HRC spectrometers. Since the effective area of the latter is itself uncertain by as much as 25\\% at intermediate wavelengths around 60\\AA\\ \\citep[e.g.][]{peaseetal03}, we found it necessary to perform a more detailed cross-calibration of the three instruments. As a result, we are able to determine accurate soft X-ray fluxes of the hot white dwarfs \\ohz\\ and \\osi\\ and of the neutron star \\orx\\ (henceforth \\rxj) and thereby establish them as standards in the soft X-ray regime\\footnote{We consider neither quasars and BL\\,Lacs because of their inherent variability nor supernova remnants because of their extended nature.}. Although one of the two white dwarfs (\\ohz) is a well known standard for wavelengths $\\lambda > 912$\\AA\\ \\citep{bohlinetal01} and the optical flux of the other (\\osi) is reasonably well measured \\citep{barstowetal05}, the extrapolation of the optically adjusted model spectra into the soft X-ray regime suffers from the wagging-tail problem. We solve this problem by tying the soft X-ray tail of the spectrum of \\ohz\\ to that of \\rxj.  This adds a second fix to the spectrum of \\ohz\\ besides that in the optical and effectively constrains its spectral shape.  The organization of the paper is as follows. In Sect.~2, we describe the analysis of the LETGS and PSPC data and list the EUVE archive data used. Sect.~3 is devoted to the description of our theoretical white dwarf spectra and summarizes the available information on \\teff\\ and \\logg\\ of \\ohz. Our principal results are provided in Sect.~4. It starts with information on the optical fluxes to which our model spectra are normalized (Sect.~4.1). In Sect.~4.2, we derive an accurate spectrum of \\rxj\\ that serves as the short-wavelength fix for \\ohz. In this context, we discuss interstellar absorption near the carbon K-edge at 43.7\\AA\\ to some detail. The essence of the paper is then presented in Sect.~4.4, the simultaneous fit of \\rxj, \\ohz, and \\osi\\ that results in an internally consistent set of photon spectra and simultaneously in the wavelength-dependent correction functions to the \\mbox{LETG+HRC-S} effective areas at wavelengths $\\lambda \\ge 44$\\AA. Here, we discuss also the systematic errors of our approach. The stellar parameters implied by our fit are presented and discussed in Sect.~4.5. Sect.~4.6 compares our LETGS calibration to that of the \\euve\\ short wavelength spectrometer and Sect.~4.7 finally demonstrates the accurate calibration of the \\rosat\\ PSPC. The conclusions are summarized in Sect.~5. ",
        "conclusions": "We have established the neutron star \\rxj\\ and the hot white dwarfs \\ohz\\ and \\osi\\ as standard stars in the soft X-ray regime by fitting the LETG+HRC spectra of all three stars simultaneously with the best available model spectra for the white dwarfs and a two-blackbody model for the neutron star. Our method ties the soft X-ray spectra of the two white dwarfs to that of the neutron star \\rxj\\ and thereby to be better calibrated effective areas of the \\chandra\\ LETG+HRC spectrometer shortwards of 40\\AA. The spectrum of \\ohz\\ (and indirectly that of \\osi) is thus fixed at soft X-ray \\emph{and} at optical wavelengths, a procedure which avoids the wagging-tail problem of previous calibration attempts using white dwarfs. Our second major result is the self-consistently determined correction to the LETG+HRC-S effective areas longwards of 44\\AA. We find that the areas for both dispersion directions need a maximum correction of about $-23$\\% at 65\\AA\\ and are correct to within $\\sim10$\\% at wavelengths exceeding 80\\AA. These corrections are valid relative to the short-wavelength effective areas. The largest remaining uncertainty, which is not addressed here, lies in the absolute calibration of the LETG+HRC-S at short wavelengths.  Our photospheric parameters of the three stars are in good agreement with those derived from independent optical/ultraviolet and soft X-ray studies and, in part, improve on them. The soft X-ray flux of \\ohz\\ can be produced by a range of \\teff-\\logg\\ combinations and is better defined than either of these two parameters.  We find that the calibration of the \\euve\\ short-wavelength spectrometer differs from that of the \\chandra\\ LETG+HRC-S by about $15\\pm 7$\\% in the sense that the original ground calibration seems to have been more adequate than the adopted in-flight calibration.  We demonstrate that the ground calibration of the \\rosat\\ PSPC is correct to within a few percent and that previous reports of a major miscalibration are unfounded. The observed PSPC spectra of our three stars agree within a few percent with the LETGS-based soft X-ray spectra folded through the PSPC response. This internal consistency, in turn, supports our adjustment of the \\mbox{LETG+HRC-S} calibration at long wavelengths and the basic correctness of the nominal calibration at short wavelengths."
    },
    "0608/astro-ph0608301_arXiv.txt": {
        "abstract": "We present optical and near-infrared images, \\HI\\ 21~cm emission maps, optical spectroscopy, and {\\em Hubble Space Telescope}/Space Telescope Imaging Spectrograph ultraviolet spectroscopy of the QSO/galaxy pair \\sbs/\\ic. The QSO sight line lies at a position angle of $27\\degr$ from the minor axis of the nearby dwarf starburst galaxy \\ic\\ ($\\czgal = 1204\\pm3$~\\kms, $\\LB \\sim 0.09$~\\Lstar, current star formation ${\\rm rate} = 0.08$--0.24\\hsq~\\Msun~\\peryr) and 33\\h~kpc ($6\\farcm6$) from its nucleus. We find that \\ic\\ has an \\HI\\ mass of $\\MHI = (3.6\\pm0.1)\\times10^8$~\\Msun\\ and a dynamical mass of $\\Mdyn = (3.1\\pm0.5)\\times10^{10}\\h$~\\Msun. The UV spectrum of \\sbs\\ shows a metal-line (\\Lya+\\CIV) absorber near the redshift of \\ic\\ at $\\czabs = 1110\\pm30$~\\kms. Since \\ic\\ is a dwarf starburst and the \\sbs\\ sight line lies in the expected location for an outflowing wind, we propose that the best model for producing this metal-line absorber is a starburst wind from \\ic. We place consistent metallicity limits on \\ic\\ ($[Z/\\Zsun] \\sim -0.7$) and the metal-line absorber ($[Z/\\Zsun] < -0.3$). We also find that the galaxy's escape velocity at the absorber location is $\\vesc = 80\\pm10$~\\kms\\ and derive a wind velocity of $\\vw = 160\\pm50$~\\kms. Thus, the evidence suggests that \\ic\\ produces an unbound starburst wind that escapes from its gravitational potential to transport metals and energy to the surrounding intergalactic medium. ",
        "introduction": "\\label{intro}  Metals have been distributed throughout much of the intergalactic medium (IGM) at all observed redshifts. Transport by galactic starburst winds is a leading theory for explaining the existence of these IGM metals, which are often found far \\citep[several hundred kiloparsecs;][]{stocke06b} from any sites of ongoing star formation. However, it is not clear whether starbursts associated with luminous massive galaxies or smaller but more numerous dwarf galaxies produce the winds that are primarily reponsible for enriching the IGM with metals and energy. Statistically, it appears that 0.1~\\Lstar\\ galaxies must contribute to IGM enrichment unless \\Lstar\\ galaxies can enrich regions up to $\\sim 1$~Mpc in radius \\citep{tumlinson05a}.  Starburst winds have been studied around nearby galaxies in emission using narrowband nebular emission line (e.g., \\Ha) and thermal X-ray continuum images \\citep*{watson84,martin99,martin02} as well as ultraviolet and optical absorption-line spectroscopy \\citep*{heckman00,heckman01,rupke05a,rupke05b,martin05,martin06,keeney05,keeney06}. Both imaging and absorption-line spectroscopy have limitations, but the two techniques are complementary in many ways. Emission-line studies can address the extent and morphology of the outflow, but only in its densest regions. These studies cannot measure the velocity or temperature of the outflow in the diffuse halo, so they are incapable of determining whether ejecta in the halo are bound to the galaxy. Thus, emission-line studies alone cannot determine which galaxies produce outflows that enrich the surrounding IGM. Absorption-line studies, on the other hand, are much more sensitive to diffuse gas and therefore can determine whether ejecta in the halo are bound to the galaxy, but only along one line of sight so they are unable to study the outflow morphology.  Most absorption-line studies also introduce an ambiguity in the distance between the background source and the absorbing gas since they typically use the stellar continuum of the starburst itself as their source \\citep[e.g.,][]{heckman00,heckman01,rupke05a,rupke05b,martin05,martin06}. Consequently, they cannot distinguish a high-velocity outflow in the galactic halo from one in the starburst region itself. Using outflow velocities measured in the galactic disk to predict whether the absorbing gas is gravitationally bound to the galaxy is problematic because ejecta within several kiloparsecs of the starburst region will decelerate under the combined effects of gravity and mass loading \\citep[e.g.,][]{suchkov96,martin02,strickland04}. This problem can be circumvented by probing the outflow well away from the starburst region, which is only feasible if there is a bright background QSO projected near the galaxy of interest \\citep[e.g.,][]{stocke04,keeney05,keeney06}.  Massive galaxies have high rates of metal production and thus large amounts of potential fuel for IGM enrichment. Luminous low-redshift starbursts have star formation rates (SFRs) in the range 1--10~\\Msun~\\peryr, although rates of $\\sim 100$--1000~\\Msun~\\peryr\\ can be triggered by mergers \\citep[\\citealp*{heckman90};][]{martin03,rupke05a,rupke05b}. These galaxies generate winds with average outflow velocities of 300--400~\\kms, but wind speeds of $\\gtrsim 1000$~\\kms\\ are not unheard of \\citep*{heckman00,filippenko92,veilleux94,veilleux05}. Only the most luminous galaxies can be detected at high redshift, so studies of the earliest epochs of star formation and metal enrichment are inherently biased toward the most massive starbursts \\citep*{steidel96a,steidel96b,steidel99,steidel01,pettini01,pettini02,adelberger03}.  Dwarf galaxies produce fewer metals per galaxy than their massive counterparts but they are much more numerous and their cumulative effects could dominate IGM enrichment \\citep{stocke04,tumlinson05a}. Dwarf starburst galaxies have typical SFRs of 0.1--1~\\Msun~\\peryr\\ \\citep{martin03} and produce winds with outflow velocities of $\\sim 50$--200~\\kms\\ \\citep{marlowe95,martin98,schwartz04}. However, starbursts of all luminosities have the same maximum areal SFR \\citep[$\\sim 45$~\\Msun~\\persqkpc~\\peryr;][]{meurer97} and produce winds with the same X-ray temperature \\citep{martin99}. These results indicate that dwarf starbursts may be more efficient than their massive counterparts at transporting the metals entrained in their winds to the IGM due to their shallower gravitational potentials with lower escape velocities.  Recent wind studies that utilize nearby QSOs to probe starburst winds several kiloparsecs from the host galaxy find that luminous galaxies produce bound winds while dwarf galaxies produce unbound winds. \\citet{keeney05} found that the nearby luminous starburst galaxy NGC~3067 ($L \\approx 0.5$~\\Lstar, ${\\rm SFR} \\approx 1.4$~\\Msun~\\peryr) produces a bound wind along the line of sight to 3C~232, which is located near the minor axis of NGC~3067 and 11\\h~kpc from the galactic plane. The Milky Way also produces a bound starburst wind along the lines of sight to two high-latitude AGN (Mrk~1383 and PKS~2005--489) at $l=350\\degr$ that probe the regions on either side of the Galactic Center at heights up to 12.5~kpc \\citep{keeney06}. On the other hand, \\citet{stocke04} found that a dwarf poststarburst galaxy produced an unbound wind $\\sim 3.5$~Gyr ago that is now observed as the 1586~\\kms\\ metal-line absorber in the spectrum of 3C~273, which is 71\\h~kpc away in projection on the sky.  In our on-going study to determine what type of galaxies enrich the IGM, this paper presents an ultraviolet spectrum of the QSO \\sbs, as well as optical and near-infrared images, an \\HI\\ 21~cm emission map, and an optical spectrum of the blue compact galaxy \\ic, which is $6\\farcm6$ away on the sky. These observations will be used to study whether the starburst wind produced by \\ic\\ can escape its gravitational potential to enrich the surrounding IGM. In \\S\\,\\ref{obs} we describe the acquisition and reduction of our multi-wavelength observations. The connection between the metal-line absorber in the spectrum of \\sbs\\ and \\ic\\ is discussed in \\S\\,\\ref{galabs}.  In \\S\\,\\ref{escape} we examine whether the starburst wind produced by \\ic\\ can escape from its gravitational potential. We summarize our results and discuss their implications for IGM enrichment in \\S\\,\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  We have used HST to detect \\Lya\\ and \\CIV\\ absorption ($\\czabs = 1110\\pm30$~\\kms) from the nearby dwarf starburst galaxy \\ic\\ ($\\czgal = 1204\\pm3$~\\kms) in the spectrum of the QSO \\sbs. Narrowband \\Ha\\ images show that \\ic\\ is currently forming stars at a rate of 0.08--0.24\\hsq~\\Msun~\\peryr, which is comparable to the SFRs found in other nearby dwarf starburst galaxies \\citep{martin03}. A long-slit optical spectrum of \\ic\\ reveals \\Hb, H$\\gamma$, [\\OII], [\\OIII], [\\NII], and [\\SII] emission in addition to \\Ha. The H-band surface brightness profile of \\ic\\ is well-fit by a \\citet{devaucouleurs48} $R^{1/4}$ profile, making it photometrically similar to elliptical galaxies despite its current starburst. An \\HI\\ 21~cm emission map obtained with the VLA shows that \\ic\\ is interacting with \\sdsslsb, a LSB galaxy to the north (see Fig.~\\ref{fig:HI}), which could have triggered its current epoch of star formation. Eliminating as much tidal material as possible by-eye, \\ic\\ has a \\HI\\ mass of $\\MHI = (3.6\\pm0.1)\\times10^8$~\\Msun\\ and its \\HI\\ 21~cm rotation curve (Fig.~\\ref{fig:rotcur}) implies that it has a dynamical mass of $\\Mdyn = (3.1\\pm0.5)\\times10^{10}\\h$~\\Msun.  Since \\ic\\ is a dwarf starburst galaxy and the \\sbs\\ sight line lies in the expected direction of a starburst wind \\citep{heckman90,veilleux05}, we suggest that the metal-line absorber at $\\czabs = 1110\\pm30$~\\kms\\ is caused by an outflowing starburst wind from \\ic. If this is indeed the case, then the wind produced by \\ic\\ will escape its gravitational potential to enrich the surrounding IGM since the escape velocity at the absorber location is $\\vesc = 80\\pm10$~\\kms\\ and the wind velocity is $\\vw = 160\\pm50$~\\kms. If the absorbing gas was ejected from \\ic\\ when it was forming stars at its current rate in a burst lasting $\\sim 10^8$~years, then it could accelerate $\\gtrsim 10^7$~\\Msun\\ of material to the current absorber velocity. A rough estimate of a subsolar metallicity ($[Z/\\Zsun] < -0.3$) for this absorber is also consistent with the metallicity limit of $[Z/\\Zsun] \\sim -0.7$ found for \\ic.  Our conclusion that \\ic\\ produces an unbound starburst wind agrees with the results of \\citet{stocke04}, who found that an unbound wind from a dwarf poststarburst galaxy could be responsible for the $\\czabs = 1586$~\\kms\\ metal-line system in the 3C~273 sight line 71\\h~kpc away. This dwarf poststarburst galaxy could be representative of a later stage in the evolution of the \\sbs/\\ic\\ system. Once all of the \\HI\\ in \\ic\\ has been exhausted via star formation and ejection it will develop a poststarburst spectrum and fade in luminosity as envisioned by \\citet{babul92} to the current brightness of the 3C~273 dwarf ($\\MB = -13.9$) within $\\sim 1$~Gyr after the \\HI\\ is exhausted \\citep{bruzual93,bruzual03,babul96}. Meanwhile, the unbound wind will continue to propagate into the surrounding IGM and increase the distance that \\ic\\ has distributed metals.  On the other hand, more luminous galaxies appear to produce bound winds. \\citet{keeney05} found that the nearby luminous (0.5~\\Lstar) starburst galaxy NGC~3067 (${\\rm SFR} \\approx 1.4$~\\Msun~\\peryr) produces a bound wind along the 3C~232 sight line, which probes the halo of NGC~3067 near its minor axis and 11\\h~kpc from the plane. Similarly, \\citet{keeney06} found that the Milky Way produces a bound wind toward two high-latitude AGN sight lines near $l = 0\\degr$ (Mrk~1383 and PKS~2005--489) that probe regions directly to the north and south of the Galactic Center at heights of up to 12.5~kpc. In both NGC~3067 and the Milky Way, the bound winds have the same spectral signature as high-velocity clouds and share many of their properties.  Statistical studies of QSO-absorber pairs in large samples of low-$z$ \\Lya+metal line absorbers \\citep[e.g.,][]{stocke06b} also suggest that IGM metals are spread primarily by dwarf galaxies. For example, \\citet{tumlinson05a} found that enriched regions of gas must extend $\\sim 1$~Mpc from \\Lstar\\ galaxies to be due to $\\geq \\Lstar$ galaxy superwinds. Enrichment regions of 100--150~kpc are much more plausible based upon observed absorber-galaxy distances \\citep{stocke06b}, but require that enrichment is due primarily to dwarf ($\\lesssim 0.1$~\\Lstar) galaxies \\citep{tumlinson05a}. The \\sbs/\\ic\\ system supports this conclusion.  \\ic\\ is not the only blue compact galaxy that shows evidence for an outflowing wind. The Fornax Cluster galaxy FCC~35 shows strong \\Ha\\ and [\\OIII] emission and an unusual single-dish \\HI\\ 21~cm profile in which a rotationally supported \\HI\\ disk is superimposed with an irregularly shaped \\HI\\ cloud with no optical counterpart \\citep{putman98}. The disk and the cloud have roughly the same \\HI\\ mass ($\\MHI = 2.2\\times10^8$~\\Msun) and overlap spatially but not in velocity, with the cloud blueshifted by $\\sim 150$~\\kms\\ with respect to the systemic velocity of the disk. While the cloud has no velocity structure, the rotation curve of the disk indicates that it has a truncated mass distribution. \\citet{putman98} suggest that the \\HI\\ cloud is triggering the current burst of star formation in FCC~35, but by analogy with \\ic\\ the \\HI\\ cloud could also be ejecta from the starburst wind of FCC~35, which would explain the truncated mass distribution of its disk. FCC~35 has an \\HI\\ mass comparable to the galaxies with the lowest gas masses and smallest gas depletion timescales in the \\citet{pisano01} sample of nearby blue compact galaxies. These low-mass galaxies may be undergoing their last star formation event and will likely fade in luminosity as envisioned by \\citet{babul92} once their current star formation ceases \\citep{pisano01}.  Collectively, these results imply that starburst winds escape more easily from dwarf starburst galaxies than their more massive counterparts. This is to be expected since \\citet{meurer97} found a maximum areal SFR of $\\sim 45$~\\Msun~\\persqkpc~\\peryr\\ for starbursts of all luminosities and \\citet{martin99} found that the temperature of starburst winds is nearly constant as a function of galaxy mass, both of which imply that the strength of a starburst wind is independent of galaxy size. Thus, there is growing evidence that the metals and energy expelled by massive galaxies are retained in their bound halos and the weaker but more numerous dwarfs are primarily responsible for enriching the IGM.  We have argued that an outflowing unbound starburst wind from \\ic\\ is the best model for explaining the origin of the metal-line absorber in the spectrum of \\sbs\\ due to: (1) the proximity of the QSO sight line to the galaxy's minor axis, (2) the on-going star formation in \\ic, (3) the energetic ability of the current \\ic\\ starburst to eject $\\sim 10^7$~\\Msun\\ of material to the observed wind velocity after escaping the galaxy's gravitational potential well, and (4) the consistent subsolar metallicities derived for \\ic\\ and the metal-line absorber. However, other plausible models for the absorber origin exist. In particular, we cannot rule out the possibility that the absorber is caused by diffuse tidal debris from the recent interaction of \\ic\\ and \\sdsslsb, although the observed geometry of the \\HI\\ 21~cm emitting gas does not obviously support this hypothesis. Future observations of \\sbs\\ as well as fainter AGN near \\ic\\ with the Cosmic Origins Spectrograph (COS) would allow us to distinguish between these models and sensitively search for absorption in other ions at $cz \\sim 1110$~\\kms.   {"
    },
    "0608/astro-ph0608137_arXiv.txt": {
        "abstract": "We extend the study of the performance of a prototype two-phase liquid xenon WIMP dark matter detector to recoil energies below 20 keV.   We demonstrate a new method for obtaining the best estimate of the energies of events using a calibrated sum of charge and light signals and introduce the corresponding discrimination parameter, giving its mean value at 4~kV/cm for electron and nuclear recoils up to 300 and 100~keV, respectively.  We show that fluctuations in recombination limit discrimination for most energies, and reveal an improvement in discrimination below 20 keV due to a surprising increase in ionization yield for low energy electron recoils.  This improvement is crucial for a high-sensitivity dark matter search. ",
        "introduction": "\\label{intro}  Two-phase time-projection-chambers are a powerful technique for combining the high mass of a liquid-phase detector with the highly-sensitive gas-phase readout of electrons via proportional scintillation light \\cite{Bolozdynya:1999dualphase}.   Recently, there has been a large effort to develop Xe and Ar two-phase detectors to search for WIMP dark matter \\cite{Smith:2003zeplin,Kudryavtsev:2003zeplin,Yamashita:2003xmassdm,Brunetti:2004warp,CaseBrownColumbia:2006NucRec}, including by the XENON collaboration \\cite{Aprile:2002xenon}, in which our group is a member.  These detectors consist of a liquid target with a gas region above and grid structures to create electric fields which drift electrons to the liquid surface, extract them into the gas phase, and induce them to create proportional scintillation (``S2'').   The initial (``S1'') scintillation light from the event site and the S2 light are measured by PMTs; the time difference between these signals gives the event depth.  A key feature of these detectors is discrimination between electron recoils (from radioactive backgrounds) and nuclear recoils (from WIMPs) based on the different amounts of charge recombination for these two types of interactions.  This was recently demonstrated in prototype detectors by our group and by the Columbia and Brown XENON groups \\cite{CaseBrownColumbia:2006NucRec}.   This note extends these measurements to below 5~keV, explores the dominant role played by fluctuations in recombination, and reveals a surprising improvement in discrimination below 20 keV.  Our detector has an active liquid region 3.7 cm $\\O$ and 0.96 cm deep, with stretched wire grids to define the drift field and a 10 kV/cm field region in the gas phase.  Two 178~nm xenon scintillation-light sensitive PMTs  (Hamamatsu 9288) are located below and above the active volume.  We estimate a light collection efficiency of $\\sim$50\\%, resulting in $\\sim$1~photoelectrons (pe) per keV for nuclear recoils ($\\sim$5~pe/keV for electron recoils at zero field).   The trigger threshold was $\\sim$5 electrons (in the S2 signal), with $\\sim$50\\% acceptance for S1 single photo electrons.  The detector employs a temperature-controlled, liquid nitrogen coldfinger-cooled cryostat, with $\\leqslant \\pm0.2$~K stability.  The level of the liquid surface was measured using three parallel plate capacitors, allowing alignment of the liquid level and grids to $\\sim 0.1^{\\circ}$.   The detector was calibrated with 122~keV gammas from $^{57}$Co, and discrimination measurements were made using gammas from \\(^{133}\\)Ba and neutrons from a 25 $\\mu$Ci \\(^{252}\\)Cf fission source.   Inelastic neutron reactions gave a prompt 40~keV line from  \\(^{129}\\)Xe.  Purity of the Xe, maintained with a heated getter purifier, was measured to $\\lesssim$3\\% charge loss over 1~cm via the depth dependence of the charge signal for 122 keV gammas. ",
        "conclusions": ""
    },
    "0608/astro-ph0608651_arXiv.txt": {
        "abstract": "{This work is part of a systematic X-ray survey of the Taurus star forming complex with XMM-Newton.} {We study the time series of all X-ray sources associated with Taurus members, to statistically characterize their X-ray variability, and compare the results to those for pre-main sequence stars in the Orion Nebula Cluster and to expectations arising from a model where all the X-ray emission is the result of a large number of stochastically occurring flares.} {The analysis of the lightcurves is based on a maximum likelihood algorithm that segments the time series in intervals of constant signal without the need of binning. Flares are defined with criteria that take into account the amplitude and the derivative of the segmented lightcurves. Variability statistics are evaluated for different classes of pre-main sequence stars (protostars, cTTS, wTTS, brown dwarfs), and for different spectral type ranges. Flare frequency and energy distribution are computed.} {We find that roughly half of the detected X-ray sources show variability above our sensitivity limit, and in $\\sim 26$\\,\\% of the cases this variability is recognized as flares. Variability is more frequently detected at hard than at soft energies. The variability statistics of cTTS and wTTS are undistinguishable, suggesting a common (coronal) origin for their X-ray emission. The frequency of large flares ($E > 10^{35}$\\,erg) on Taurus members is $1$ event per star in $800$\\,ksec. The typical duration of these flares -- probably biased by the finite observing time -- is about $10$\\,ksec. We have for the first time applied a rigorous maximum likelihood method in the analysis of the number distribution of flare energies on pre-main sequence stars. In its differential form this distribution follows a power-law with index $\\alpha = 2.4 \\pm 0.5$, in the range typically observed on late-type stars and the Sun. } {The signature of the X-ray variability in the pre-main sequence stars in Taurus and Orion provides twofold support for coronal heating by flares: (i) The correlation between the maximum variability amplitude and the minimum emission level indicates that both flare and quiescent emission are closely related to the coronal heating process. (ii) The power-law index $\\alpha$ derived for the flare energy distribution is large enough to explain the heating of stellar coronae by nano-flares ($\\alpha > 2$), albeit associated with a rather large uncertainty that leaves some doubt on this conclusion. } ",
        "introduction": "\\label{sect:intro}  X-ray variability is a characteristic property of magnetically active stars, including such diverse objects as the Sun, dMe flare stars, RS\\,CVn variables, and pre-main-sequence (pre-MS) T Tauri stars \\citep{Feigelson99.1, Favata03.1, Guedel04.1}. Variability on long time-scales may comprise modulation by rotating active regions (with time-scales on the order of days corresponding to the rotation period of the star) or even activity cycles (with expected time-scales on the order of years); see \\citet{Marino03.1, Favata04.1} for examples. Most of the observed variations have short time-scales (on the order of hours), are accompanied by an increase in electron temperature, and are therefore attributed to flares resulting from magnetic reconnection events.  The intense X-ray radiation from flares is ascribed to the filling of magnetic loops with heated gas that has been driven from the chromospheric layers along the magnetic field lines into the corona \\citep[the `chromospheric evaporation scenario', e.g.,][]{Antonucci84.1}. Various authors have therefore speculated that the overall X-ray emission seen in magnetically active stars is in fact the result of a large number of stochastically occurring flares. In analogy to the `microflare heating' hypothesis in solar physics \\citep[see e.g.][]{Hudson91.1}, these events are believed to be distributed in energy such that the large majority releases only small amounts of energy, making their individual identification in light curves impossible while we only measure the integrated emission \\citep[e.g.,][]{Audard00.1, Guedel02.1, Guedel03.1}.  Extensive studies yielding statistical properties of X-ray variability for different classes of active stars have long been impeded, mostly due to observational restrictions deriving from insufficient monitoring time. Further limitations, besides sensitivity, are imposed by the poor temporal coverage of X-ray instruments in low-Earth orbits that result in frequent occultations of the target by the satellite revolution. {\\em XMM-Newton} and {\\em Chandra} have opened the field for systematic investigations on time-scales typical for magnetic activity. This is due to the eccentric orbits of both satellites, enabling uninterrupted observations of up to $\\sim 2$\\,d. Even then, variability studies are difficult, because in practice it is nearly impossible to maximize both sample size and exposure time per target. An outstanding study in this respect is the nearly 13\\,d long observation of the Orion Nebula Cluster (ONC) by {\\em Chandra}, termed the {\\em Chandra} Orion Ultradeep Project (COUP). The COUP yielded not only the longest uninterrupted X-ray time series for pre-MS stars so far, but owing to the dense population of the ONC it allowed also to examine in a homogeneous way the X-ray variability of the largest sample of pre-MS stars: More than $1600$ X-ray sources were detected in a single {\\em Chandra} field, mostly members of the ONC. Amongst others, the COUP resulted in a systematic study of flares on `young Suns' \\citep{Wolk05.1}, rotational modulation \\citep{Flaccomio05.1}, the geometry of flare loops \\citep{Favata05.1}, and X-ray variability of hot stars \\citep{Stelzer05.1}. %  In contrast to the ONC, the Taurus star forming complex spans a large angle on the sky \\citep[see Fig.~1 in ][]{Guedel06.1}. This implies that observing a significant fraction of its stellar members with a telescope operating in pointing mode requires many different exposures. Realistically, the whole population can be sampled only by an all-sky monitoring. \\citet{Neuhaeuser95.1} have used the data obtained during the {\\em ROSAT} All-Sky Survey (RASS) to study the X-ray emission of the pre-MS stars in Taurus. However, the RASS consisted in snapshots of a few seconds each separated by $\\sim 90$\\,min, corresponding to the Earth revolution of the {\\em ROSAT} satellite, and is inappropriate for timing studies. A systematic study of the X-ray variability of the pre-MS stars in Taurus-Auriga based on more than $100$ {\\it pointed} {\\em ROSAT} PSPC observations was presented by \\citet{Stelzer00.1}. That investigation yielded the first estimate of flare rates for these stars, which were found in outburst during about $1$\\,\\% of the observing time. Various uncertainties compromise this measurement, e.g. the {\\em ROSAT} data set % was (i) inhomogeneous in sensitivity due to the vastly different exposure times of the individual observations, (ii) affected by Earth blocks due to the satellite's revolution, and (iii) limited to soft energies.  Here we present a variability study for the pre-MS stars in the Taurus star forming region based on extensive {\\em XMM-Newton} observations. We discuss data from the {\\em XMM-Newton Extended Survey of the Taurus Molecular Cloud} (XEST). This project was devised to comprise observations with roughly uniform and uninterrupted exposures of $\\sim 30 - 40$\\,ksec. Thus, it represents a major improvement with respect to previous X-ray data of the Taurus Molecular Cloud (TMC). A further advantage is the extension of the hard energy band beyond $2$\\,keV. A detailed description of the XEST observations is found in \\citet{Guedel06.1}. Here we present a systematic time series analysis of the pre-MS population detected in the XEST.  In Sect.~\\ref{sect:analysis} the methods used in the analysis of the time series are described, and the results are summarized. Sect.~\\ref{sect:var_quant} deals with a quantitative investigation of the lightcurves. The relatively short duration of the XEST observations imply that the most likely cause of the variability is flaring. We introduce our way of defining flares and present the results from the flare detection process in that section. Variability statistics for different classes of pre-MS stars in the TMC are given in Sect.~\\ref{sect:results_var}, including a discussion of observational biases. The frequency and the energy distributions of the detected flares are examined in Sect.~\\ref{sect:flare_tau} and Sect.~\\ref{sect:flare_energy}. Sect.~\\ref{sect:ctts_wtts} presents a detailed comparison of the variability characteristics of cTTS and wTTS. We summarize our results in Sect.~\\ref{sect:conclusions}. ",
        "conclusions": ""
    },
    "0608/astro-ph0608467_arXiv.txt": {
        "abstract": "We quantify the effects of electron thermal conduction on the properties of hot accretion flows, under the assumption of spherical symmetry. Electron heat conduction is important for low accretion rate systems where the electron cooling time is longer than the conduction time of the plasma, such as Sgr A* in the Galactic Center.  For accretion flows with density profiles similar to the Bondi solution ($n[r] \\propto r^{-3/2}$), we show that heat conduction leads to super-virial temperatures, implying that conduction significantly modifies the dynamics of the accretion flow.  We then self-consistently solve for the dynamics of spherical accretion in the presence of saturated conduction and electron heating.  We find that the accretion rate onto the central object can be reduced by $\\sim 1-3$ orders of magnitude relative to the canonical Bondi rate.  Electron conduction may thus be an important ingredient in explaining the low radiative efficiencies and low accretion rates inferred from observations of low-luminosity galactic nuclei.  The solutions presented in this paper may also describe the nonlinear saturation of the magnetothermal instability in hot accretion flows. ",
        "introduction": "The conversion of gravitational binding energy of accreting matter into radiation is thought to account for the high luminosities observed from many compact objects (e.g., AGN; \\citealt{lb69}). Most of the time, however, systems that appear to harbor compact objects (low luminosity AGN, X-ray binary systems in quiescence) do not radiate at the levels one would expect from simple estimates of the efficiency at which binding energy is converted to radiation. This implies that these systems have low radiative efficiencies and/or low accretion rates, and that the accreting gas may thus retain most of its binding energy in the form of heat. The ultimate energy sink for such hot radiatively inefficient accretion flows (RIAFs) remains an area of active research. The energy may be primarily advected in with the flow \\citep{ny94}, or it may be directed into driving bulk convective motion \\citep{qg00a, nar00} or a global outflow \\citep{bb99}.  Numerical simulations favor the latter possibilities (e.g., \\citealt{ia99}; \\citealt{sp01}; \\citealt{hb02}).  Another possible outlet for the accretion energy is a heat flux due to thermal conduction.  In order to maintain a hot accretion flow, the timescale for electrons and protons to exchange energy by Coulomb collisions exceeds the timescale for the plasma to flow into the central black hole (e.g., Rees et al. 1982).  Thus the inflowing plasma is collisionless and energy transport by heat conduction may be dynamically important.  Direct observational estimates of the electron mean-free path on scales resolved by {\\it Chandra} in nearby galactic nuclei -- including our Galactic Center -- give values $\\sim 0.02-1.3$ times the gravitational radius of influence of the central black hole ($\\sim 10^5-10^6$ Schwarzschild radii).  Thus the plasma is relatively collisionless even at large distances from the black hole \\citep{q04, tm06}.  Under these conditions, the magnitude of the heat flux is uncertain and depends on the geometry of the magnetic field and the effective collision rate due to plasma waves and instabilities.  It is plausible, however, that thermal conduction will occur at close to its maximal saturated rate, with a heat flux $F \\propto \\rho v_e^3$, where $\\rho$ is the density and $v_e$ the electron thermal speed \\citep{cm77}. This appears to be true in the solar wind \\citep{salem03}, a well-studied example of a macroscopically collisionless plasma.  For non-relativistic electrons, the ratio of the inflow time to the electron conduction time in a hot accretion flow may be as large as $\\sim v_e/v_r$ (see equation [\\ref{ENTe}] below), where $v_r$ is the radial flow velocity. For subsonic inflow (with respect to the ions), this ratio is $\\gtrsim \\sqrt{m_p/m_e} \\sim 40$, highlighting the possible importance of conduction in hot accretion flows.  In addition, \\citet{bal00} has shown that low collisionality conducting plasmas are unstable to convective-like motions in the presence of an outwardly decreasing temperature gradient (the MTI or magnetothermal instability).  Although the interplay between the MTI and the magnetorotational instability (MRI) is uncertain, it is plausible that the MTI will lead to a dynamically important heat flux in hot accretion flows.  Any significant heat flux from the inner (hotter) to the outer (colder) parts of an accretion flow can act to reduce the inflow rate. This is true if energy is transported by radiation (e.g., \\citealt{owym76}), viscous stresses (e.g., \\citealt{bb99}), convection (e.g., \\citealt{qg00a}), or thermal conduction, the focus of this paper.  Gas in the outer region of a spherical accretion flow near the sonic point has a thermal energy comparable to its gravitational binding energy.  Additional heating will drive the temperature of the gas above the virial temperature. If the gas is unable to cool efficiently, this increase in temperature will lead to a decrease in accretion rate by either decreasing the radius of the sonic point or by driving a thermal outflow.  Motivated by these considerations, this paper assesses the importance of thermal conduction using a simple one-dimensional (spherical) model of hot accretion flows with saturated conduction.  \\cite{tm06} have carried out a related analysis and find that the accretion flow can spontaneously produce thermal outflows driven in part by conduction (as in early models of the solar wind).  Their analysis is two-dimensional but self-similar in radius, whereas ours is one-dimensional but does not assume self-similarity; we also treat saturated heat conduction somewhat differently than they do (see \\S\\ref{EE}).  Tanaka \\& Menou's assumption of self-similarity enforces a density profile that varies as $r^{-3/2}$, whereas simulations of RIAFs consistently find density profiles shallower than this (e.g., \\citealt{spb99}, \\citealt{ia99}, \\citealt{sp01}, \\citealt{hb02}, \\citealt{ina03}).  We begin in \\S\\ref{EE} by outlining the energy equation in the presence of saturated heat conduction.  For completeness we consider both the ion and electron energy equations, though our focus will be on the electrons throughout this paper; we also quantify when electron heat conduction is likely to be important as an energy transport mechanism relative to radiative cooling.  We then present order of magnitude arguments that demonstrate the importance of heat conduction in hot accretion flows.  In particular, we solve for the electron temperature in the presence of heat conduction under the simplifying assumption that the density profile of the accretion flow is a fixed power-law in radius (\\S\\ref{FDP}).  In \\S \\ref{section:bondi} we relax the power-law density profile assumption of \\S \\ref{FDP} and self-consistently solve for the density, temperature, and radial velocity of the flow neglecting angular momentum; this is a version of the classic \\citet{bon52} accretion problem.  We summarize and discuss the implications of our findings in \\S \\ref{DISC}. ",
        "conclusions": "\\label{DISC}  We have shown that electron heat conduction can have a significant effect on the properties of radiatively inefficient accretion flows (RIAFs).  We considered two model problems to demonstrate this, both of which assume spherical symmetry.  First, for the canonical Bondi problem, we find that if a few percent of the accretion power is transferred to large radii by heat conduction, the heat flux can unbind the gas at large radii (eq. [\\ref{trbondi}]).  Generalizing this calculation to allow for a power-law density profile with $n(r) \\propto r^{-p}$ (\\S \\ref{FDP} and Figs. \\ref{f1}-\\ref{f3}), we solved for the electron temperature profile in the presence of saturated heat conduction; $p = 3/2$ is appropriate for the spherical Bondi problem, while $p \\sim 1/2-1$ is typically found in numerical simulations of rotating radiatively inefficient accretion flows.  Our calculations show that for density profiles with $p \\gtrsim 1$, and for plausible values of the electron conductivity and electron heating, the electron temperature is super-virial at large radii (see Fig.~\\ref{f1}-\\ref{f3}).  This indicates that the effects of conduction on the dynamics of the accretion flow must be self-consistently taken into account.  Motivated by this initial calculation, we solved the problem of spherical accretion in the presence of heat conduction and electron heating.  For appreciable electron heating rates ($f_e \\sim 0.1$) and dimensionless conductivities $\\alpha_c \\sim 0.1-1$ (see \\S \\ref{EE} for the details of the electron energy equation), the accretion rate is reduced by $\\sim 1-3$ orders of magnitude relative to the canonical Bondi rate (\\S\\ref{section:bondi}; Fig.~\\ref{f8}).  Physically, this is because some of the energy dissipated as heat is conducted to large radii where the increase in pressure stifles the inflow of matter. This decrease in $\\dot M$ relative to the Bondi accretion rate is consistent with observations of the Galactic Center, which indicate that the accretion rate is $2-3$ orders of magnitude lower than the Bondi rate \\citep{qg00b,bower,marrone}.  As discussed in \\S\\ref{EE}, electron heat conduction is only important at very low accretion rates $\\lesssim 10^{-5}-10^{-6} \\, \\dot M_{\\rm EDD}$, when the cooling time of the electrons via synchrotron radiation is longer than the conduction time in the plasma (eq. [\\ref{MDUL}]). Such accretion rates are appropriate for very low-luminosity AGN such as Sgr A* in the Galactic Center, massive black holes in early type galaxies (e.g., \\citealt{TD}), and some X-ray binaries in quiescence. For higher accretion rates, the electrons cool more rapidly than energy can be redistributed by conduction, and so electron conduction is unimportant. Ion heat conduction could in principle be relevant at all accretion rates since the ion cooling time in RIAFs is always longer than the inflow time (by construction); this is, however, harder to quantify because the ion conduction time is at best comparable to the inflow time of the plasma.  All of the calculations in this paper assume spherical symmetry.  This is formally appropriate only in the absence of significant angular momentum, when the circularization radius of the inflowing material is small compared to the gravitational sphere of influence of the central object.  We suspect, however, that the calculations presented here are relevant to rotating accretion flows as well, given the importance of energy transport in determining the structure of RIAFs (e.g., \\citealt{bb99}).  In a rotating accretion disk, however, the conducted energy may be transported to the surface of the disk, rather than to large radii.  This conductive flux might contribute to driving a thermal outflow. Indeed, the two-dimensional calculations of \\cite{tm06} indicate that conduction helps drive a bipolar outflow. Numerical simulations of RIAFs with heat conduction are necessary to further quantify the importance of conduction and to assess whether any resulting heat flux is primarily radial or vertical.  Existing simulations of RIAFs -- which do not include thermal conduction -- find density profiles with $p \\approx 1/2-1$, in which case our calculations indicate that conduction is unlikely to lead to super-virial temperatures which would be dynamically important (see Figs.~\\ref{f1}-\\ref{f3}).  Even in this case, however, conduction would be important for determining the electron temperature profile and thus the observed radiation.  Throughout this paper we have treated the dimensionless conductivity $\\alpha_c$ as a free parameter in order to explore the effects of conduction in RIAFs.  A precise determination of $\\alpha_c$ in this context requires understanding the effective conductivity in a magnetized, turbulent plasma.  Wave-particle scattering by high frequency turbulence can limit the mean free path of particles in low collisionality plasmas, and thus the conductivity.  In the Appendix, we argue that instabilities generated by the nonlinear evolution of the MRI in RIAFs lead to an upper limit on the ion conductivity of $\\alpha_c \\lesssim 0.1$, in which case ion conductivity is unlikely to be dynamically important.  We find that there is no corresponding limit on the electron conductivity, although a large heat flux itself may be self-limiting by generating whistler instabilities (e.g., Pistinner \\& Eichler 1998).  Independent of wave-particle scattering, electron transport perpendicular to a static magnetic field is suppressed due to the small value of the electron Larmor radius. However, the rapid separation of neighboring field lines in a turbulent medium results in an effective isotropic conductivity comparable to the field-free conductivity \\citep{jok73,nm01,cl04,cm04,mcb04}. This suggests that electron heat conduction will by dynamically important even in the presence of a tangled magnetic field, with $\\alpha_c \\sim 0.1-1$.  With this value of the electron conductivity, we find that conduction significantly modifies the dynamics of Bondi accretion (see Figs. \\ref{f7}-\\ref{f8}).  \\citet{bal00} has shown that low collisionality conducting plasmas are unstable to convective-like motions in the presence of an outwardly decreasing temperature gradient (the MTI).  In a non-rotating atmosphere, \\cite{ps05} demonstrate that the MTI leads to magnetic field amplification and a substantial heat flux down the temperature gradient.  It is natural to speculate that our solutions represent the non-linear saturation of the MTI in accretion disks.  However, simulations of the MTI in RIAFs are necessary to test this hypothesis since it is not clear what effect differential rotation and the MRI will have on the results of \\cite{ps05}.  The MTI will also be important in purely spherical accretion, as indicated by Figure~\\ref{f9}, which shows the growth rate of the MTI in units of the inverse infall time calculated from one of our spherical accretion solutions with $f_e = 0$ and $\\alpha_c = 0.1$.  As Figure ~\\ref{f9} shows, the MTI will be particularly important at large radii and will likely significantly modify existing results on the impact of magnetic fields on spherical accretion (e.g., those of \\citealt{in02}).  Although our focus in this paper is on the electron energetics in RIAFs, we have also carried out some preliminary calculations including the ions (neglecting ion conduction for reasons discussed above and in \\S \\ref{EE}).  For spherical accretion with significant ion heating, the accretion rate is reduced relative to the Bondi rate by even more than is shown in our electron-only calculations in Figure \\ref{f8}.  This is consistent with the arguments in \\S \\ref{section:bondi}. The ions are non-relativistic throughout most of the flow, and therefore have $\\gamma_i \\simeq 5/3$. In the presence of ion heating, the effective adiabatic index of the flow is $\\gtrsim 5/3$, in which case there is no transonic Bondi solution in Newtonian gravity.  Instead, the inflowing matter is in approximate hydrostatic equilibrium throughout the bulk of the flow and the solution never achieves the free-fall density profile of $\\rho \\propto r^{-3/2}$ (see the discussion around eq.  [\\ref{GAMC}]). The further reduction in accretion rate due to the inclusion of ion heating (even absent ion conduction), while strengthening our conclusions, is a consequence of the singular nature of the $\\gamma_i \\simeq  5/3$ Bondi solution, for which $\\rS \\simeq 0$. Including additional physical effects such as rotation will partially remove this singular behavior and may be required to realistically assess the impact of ion heating and conduction on the dynamics of hot accretion flows.  As noted in the Introduction, radiation emitted from close to the accreting object can play a role analogous to that of conduction in heating up the outer region of the accretion flow to super-virial temperatures. This mechanism has been studied in the context of RIAFs by \\cite{po98,po99,po01}. The accretion flow in that case can exhibit time dependence due to the lack of a steady-state solution or to thermal instability \\citep{owym76,cos78,bb80,kl83}. We do not expect time dependence in our problem, since our calculations show that a steady-state solution does exist and conduction tends to stabilize thermal instabilities \\citep{kl83}; a time-dependent calculation or a stability analysis would, however, be required to confirm this.  We also note that our calculations are only relevant for systems with sufficiently low luminosities as to be unaffected by radiative preheating (as discussed above, our solutions require low luminosities so that a significant fraction of the accretion energy can be conducted outwards rather than being radiated away)."
    },
    "0608/astro-ph0608184_arXiv.txt": {
        "abstract": "Model selection aims to determine which theoretical models are most plausible given some data, without necessarily asking about the preferred values of the model parameters. A common model selection question is to ask when new data require introduction of an additional parameter, describing a newly-discovered physical effect. We review several model selection statistics, and then focus on use of the Bayesian evidence, which implements the usual Bayesian analysis framework at the level of models rather than parameters. We describe our \\emph{CosmoNest} code, which is the first computationally-efficient implementation of Bayesian model selection in a cosmological context. We apply it to recent WMAP satellite data, examining the need for a perturbation spectral index differing from the scale-invariant (Harrison--Zel'dovich) case. ",
        "introduction": "Cosmologists are becoming very good at determining the parameters of the Universe. Within the last few years observational results, exemplified by the microwave background anisotropy measurements by the Wilkinson Microwave Anisotropy Probe (WMAP), have introduced precision into cosmological modelling. Considerable sophistication is now required both in deriving theoretical predictions from models and in carrying out data analysis procedures able to squeeze the best from the data.  Within a cosmological model, the parameters indicate the importance of different effects. For instance, they describe the relative amounts of different types of material in the Universe, the geometry and expansion rate of the Universe, and the properties of the initial irregularities in the Universe which led to the formation of structure. Such parameters are not predicted by fundamental theories, but rather must be fit from data in order to decide which combination, if any, is capable of describing our Universe. A variety of cosmological data are currently well fit by a model of the Universe that is homogeneous, isotropic and spatially flat, contains cold dark matter in greater proportions than ordinary baryonic matter, and in which tiny initial perturbations evolved under Einstein's theory of General Relativity into the structures of today. Current data constrain the parameters of this model rather well, many at the 10\\% level or better.  However, the presence of good data leads to a different problem, one of knowing when to stop fitting. Two different and competing models of the Universe may explain the data equally well, so how do we choose between them? The solution is one proposed by William of Occam, that the simpler model should be preferred. This is known as \\emph{Occam's razor}. So a complicated model that explains the data slightly better than a simple one should be penalized for the extra parameters it introduces, because the extra parameters bring with them a lack of predictability. On the other hand, if a model is too simple, and cannot fit certain data well, then it can be discarded.  This is a rather common type of statistical problem, both in cosmology and in other fields of astrophysics: each available parameter within a model describes some piece of physics that might be relevant to our Universe, but until measurements are made we don't know which.  Cosmological \\emph{model selection} refers to comparing different model descriptions of the data. It doesn't care particularly about the actual values of parameters, but rather aims to determine which {\\em set} of parameters gives the preferred fit to observational data. ",
        "conclusions": "Many of the most interesting cosmological questions are ones of model selection, not parameter estimation.  With the growing precision of cosmological data, it is imperative to deploy proper model selection techniques to extract the best robust conclusions from data.  Application to the post-WMAP3 cosmological data compilation continues to indicate that the data can be well fit by quite minimal cosmological models. Five fundamental parameters are definitely required, and WMAP3 has provided suggestive indications that a sixth, the density perturbation spectral index, needs to be added to the set. According to the Bayesian evidence, however, the case for inclusion of the spectral index has yet to become compelling.  As the data improve in sensitivity we expect new model selection based questions to be both raised and answered in the next decade.  These may be about the nature of dark energy, the model for reionization, the nature of inflation, the case for primordial gravitational waves, or the nature of cosmic topology.  Model selection, of course, will have further applications in astrophysics and beyond."
    },
    "0608/astro-ph0608478.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {} % aims heading (mandatory) {We have estimated the age of the young moving group TW Hydrae Association, a cohort of a few dozen stars and brown dwarfs located near the Sun which share the same kinematic properties and, presumably, the same origin and age.} % methods heading (mandatory) {The chronology has been determined by analyzing different properties (magnitudes, colors, activity, lithium) of its members and comparing them with several well-known star forming regions and open clusters, as well as theoretical models. In addition, by using medium-resolution optical spectra of two M8 members of the association (2M1139 and 2M1207 --an accreting brown dwarf with a planetary mass companion), we have derived spectral types and measured  H$\\alpha$ and lithium equivalent widths. We have also estimated their     effective temperature and gravity, which were used to produce an independent age estimation for these two brown dwarfs. We have also  collected spectra of  2M1315, a candidate member with a L5 spectral type and measured its H$\\alpha$ equivalent width.} % results heading (mandatory) {Our age estimate for the association, 10$^{+10}_{-7}$ Myr, agrees with previous values cited in the literature. In the case of the two brown dwarfs, we have derived an age of  15$^{+15}_{-10}$ Myr,  which also agree with our estimate for the whole  group.} % conclusions heading (optional), leave it empty if necessary {We compared our results with recent articles published on the same subject using other techniques, and  discuss the limits of the age-dating techniques.} ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%    If we want to put any astronomical phenomenon in context, we need to know its evolutionary status, i.e. its age. This is even more important when the phenomenon is a benchmark in a given discipline. This is the case  of the TW Hydrae Association or TWA (de la Reza et al.1989; Kastner et al. 1997), a young, nearby moving group which includes a few dozen stars and brown dwarfs (see the recent review by Zuckerman \\& Song 2004). The age has been estimated  in the range 8-20 Myr (Kastner et al. 1997; Stauffer,  Hartmann \\& Barrado y Navascu\\'es 1995; Soderblom et al.  1998; Hoff et al. 1998; Weintraub  et al. 2000; Makarov \\& Fabricius 2001; Makarov et al. 2005; de la Reza, Jilinski, Ortega 2006), and the individual  distances,  from Hipparcos, are  in the range 47-67 pc (ESA 1987). Due to these  characteristics,  the TWA is becoming a reference for such phenomena as accretion and circumstellar disks, rotation and activity,  the formation of planetary systems and the direct detection of planets.  Regarding the latter issue,  members of the association have been the targets of several searches for planetary companions, using AO techniques. Of those studies which have been  published, the most comprehensive one has been that by Masciadri et al. (2004), who observed 30 targets with VLT/NACO, including six stellar members  of TWA, with no positive result. However,  Chauvin et al.  (2004) announced the discovery of a candidate companion of the brown dwarf 2M1207334-393254 (2M1207 hereafter), which has a mass estimate within the planetary domain (5 M$_{jup}$).  One year later, the same team (Chauvin et al. 2005) measured the proper  motion and  verified that the brown dwarf and the companion were moving together, confirming  the nature of this object, the first planet whose detection had been achieved by direct imaging.  The 2M1207 system is full of surprises, since the brown dwarf is still undergoing active accretion as shown by Mohanty, Jayawardhana \\& Barrado y Navascu\\'es (2004) from high--resolution optical spectra around the H$\\alpha$ lines. It harbors a  circumsubstellar disk, seen as a mid-IR excess by Sterzik et al. (2004), which can be explained   by optically thick dust disk models. Rioz, Gizis \\& Hmiel (2006) have also derived the IR excess by using Spitzer data. However, no radio emission has been detected (Osten \\& Jayawardhana 2006). Mohanty (2006, priv. comm.) has suggested that the secondary --the planetary mass companion-- might have a disk too. Age is a key parameter in  understanding the nature of the faint companion and the properties of the system.   In the present analysis, we will present in section 2 the new spectroscopic data collected for several among the lowest mass candidate members of the association (two M8 and a L5), together with an estimate of the age and masses. Section 3 will  revisit the data regarding the known members of the association, to estimate an accurate age by using different methods. We will end the paper with the conclusions and a discussion about the limits of the age-estimate techniques. ",
        "conclusions": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  The TW Hydrae association is becoming a cornerstone  in our understanding of phenomena such as stellar  accretion and the formation of planetary systems, due to the fact that it contains several dozen candidate members (only a fraction of them confirmed), is nearby, and its age is critical for the evolution of several stellar properties. Our main goal has been to establish an accurate age for the whole association and, in the process, for the brown dwarf 2M1207, an accreting member which contains, at least, a planetary mass object located at a projected distance of  about $\\sim$50 AU.  We have used the available data for the {\\it bona-fide} members and compared  them with the properties of well-known star forming regions and young clusters, as  well as theoretical models. The methods we have used include HR and Color-Magnitude Diagrams, surface gravity,  activity, and lithium. Table 2 summarizes all the age estimates, including minimum and maximum values. Depending on the method (the property we have used to estimate the age), the derived values are in the range of 5 Myr (Color-Magnitude Diagram with 2MASS data and theoretical models, not very reliable) to 15 Myr (for the gravity). The ''best'' value is 10 Myr. The minimum age is imposed by Color-Magnitude and HR Diagrams, as well as by the lithium abundance, whereas  an upper limit of 20 Myr can be derived from the HR Diagram. Therefore, we can say that the age of the association is 10$^{+10}_{-7}$ Myr.  In the case of 2M1207, the gravity, as measured by mid-resolution optical spectroscopy, and its comparison with theoretical models by Baraffe et al. (1998) indicates that the age is 15$^{+15}_{-10}$ My, in agreement with the age of the association. As a comparison, Mamajek (2005)  derived an age for TWA  based on the moving cluster distance method, an expansion age. After confirming the membership of 2M1207, he  obtained a lower limit of 10 Myr. This value is in contradiction with an estimation  based on the dynamical evolution of the association (Makarov \\& Fabricius 2001; de la Reza, Jilinski, \\& Ortega 2006), 8.3$\\pm0.8$ Myr.  Moreover, Makarov et al. (2005) have found that three  members (TWA 1, 4 and 11) might be younger, since they  derived an expansion age of 4.7$\\pm$0.6 Myr for them. However, our own estimation is compatible with these  determinations.  In this analysis, we have assumed the coevality of the TW Hydrae Association. This is an hypothesis which might not be true. Lawson \\& Crause (2005) have found a bimodal distribution in the rotation periods, and their interpretation is that there are two distinct populations, which are 10 and 17 Myr old. As we have shown, the almost identical 2M1207 and 2M1139 show small differences in their spectrum, which might be due to different gravities,  which would indicate different  ages. Therefore, non-coevality is a possibility.  We have several  caveats regarding these age estimations. First, age estimation is  difficult. Each individual method  has its own problems. Some among them depend on the direct comparison with theoretical models (such as the HR diagram), other include not well-understood properties (such as the process of accretion and its evolution evolution or the lithium abundance). There are also problems transferring observational into theoretical quantities, such as from magnitudes and colors into effective temperatures and luminosities (Stauffer, Hartmann \\& Barrado y Navascu\\'es 1995). Moreover, colors can be affected by activity and related effects (see the  case of the Pleiades,  Stauffer et al. 2003). On the other hand, different methods can provide very different ages. A case study is lithium depletion chronology for young clusters such as NGC2547, IC2391, Alpha Per and the Pleiades, which produces ages about 50\\% larger (Stauffer et al. 1998, 1999, Barrado y Navascu\\'es et al. 1998, 1999, 2004; Oliveira et al. 2003)  than canonical values, as derived from isochrone fitting in the lower Main  Sequence or the turn-over method.  Therefore, our age --or any age-- estimation should be  understood in the context   of a particular age scale, in  an evolutionary sequence. The important fact is that we can sort different star forming regions,  moving groups, associations and open clusters by their age, stating, for instance, that Taurus is younger than Lambda Orionis, which is younger than IC2391. We can be sure that TWA is in between these two latter clusters. Since we have analyzed several associations, trying to understand different, although related, phenomenology, in the same way (see, for instance, the case of protoplanetary disks of $\\beta$ Pic and HR4796 and Fomalhaut (Barrado y Navascu\\'es et al. 1999, 2001 and references therein), we  are trying to link a whole array of properties in a coherent, but relative, evolutionary context. However, much work remains to be done.  An absolute age-scale (or scales), based on homogeneous datasets (same methodology to acquire and analyze  the data), is needed, coupled with theoretical work, in order to reach an agreement between the estimates provided by different properties.            %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0608/astro-ph0608521_arXiv.txt": {
        "abstract": "Momentum feedback from super-Eddington accretion offers a simple explanation for the observed $M$--$\\sigma$ and $M$--$M_{\\rm spher}$ relations between supermassive black holes (SMBH) and the spheroids of their host galaxies. Recently \\citeauthor{FerrareseEtal06} and \\citeauthor{WehnerHarris06} observed analogous relations between the masses of central star clusters and their hosts. We show that stellar winds and supernovae from such nuclear clusters (NC) give similar feedback explanations for this case also, and we discuss the connection to the Faber-Jackson relation for the spheroids themselves. ",
        "introduction": "\\label{sec:intro}  It is well known that the masses of supermassive black holes (SMBHs) in the nuclei of early-type galaxies and late-type bulges correlate tightly with the velocity dispersions of the stellar spheroids: $M_{\\rm BH} \\propto \\sigma^{x}$, with $x \\simeq 4.0$--4.5 \\citep{TremaineEtal02,FerrareseFord05}. $M_{\\rm BH}$ also increases nearly linearly with galaxy spheroid mass (e.g., \\citealt{HaringRix04}). It is now becoming clear that many galaxies have nuclear star clusters (NCs) with masses similarly connected to the host properties.  A recent HST/ACS survey of 100 early-type galaxies in Virgo has found that $\\sim\\! 70\\%$--80\\% of systems with $-20.5\\la M_B\\la -15$ contain dense nuclear components that are resolved by HST into star clusters with luminosities $L\\simeq 10^5\\,L_\\odot$--$5\\times10^7\\,L_\\odot$, half-light radii of about 4~pc (possibly with a size-luminosity relation among the brighter nuclei), and colors suggesting ages of $\\sim\\! 2$--10~Gyr and metallicities ${\\rm [Fe/H]}\\approx -0.5 \\pm 1$ \\citep{CoteEtal06}. Most spiral bulges and bulgeless disk galaxies also contain nuclear clusters with rather similar properties \\citep{PhillipsEtal96, CarolloEtal98, BoekerEtal02, WalcherEtal05}.  All known NCs have masses less than a few $10^8\\,M_\\odot$, while most measured SMBHs have $M_{\\rm BH}\\ga 10^8\\,M_\\odot$. The apparent paucity of low-mass SMBHs is at least partly a selection effect, but the upper limit on the nuclear clusters may well be real. \\citet{CoteEtal06} find that NC luminosity increases with spheroid luminosity, such that $M_{\\rm NC}\\ga 2\\times10^8\\,M_\\odot$ is expected for galaxies with $M_B\\la -20.5$. But although they have HST surface photometry of the cores of {\\it all} such galaxies in Virgo, \\citeauthor{CoteEtal06} find no evidence for nucleation in any of them.  \\citet{FerrareseEtal06} have obtained long-slit spectra for 29 of the nucleated ellipticals in Virgo. They find that the NC masses correlate well with the galaxies' velocity dispersions averaged over an effective radius ($\\approx1$--2 kpc): $M_{\\rm NC}\\propto \\sigma^x$ with $x=4.3\\pm0.6$, essentially the same as for the SMBH relation. However, they also find an {\\it offset} between the cluster and black-hole scalings. Fitting power-law $M$--$\\sigma$ relations with $x\\equiv 4$ to the \\citeauthor{FerrareseEtal06} NC data and to SMBH data from the literature yields (L.~Ferrarese, private communication) \\medskip \\begin{equation} \\begin{array}{lcl} \\log\\,M_{{\\rm NC}, 8} & = & (1.25 \\pm 0.55) \\ + \\ 4\\,\\log\\,\\sigma_{200}  \\\\ \\log\\,M_{{\\rm BH}, 8} & = & (0.25 \\pm 0.33) \\ + \\ 4\\,\\log\\,\\sigma_{200}  \\ , \\end{array} \\label{eq:msigma4} \\end{equation}  \\medskip \\noindent where $M_8\\equiv M/10^8\\,M_\\odot$ and $\\sigma_{200}\\equiv \\sigma/(200\\ {\\rm km\\,s}^{-1})$. A limit of $M_{\\rm NC}\\la 2\\times10^8\\,M_\\odot$ thus corresponds to $\\sigma\\la 120$ km~s$^{-1}$. The nuclei of spheroids with velocity dispersion less than this are dominated by stellar clusters, with the mass of any SMBH that might be present expected to be $\\simeq\\! 10$ times smaller. More massive galaxies apparently always contain nuclear SMBHs but, so far as is known, not NCs.  \\citet{WehnerHarris06} use photometric data in the literature for about 40 dwarf-elliptical nuclei to show that $M_{\\rm NC}$ increases almost linearly with galaxy spheroid mass. \\citet{FerrareseEtal06} also find this for their nucleated Virgo galaxies with dynamical mass estimates. Moreover, both studies conclude that the $M_{\\rm NC}$--$M_{\\rm spher}$ and $M_{\\rm BH}$--$M_{\\rm spher}$ relations meet almost seamlessly at a mass scale $\\sim\\! 10^8\\,M_\\odot$, i.e., there is no large offset as in the $M$--$\\sigma$ relations. These authors therefore refer to nuclear clusters and supermassive black holes together as ``central massive objects,'' or CMOs. We adopt this term here.  A derivation of the $M_{\\rm BH}$--$\\sigma$ scaling has been given by \\citet{King03, King05} (see also \\citealt{Fabian99}; \\citealt{MurrayEtal05}; \\citealt{BegelmanNath05}). He considers super-Eddington accretion on to a seed SMBH at the center of an isothermal dark-matter halo. Accretion feedback produces a momentum-driven superbubble that sweeps ambient gas into a thin shell, which expands into the galaxy. Eventually the shock cooling time becomes so long that the shell becomes energy-driven and accelerates to escape the galaxy. This truncates accretion and freezes in a relation of the form $M_{\\rm BH}\\propto \\sigma^4$ that, with no free parameters, matches the observed one remarkably well. A roughly linear relation between $M_{\\rm BH}$ and $M_{\\rm spher}$ is also established as part of this process.  In this {\\it Letter}, we examine the possibility that the mass of a central star cluster in a protogalaxy might be similarly self-regulated, by feedback from stellar winds and supernovae.\\footnotemark \\footnotetext{After this paper was submitted, a preprint appeared by \\citet{LiEtal06}, which uses numerical simulations to study the development of an apparently common $M_{\\rm CMO}$--$M_{\\rm spher}$ relation for both NCs and SMBHs. Unlike that work, in our analytical model we consider stellar feedback explicitly to explain this result, and moreover we attempt to understand the {\\it offset} between the $M$--$\\sigma$ scalings for NCs vs.~SMBHs.} ",
        "conclusions": "\\label{sec:discussion}  Current data probably do not rule out the idea that nuclei of galaxies with $\\sigma\\la 120$ km~s$^{-1}$ could harbor {\\it both} star clusters {\\it and} SMBHs some $\\sim\\!10$ times less massive, and thus it will be important to ask how such galaxies might choose between SMBH and NC feedback channels in regulating their formation. For now, a possibly more straightforward question is why there are {\\it no} nuclear clusters in larger galaxies with $\\sigma\\ga120$ km~s$^{-1}$ ($M_{\\rm spher}\\ga 2$--$3\\times10^{10}\\,M_\\odot$), which apparently all contain black holes.  For our basic scenario to be self-consistent, we evidently require that the dynamical time of the superbubble always be shorter than that of the entire halo, until the point of blow-out when $t_{\\rm flow}=t_{\\rm cool}$. Equivalently, the radius of the shell must always be less than the halo virial radius, and in particular $R_{\\rm cool}\\la R_{\\rm vir}$ is required for the growth of a CMO (whether an SMBH or an NC) to be self-regulated as we envision. From the relations in \\citet{BryanNorman98}, if $\\Omega_{m,0}=0.3$, $\\Omega_\\Lambda=0.7$, and $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$, then $R_{\\rm vir}\\approx540\\ {\\rm kpc}\\ \\sigma_{\\rm 200}\\,(1+z)^{-1.1}$ (accurate to better than 10\\% for $z\\le 2$). Combining this with equation (\\ref{eq:Ratom}), our picture can work for NCs forming at redshift $z_{\\rm NC}$ only in halos with $\\sigma \\la 160\\ {\\rm km~s}^{-1}\\,(1+z_{\\rm NC})^{-0.55}$.  If nuclear clusters are typically  $\\sim\\! 5$ Gyr old, $z_{\\rm NC}\\simeq 0.5$ and this upper limit becomes $\\sigma\\la 130$~km~s$^{-1}$. In halos with velocity dispersion higher than this, the superbubble blown by an NC reaches the halo virial radius before it becomes energy-conserving and can accelerate to escape. It is presumably then held there, or even driven into collapse, by the infall of material from beyond $R_{\\rm vir}$. The growth of the central star ``cluster'' is never choked off, and it ultimately becomes indistinguishable from the galaxy spheroid itself. When the CMO is a black hole, the much longer Compton-cooling time found in \\citet{King03} implies that $R_{\\rm cool}\\la R_{\\rm vir}$ for all $\\sigma\\la 500\\,{\\rm km~s}^{-1}\\,(1+z_{\\rm BH})^{-0.55}$, so it is only in very massive systems indeed that the self-regulated $M_{\\rm BH}$--$\\sigma$ relation of equation (\\ref{eq:Mcmo1}) breaks down."
    },
    "0608/astro-ph0608001_arXiv.txt": {
        "abstract": "{A summary is given of planetary nebulae abundances from ISO measurements. It is shown that these nebulae show abundance gradients (with galactocentric distance), which in the case of neon, argon, sulfur and oxygen (with four exceptions) are the same as HII regions and early type star abundance gradients. The abundance of these elements predicted from these gradients at the distance of the Sun from the center are exactly the solar abundance. Sulfur is the exception to this; the reason for this is discussed. The higher solar neon abundance is confirmed; this is discussed in terms of the results of helioseismology. Evidence is presented for oxygen destruction via ON cycling having occurred in the progenitors of four planetary nebulae with bilobal structure. These progenitor stars had a high mass, probably greater than 5M\\smallsun. This is deduced from the high values of He/H and N/H found in these nebulae. Formation of nitrogen, helium and carbon are discussed. The high mass progenitors which showed oxygen destruction are shown to have probably destroyed carbon as well. This is probably the result of hot bottom burning. ",
        "introduction": "Planetary nebulae (hereafter PNe) are an advanced stage of stellar evolution of low and intermediate mass stars. Their abundances have been changed by various processes which have occurred in these objects since their formation. Not all elements have been affected in the same way or to the same extent; some elements have not been affected in the course of evolution. Our present purpose is to determine which elements have been processed in the course of evolution, and which are the affected nebulae. This is done using the abundances determined with the help of the Infrared Space Observatory observations (ISO, Kessler et al.\\,1996). In particular the result of the Short and Long Wavelength Spectrometers have been used (hereafter SWS and LWS respectively).  In order to determine which abundances have changed in the course of evolution it is necessary to find the initial abundances before evolution began. These values are taken from three sources: HII regions, young (B and O type) stars, and the Sun. For the first two groups it is known that they show a gradient in abundance as a function of their distance from the Galactic Center (see for example Shaver et al.\\,\\cite{shaver} and Rolleston et al.\\,\\cite{rolle}). Thus in comparing planetary nebulae abundances with that in HII regions and young stars this gradient must be taken into account. This is done by plotting the abundances as function of the distance from the galactic center and then comparing these plots.  Two potential problems arise. Firstly, the distance to individual objects from the Sun is often uncertain. This is especially a problem for PNe whose distances are sometimes uncertain by more than 50\\%. But because the distance to the Galactic Center depends in an important way on the direction of the nebula from us, the distance to the Galactic Center is actually considerably less uncertain than the distance to the Sun.  The second problem is the following. It is assumed that the galactic abundance gradient arises because the material out of which the object is formed has a different (initial) abundance at different distances from the galactic center. The abundance in the objects measured reflects the initial abundance at the position in which they are measured. This is easy to defend in the case of HII regions and O and B stars because it is unlikely that these young objects have moved far (radially) from the position in which they have been formed.  This is not obvious for planetary nebulae since they are older objects; they may have orbited the galaxy one or several times since they were formed. Only if they are on nearly circular orbits will their present position reflect the position of formation. We will assume that this is true for the nebulae to be considered, as is usually done (e.g. Maciel \\& Costa 2003). We demonstrate in this paper that this leads to consistent results.  The subject of nebular abundances has been discussed often in the past 25 years, especially from the theoretical point of view. A recent summary has been given by Lattanzio \\cite{latt} at the 2001 IAU PNe Symposium. In ending this summary Lattanzio remarks the following: 'What is really needed is a complete test of the model predictions against observations, which has so far been rather limited...'. This limitation lies in a general feeling that the theory qualitatively explains the observations, while quantitatively little has been done. The reason for this is, on the one hand, the uncertainty of the theoretical results, and on the other hand the inaccuracy of the observational abundances.    That the abundances derived from observations are often not very accurate can be seen by a comparison of abundances derived for well studied nebulae. Abundances derived by different authors for the same nebula often show disagreements of a factor of 2; sometimes considerably larger differences are found. These differences are usually not due to differences in the observations used, and they occur in spite of the fact that often the different authors used the same or similar observations. The differences occur for two reasons. The first is that the electron temperature ($T_e$) is usually assumed to be constant in the nebula. This assumption is made because using the optical spectrum only two ions have a sufficient number of lines to determine $T_e$: these are \\ion{[O}{iii]} and \\ion{[N}{ii]}. These two ions usually give somewhat different temperatures. The reason that one still chooses to use a single temperature is probably that one is not certain enough of the reality of this difference. Only after the ISO infrared spectra became available could these temperature variations be verified, because enough lines are now measured in many nebulae so that electron temperatures from five to ten ions may be obtained for a given nebula. In addition, these temperatures are correlated with the ionization potential required to reach the observed ionization stage. Thus a temperature gradient is often found; the higher temperature regions occuring for the ions with the higher ionization potentials; these probably lie closer to the center of the nebula. From this it is now possible to measure (or predict) more accurately the electron temperature to be assigned to each measured ion. This in turn leads to an important improvement in the abundance of the ion considered. Actually when the infrared lines are used, the derived abundances are relatively insensitive to the electron temperature since the lines are formed from very low lying energy levels.  There is a second reason why the inclusion of the infrared observations leads to increased accuracy. This is because many more ions are now measured, and therefore there are many less unobserved stages of ionization. This improvement in abundance accuracy is especially important for neon, argon, nitrogen and sulfur.  One of the difficulties which are encountered in the analysis of the SWS data (see de Graauw et al.\\,1996) occurs when the nebula is larger than the aperature of the instrument and a correction must be made. This is usually done by relating the hydrogen and helium lines in the SWS spectrum to lines of the same elements in other spectral regions. The theoretical line ratios then relate the lines in the various spectral regions. The errors introduced in this correction are rather limited. The LWS ISO observations (see Clegg et al.\\,1996) are made with a larger diaphragm so that no diaphragm correction is necessary. A more important difficulty involve the density dependence of a few of the infrared lines with very low transition probabilities. In practice this involves the \\ion{[N}{ii]} line at 121$\\mu$m, the \\ion{[O}{iii]} lines at 51 and 88$\\mu$m and to a lesser extent the \\ion{[N}{iii]} line at 57$\\mu$m and the \\ion{[S}{iii]} line at 33$\\mu$m. Errors can occur in low density nebulae if the density is not well determined or if large density gradients are present. For three of these ions other lines are available to check this effect.  In this paper the abundances of 26 nebulae are considered. They are listed in the Appendix, together with the references from which they were taken. Their abundances have all been determined including ISO SWS and LWS spectra and we believe are the most accurate available. They were all determined in a similar manner, and all contain information about electron temperature gradients in the nebulae. These nebulae have been selected primarily because they are strong infrared emitters. They are therefore mostly nearby PNe. They are not very much larger than the ISO diaphragm of about 20\\arcsec. Because the constraints are mostly observational, this can be seen in a first approximation as a random selection of nearby nebulae. For 23 of these nebulae ultraviolet observations (IUE) are also available, without which carbon abundances cannot be determined.  This is the second time that abundance results including ISO data have been used in discussing element evolution in planetary nebulae. It was done several years ago by Marigo et al. \\cite{marigo} who used the results of 10 PNe. In that discussion new evolution models were made in discussing the results. We now have results from many more ISO PNe. In addition we approach the problem of the initial abundances of the stars which form PN in a different way. On the other hand we do not include any self-made evolutionary models.  This paper is structured as follows. First the references are given for the comparison objects (HII regions, early type stars and the Sun). These are presented in Sect.\\,2. In the next four sections oxygen, neon, sulfur and argon are discussed. These are the elements for which evolutionary abundance changes are relatively small. In Sect.\\,7 nitrogen and helium are discussed and in Sect.\\,8 carbon is discussed. In each of the individual sections a comparison is made with the abundances found in the theoretical evolutionary models made earlier by Marigo et al. \\cite{marigo}, and especially by Karakas \\cite{karakas}, which were chosen because of the large amount of detail available in these calculations. Sect.\\,9 gives a summary and conclusions. ",
        "conclusions": "\\indent 1) For neon, sulfur, argon and most of the oxygen abundances, the galactocentric PNe and HII values of composition are essentially identical. This is remarkable because the HII regions are young objects while PN are older objects. The conclusion is drawn that the galactocentric position of the PNe does not significantly change with time, at least for the nebulae which have been measured. This must mean that their galactic orbits are roughly circular.  In addition all galactocentric gradients are consistent with a single value: -0.085 dex/kpc assuming that the Sun is at 8 kpc from the Galactic Center.  2) The solar abundances of the above elements are the same as both the PNe and HII regions at 8 kpc from the center. This is shown quite strikingly for oxygen in Fig.\\,1. Sulfur is an exception, being almost a factor of two higher. Because it is difficult to find a reason for this difference, the question arises as to the correctness of either the solar sulfur abundance on the one hand or the HII and PNe abundance on the other.  3) The solar abundance of neon is the same in the PNe and HII regions, using the determination of the solar neon abundance given in Table 1. This determination is very similar to that of Feldman \\& Widing\\,\\cite{fw}, but differs from the solar neon abundance given by Asplund et al.\\, \\cite{asplund}. This has significance for the controversy concerning the consistency of solar models with helioseismological measurements (see discussions given by Antia \\& Basu \\cite{antia} and Bahcall et al.\\,\\cite{bahcall}). Our results support the higher neon abundance suggested by these authors.  4) Four of the 26 PNe studied show evidence of oxygen burning in the course of evolution. Three of these PNe have large bilobal structures, but there are other bipolar nebulae which show no evidence for oxygen burning. We have considered the suggestion of Marigo et al.\\,\\cite{marigo} that these PNe originated from material with a lower oxygen abundance but find the Karakas models, which are able to reproduce the observed N/O ratio starting with solar metalicity, a more likely expanation for the lower oxygen abundance.  5) The four PNe which show evidence of oxygen destruction all have high helium abundances. Their nitrogen abundances are somewhat higher than that of the oxygen. These four nebulae have abundances which are compatible with model predictions of Karakas\\,\\cite{karakas} for progenitor stars with masses higher than 5M\\smallsun, and a value of Z=0.02.  6) The remaining PNe show helium abundances ranging from slightly below solar, He/H=0.088 to He/H=0.15. The nitrogen increases from about solar to a factor of 10 above solar, directly proportional to the helium.  7) The carbon abundance is the most difficult to interpret because it shows the largest scatter. This is probably caused by the fact that in some nebulae carbon is created in the course of evolution while in others it is destroyed.  Carbon was only measured in three of the four high mass progenitor PNe, and in all these cases the carbon was found depleted. Of the remaining PNe about half have C/O$>$1 which is almost certainly an indication that carbon has been formed in the progenitor star in the course of evolution."
    },
    "0608/astro-ph0608237_arXiv.txt": {
        "abstract": "Neutral fluorine (F\\,{\\sc i}) lines are identified in the optical spectra of cool EHe stars. These are the first identification of F\\,{\\sc i} lines in a star's spectrum, and provide the first measurement of fluorine abundances in EHe stars. The results show that fluorine is overabundant in EHe stars. The overabundance of fluorine provides evidence for the synthesis of fluorine in these stars, that is discussed in the light of asymptotic giant branch (AGB) evolution, and the expectation from accretion of an He white dwarf by a C-O white dwarf. ",
        "introduction": "The solar system, for a long time, has been the only source of information about the fluorine abundance in the Galaxy \\citep{hall69}. The astrophysical origin of solar system's fluorine is not yet identified from the known theories of stellar nucleosynthesis. The major problem with fluorine production is that, the element has only one stable, yet rather fragile, isotope, $^{19}$F. In stellar interiors it is readily destroyed by hydrogen via $^{19}$F(p,$\\alpha$)$^{16}$O and helium via $^{19}$F($\\alpha$,p)$^{22}$Ne. The high F abundances which were measured, using infrared HF vibration-rotation transitions, in the asymptotic giant branch (AGB) stars from the \\citet{jori92} sample, provided clues for fluorine production in AGB stars. The \\citet{cunha2003} study measured F abundances in different populations: in galactic red giants (a subset of Jorissen et al. sample), in red giants in the Large Magellanic Cloud (LMC), as well as in a few stars in the globular cluster Omega Centauri. Concerning $\\omega$\\,Cen, the F abundances found in that study were quite low. One possible interpretation was that the AGB stars in $\\omega$\\,Cen did not contribute significantly to the fluorine. However, this conclusion assumes that there has been effective AGB pollution in $\\omega$\\,Cen.  Apart from the third dredge-up of AGB stars, the other two sources of fluorine are Type\\,II supernovae (SNe) explosions, and stellar winds from Wolf-Rayet (W-R) stars. \\citet{woos95} showed that in Type\\,II SNe, $^{19}$F can be created by spallation of $^{20}$Ne by neutrinos. In massive or W-R stars, $^{19}$F is probably produced during He-burning phase and, before it gets annihilated, is ejected into space by strong stellar winds \\citep{may2000}. In AGB stars $^{19}$F is predicted to be produced in the convective He-rich intershell and then dredged-up to the surface during the He-burning thermal pulses \\citep{forest92}. The reaction chain for F production in the He-burning environments of AGB and W-R stars is: $^{14}$N($\\alpha$,$\\gamma$)$^{18}$F($\\beta^+$)$^{18}$O(p,$\\alpha$)$^{15}$N($\\alpha$,$\\gamma$)$^{19}$F. Protons are provided by the $^{14}$N(n,p)$^{14}$C reaction with neutrons liberated from $^{13}$C($\\alpha$,n)$^{16}$O. The $^{14}$N and $^{13}$C nuclei are the result of H-burning by CNO cycling, and the initial $^{13}$C stock acts as a limiting factor for the $^{19}$F yield.  The Cunha et al.'s same fluorine abundance results for the Milky Way were later discussed in \\citet{renda2004} in light of a chemical evolution model for the Galaxy that took into account all possible sources for fluorine production. This chemical evolution model, however, cannot be applied to the peculiar globular cluster $\\omega$\\,Cen and the low fluorine abundance results for $\\omega$\\,Cen remain puzzling. Renda et al. suggest that both W-R and AGB stars are significant sources of fluorine, and that is supported by the evidence of enhanced F abundance from F\\,{\\sc v} and F\\,{\\sc vi} absorption lines in the far-UV spectra of hot post-AGB PG\\,1159 stars \\citep{werner2005}. Also, \\citet{zhang2005} determine F abundances from the nebular emission lines: [F\\,{\\sc ii}] line at 4789\\AA\\ and [F\\,{\\sc iv}] line at 4060\\AA, for a sample of Planetary nebulae (PNe). Fluorine is abundant in PNe, and this provides evidence for the synthesis of fluorine in the AGB phase. No evidence of enhanced fluorine resulting from spallation in Type\\,II SNe is suggested by the observations of F\\,{\\sc i} interstellar absorption line at 955\\AA\\ in two sight lines towards the Cep OB2 association using $Far$ $Ultraviolet$ $Spectroscopic$ $Explorer$ \\citep{feder2005}.  Extreme helium (EHe) stars, see for \\citet{pan2006} and the references therein, are suggested to have gone through AGB phase in their earlier evolution, hence, fluorine should be present in their atmospheres. Presence of fluorine in EHe's atmosphere can serve as a test bed for fluorine production in AGB stars. In this letter, we present fluorine abundances for the sample of cool EHe stars from F\\,{\\sc i} lines. ",
        "conclusions": "A good number of F\\,{\\sc i} lines were identified for the first time in the optical spectra of the sample of cool EHe stars. An abundance analyses of these spectra provide the first abundances of fluorine for cool EHe stars. The enhanced F abundances (see Figure 4) in cool EHes are evidence of nucleosynthesis in their earlier evolution. The abundances for the other elements presented in Figure 4 are adopted from \\citet{pan2001} for FQ\\,Aqr, LS\\,IV-14$^{\\circ}$\\,109, BD\\,-1$^{\\circ}$\\,3438, and LS\\,IV-1$^{\\circ}$\\,2, from \\citet{panred2006} for LSS\\,3378, and from \\citet{wall67} and \\citet{pana2006} for KS\\,Per. The F, C, and O abundances in cool EHes are independent of the star's initial metallicity (Fe abundance); the N abundances, however, show a clear trend with Fe (see Figure 4). These trends clearly suggest that F, C, and O abundances in cool EHes are associated with He-burnt material, and the N abundances are the result of H-burning via CNO cycling; the almost complete conversion of the initial C, N, and O to N is shown in the plot [N] vs [Fe] of Figure 4. During the He-burning thermal pulses $^{19}$F is produced, and the initial $^{13}$C supply then acts as a limiting factor for the $^{19}$F yield. To account for the enhanced F abundances, in cool EHes, with no trend with Fe, requires a primary $^{13}$C supply instead of the secondary $^{13}$C supply from the CNO cycle ashes. Note the [F] vs [H] trend of Figure 4.  \\begin{figure} \\epsscale{1.00} \\plotone{f4.eps} \\caption{[N], [C], [F], [O], and [Zr] vs [Fe], and also [F] vs [H]. The sample of cool EHes is represented by filled squares; the abundances are from the references given in Section 4. KS\\,Per is represented by open triangles; the abundances are from \\citet{wall67}, and \\citet{pana2006} in preparation. $\\sun$ represents Sun; the adopted solar abundances are from Table 1 of \\citet{lod2003}. The dotted line represents conversion of the initial sum of C and N to N. The dashed line represents the locus of the sum of initail C, N, and O converted to N. [X] = [Fe] are denoted by the solid lines where X represents N, C, F, O, and Zr. \\label{fig3}} \\end{figure}  The F abundances in H-deficient PG\\,1159 stars show a strong variation from star to star, ranging from solar to 250 times solar \\citep{werner2005}. The sample cool EHes show F enhancements about 100 times solar. Note that the models of \\citet{lug2004}, for thermally pulsating AGB stars, predict that the F enrichment in the intershell is a strong function of the initial stellar mass; for solar metallicity models, log $\\epsilon(\\rm F)$, which is about 100 times solar, is maximal at M = 3$M_\\odot$. Lugaro et al. conclude that their F intershell abundances are not high enough to explain the enhanced F abundances which were measured in the AGB stars from the Jorissen et al. sample. The surface composition of these H-deficient stars: PG\\,1159 and EHe, could be a result of final He-shell flash in a single post-AGB star (FF scenario), or a merger of two white dwarfs (DD scenario). Although the FF scenario accounts for PG\\,1159 stars, present theoretical calculations \\citep{herw99} imply higher C/He and O/He ratios than are observed in EHe stars. However, DD scenario can account for the observed surface abundances of most of the elements in EHes \\citep{pan2006}.  Enrichment of $s$-process elements, and also fluorine, is not expected for the EHes resulting from a merger of a He with a C-O white dwarf (DD scenario) as discussed by \\citet{pan2004,pan2006}. However, synthesis of fluorine and $s$-process elements may occur during the merger and needs to be explored. The $s$-process elements are enriched in the cool EHe LSS\\,3378 (see for example, [Zr] vs [Fe] of Figure 4). The neutron source in EHes showing enriched $s$-process elements is not known. $^{13}$C($\\alpha$,n)$^{16}$O and $^{22}$Ne($\\alpha$,n)$^{25}$Mg are the two possible sources. If the latter reaction dominates, then F is unlikely to survive in the He-rich intershell, as the reaction rate of $^{19}$F($\\alpha$,n)$^{22}$Ne is much higher than that of $^{22}$Ne($\\alpha$,n)$^{25}$Mg \\citep{jori92}. Since $^{19}$F is overabundant in cool EHes, [an upper limit to the F abundance for the coolest hot EHe PV\\,Tel is also derived (see Table 2)], the reaction $^{22}$Ne($\\alpha$,n)$^{25}$Mg is unlikely to be the neutron source in EHe stars. Non-LTE calculations should be performed for the key elements to further improve the chemical analysis."
    },
    "0608/astro-ph0608147_arXiv.txt": {
        "abstract": "We observed 5 Hickson Compact Groups with the ESO/MPI 2.2m telescope and WFI to investigate the dwarf galaxy content and distribution in these galaxy groups.  Our deep imaging and careful selection of the candidate galaxies revealed a rich population of mainly passively evolving dwarf galaxies, which is spatially much more extended than the originally defined Hickson Compact groups. The composite luminosity function of the 5 groups shows a bimodal structure with a very steep rise in the low luminosity regime. The faint end slope is close to the predictions of CDM theory for the slope of the Dark Matter halo mass function. ",
        "introduction": "\\label{sec:1} Hickson Compact Groups (HGCs) are tight groups of galaxies selected by Hickson on POSS-I plates based galaxy density and isolation criteria as described e.g. in \\cite{Hickson1997}. Analysing the environments of HGCs it soon becomes clear, that the cores Hickson classified are often not isolated, but embedded into larger structures. As an example we found in HGC 16 an spiral galaxy with the same redshift, which is located just outside the radius defined in Hickson's isolation criterium. %  Searches for dwarf galaxies in HGCs which were conducted up to now were either deep, but limited to the area of the compact group itself (e.g. \\cite{Hunsberger1996}) or wide field, but rather shallow (e.g. \\cite{Ribeiro1994, Zabludoff2000}. We used this as a starting point to investigate the dwarf galaxy content of HGCs and to map out the structure of HGCs and their envelopes. ",
        "conclusions": ""
    },
    "0608/gr-qc0608100_arXiv.txt": {
        "abstract": "An exactly solvable bounce model in loop quantum cosmology is identified which serves as a perturbative basis for realistic bounce scenarios. Its bouncing solutions are derived analytically, demonstrating why recent numerical simulations robustly led to smooth bounces under the assumption of semiclassicality. Several effects, easily included in a perturbative analysis, can however change this smoothness. The effective theory is not only applicable to such situations where numerical techniques become highly involved but also allows one to discuss conceptual issues. For instance, consequences of the notoriously difficult physical inner product can be implemented at the effective level. This indicates that even physical predictions from full quantum gravity can be obtained from perturbative effective equations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608692_arXiv.txt": {
        "abstract": "MACHO-96-BLG-5 is a microlensing event observed towards the bulge of the Galaxy with exceptionally long duration of $\\sim 970$ days. The microlensing parallax fit parameters were used to estimate the lens mass $M=6^{+10}_{-3}$ M$_{\\odot}$ and its distance $d$ which results to be in the range $0.5$ kpc - $2$ kpc. The upper limit on the absolute brightness for main-sequence stars of the same mass is less than $1$ $L_{\\odot}$ so that the lens is a good black hole candidate. If it is so, the black hole would accrete by interstellar medium thereby emitting in the $X$-ray band. Here, the analysis of an {\\it XMM} deep observation towards MACHO-96-BLG-5 lens position is reported. Only an upper limit (at $99.8\\%$ confidence level) to the $X$-ray flux from the lens position of $9.10\\times 10^{-15}$-- $1.45\\times 10^{-14}$ erg cm$^{-2}$ s$^{-1}$ in the energy band $0.2-10$ keV has been obtained from that deep observation allowing to constrain the putative black hole accretion parameters. ",
        "introduction": "Gravitational microlensing is nowadays a well established technique to map both visible and dark matter throughout the Galaxy. The first lines of sight to be explored have been those towards the Magellanic Clouds and the bulge of the Galaxy, leading to the observation of several hundreds of microlensing events (\\citealt{alcock} and \\citealt{sumi}). The interpretation of the observations, although debated and controversial, supports the existence in the galactic halo of MACHOs (Massive Astrophysical Compact Halo Objects) with mass of $\\simeq 0.4~M_{\\odot}$. Microlensing events towards the bulge of the Galaxy are interpreted as due to lenses prevalently belonging to the known stellar populations. However, several long-duration events towards the bulge have significant probabilities of being due to Black Hole (BH) lenses (\\citealt{bennett2002} and \\citealt{mao2002}).  It is well known that the light curve of a typical microlensing event depends on four parameters (e.g. the mass $M$ of the lens, the observer-lens distance $d$, the observer-source distance $s$ and the lens transverse velocity $v_{\\bot}$), while only two quantities are available from observations (the amplification at maximum and the event duration). Therefore, it remains a degeneracy in microlensing parameters, so that the lens mass can not be univocally determined, but only a statistical information can be obtained from the analysis of a large enough number of microlensing events.  The parameter degeneracy may be solved in a few parallax events, which are events with duration long enough to make possible to estimate $d$ and/or $v_{\\bot}$. In this case, a fit procedure to the event light curve allows to extract the value of the lens mass. This is the case of at least 22 microlensing parallax events \\citep{poinsexter2005} observed by the MACHO collaboration towards the bulge of the Galaxy. Among these events three are particularly interesting: MACHO-96-BLG-5, MACHO-98-BLG-6 and MACHO-99-BLG-22. For these three events, the lens mass has been estimated to be $6_{-3}^{+10}$ M$_{\\odot}$, $6_{-3}^{+7}$ M$_{\\odot}$ \\citep{bennett2002}, and  $11_{-6}^{+12}$ M$_{\\odot}$ \\citep{mao2002,agol2002}, respectively.  The observed upper limit on the absolute brightness of these lenses is $<1$ $L_{\\odot}$ so the lenses can not be stars but are most likely BHs\\footnote{ Poindexter et al. (\\citeyear{poinsexter2005}) also concluded that MACHO-99-BLG-22 is a strong BH candidate ($78\\%$), MACHO-96-BLG-5 is a marginal BH candidate ($37\\%$) and MACHO-98-BLG-6 is a weak BH candidate ($2.2\\%$).} \\citep[see][]{bennett2002}. Stellar mass BHs there exist in the Galaxy as a consequence of the evolution of massive stars, but all the candidates known up to date are members of binary systems. On the other hand, because of the shape of the microlensing light curves, we expect that the lens is an isolated object, so that the coordinates of these events give us the direction toward which pointing an $X$-ray instrument and acquire information on the lens nature.  Indeed, the putative BH lens may be accreting interstellar gas and could be luminous in the X-ray band if it is within the thin gas layer of the galactic disk \\citep{mao2002}. A first estimate of the X-ray flux in the 1-10 keV band leads to fluxes of $\\simeq 10^{-15}$ erg cm$^{-2}$ s$^{-1}$ \\citep{agolkam}.  In this paper we report on a $100$ ks {\\it XMM} observation towards the MACHO-96-BLG-5 position, in order to search for an $X$-ray signature from the accretion process of the BH candidate.  The paper is structured as follows: In Sect. 2, we give a short description of the gravitational lensing event MACHO-96-BLG-5 and in Sect. 4 we discuss about the $X$-ray observations performed towards the target by other $X$-ray telescopes. In Sect. 4 we report our observational results. Finally, in Sect. 5 we address some conclusion.  Before closing this section we would like to mention that detecting isolated BHs is of great importance in astrophysics since this should allow to validate or not the standard model for stellar evolution and BH formation\\footnote{Indeed, it could be that nature has more difficulties in producing BHs than expected within the standard stellar evolution model as recently emerged by CHANDRA observation of the $X$-ray pulsar CXO J164710.2-455216, indicating that pulsar progenitor had mass greater than about $40M_{\\odot}$ \\citep{muno}.}. Simple estimates indicate in fact that about $10^8$ BHs should be present throughout the Galaxy (se e.g. \\citealt{st}) and this is in agrement with the chemical enrichment by supernovae indicating that about $2\\times 10^8$ BHs should have formed \\citep{samland}. The MACHO and OGLE groups claimed that three isolated BHs have been detected by gravitational microlensing indicating that at most $\\simeq 5\\times 10^8~(9M_{\\odot}/M_{BH})$ BHs reside in the Milky Way disk (\\citealt{agol2002}, \\citealt{alcock2000}). However, the number of isolated BHs present in the Galaxy should be confirmed by direct observations of their $X$-ray or radio emission \\citep{maccarone}. ",
        "conclusions": "MACHO-96-BLG-5 is a microlensing observed event towards the bulge of the Galaxy with the exceptionally long duration of $\\sim 970$ days\\footnote{We also mention that recently Poindexter et al. (\\citeyear{poinsexter2005}) presented the best fit solution for geocentric timescale for the same microlensing event and found that the event duration is $546 \\pm 165$ days or $698 \\pm 303$ days. This analysis leads to the conclusion that the lens of MACHO-96-BLG-5 is BH with a probability of about $37\\%$.}. The measure of the parallax angle \\citep{bennett2002} made it possible to estimate both the lens mass ($M=6^{+10}_{-3}$ M$_{\\odot}$) and its distance $d$ from Earth ($0.5$--$2$ kpc).  Since the observed upper limits on the absolute brightness of main-sequence stars for this lens is  $<1~L_{\\odot}$, MACHO-96-BLG-5 is a BH candidate. If it is so, then the BH would accrete by interstellar medium (via the Bondi-Hoyle mechanism) thereby emitting $X$ rays. Here, we have reported about a deep {\\it XMM} observation in the $0.2-10$ keV band towards MACHO-96-BLG-5 lens position. The analysis of the observations (see Section 4) did not show any obvious source  and it is consistent with a non-detection of the source. However, it is possible to estimate an upper limit on the source luminosity, which turns out to be in the range $9.10\\times 10^{-15}$-- $1.45\\times 10^{-14}$ erg cm$^{-2}$ s$^{-1}$, depending on the power law index $\\Gamma$.  This upper limit allows to constrain the accretion parameters of the BH candidate from a parametric study involving eqs. (\\ref{flusso1}) and (\\ref{flusso2}). In the BH mass-distance plane given in Figure \\ref{par}, the curved line corresponds to the mass-distance relation (also given in Figure \\ref{figure1}). This means that our BH candidate should lie on this curve (if its transverse velocity is $\\simeq 30$ km s$^{-1}$) and within the two dotted vertical lines. The three oblique lines correspond to the $X$-ray upper limit of $9.10\\times 10^{-15}$ erg cm$^{-2}$ s$^{-1}$ for three different values of the scale flux ($F_s$, $10^{-1}F_s$ and $10^{-2}F_s$). Obviously, varying $F_s$, it means that the accretion parameters ($\\epsilon$, $\\rho_{\\infty}$ and $v$) are changing. Thus, the region below the line labeled $F_s$ corresponds to $F_X<9.10\\times 10^{-15}$ erg cm$^{-2}$ s$^{-1}$. The same holds for the regions below the other two lines. It is noticing that a BH accreting with $\\rho_{\\infty}\\simeq 8.35\\times10^{-25}$ g cm$^{-3}$, $\\epsilon =0.1$ and $v\\simeq 30$ km s$^{-1}$ (implying $F_s\\simeq 4.45\\times 10^{-15}$  erg cm$^{-2}$ s$^{-1}$), is inconsistent with the best fit values (mass and distance) for the MACHO-96-BLG-5 lens candidate. However, for a less efficient Bondi accretion ($\\epsilon \\simeq 10^{-5}-10^{-3}$) the regions below the other two oblique lines show that the agreement is recovered.  Finally, it could be interesting to estimate the BH radiative efficiency $\\eta=L_X/L_{Edd}$ with respect to the maximum allowed BH accretion rate given by the Eddington luminosity $L_{Edd}\\simeq 1.38\\times10^{38}(M/M_{\\odot})$ erg s$^{-1}$. By using the BH mass and distance ranges, namely $3<M<16$ M$_{\\odot}$ and $0.5<d<2$ kpc, we can compare the Eddington luminosity with the $X$-ray luminosity given by $L_X=F_X4\\pi d^2$ expected for MACHO-96-BLG-5. This procedure gives $\\eta$ in the range $1.2\\times 10^{-10}$ -- $2.7\\times 10^{-9}$, in agreement with the efficiency estimates for BH accretion in quiescent galaxies and ultra-low luminous AGNs for which $\\eta = 4\\times 10^{-12}$-$~6\\times 10^{-7}$ \\citep{baganoff}.   \\begin{figure} \\epsscale{.80} \\plotone{f4.eps} \\caption{In the mass-distance plane (see Figure \\ref{figure1}), the three oblique lines define the regions (below each line) where the $X$-ray flux due to the Bondi accretion is below the upper limit flux of $9.10\\times 10^{-15}$ erg cm$^{-2}$ s$^{-1}$.}\\label{par} \\end{figure}"
    },
    "0608/astro-ph0608371_arXiv.txt": {
        "abstract": "We calculate the thermal conductivity of electrons produced by electron-electron Coulomb scattering in a strongly degenerate electron gas taking into account the Landau damping of transverse plasmons. The Landau damping strongly reduces this conductivity in the domain of ultrarelativistic electrons at temperatures below the electron plasma temperature. In the inner crust of a neutron star at temperatures $T \\lesssim 10^7$~K this thermal conductivity completely dominates over the electron conductivity due to electron-ion (electron-phonon) scattering and becomes competitive with the the electron conductivity due to scattering of electrons by impurity ions. ",
        "introduction": "\\label{introduc}  The electron thermal conductivity is an important kinetic property of plasmas in experimental devices, in metals and semiconductors and in astrophysical objects. It has been studied for a long time and described in textbooks (see, e.g., Ziman \\cite{ziman60}, Lifshitz and Pitaevski{\\u\\i} \\cite{lp81}). We will show that some aspects of this problem have to be revised.  Specifically, we will consider the thermal conductivity of degenerate electrons. It is important in metals, semiconductors as well as in degenerate cores of evolved stars (giants, supergiants and, most importantly, white dwarfs), and in envelopes of neutron stars. It is needed to study cooling of white dwarfs (e.g., Prada Moroni and Straniero \\cite{ps02}) and nuclear explosions of massive white dwarfs as type Ia supernovae (e.g.,  Baraffe, Heger and Woosley \\cite{bhw04}). In neutron star envelopes it is required for many reasons --- for simulating cooling of isolated neutron stars (see, e.g., Lattimer et al.\\ \\cite{lvpp94}, Gnedin, Yakovlev and Potekhin \\cite{gyp01}, Page, Geppert and Weber \\cite{pgw06}); for studying heat propagation and thermal relaxation in a neutron star envelope in response to the various processes of crustal energy release. In particular, these processes include pulsar glitches (e.g., Larson and Link \\cite{ll02}, and references therein), deep crustal heating of accreting neutron stars in soft X-ray transients (e.g., Ushomirsky and Rutledge \\cite{ur01}), superbursts as powerful nuclear explosions at the base of the outer crust of an accreting neutron star (e.g., Strohmayer and Bildsten \\cite{sb06}; Page and Cumming \\cite{pc05}). %  It is well known, that the electrons give the main contribution into the thermal conductivity of strongly degenerate matter. Their thermal conductivity can be written as \\begin{equation} \\kappa_e={\\pi^2 T k_{\\rm B}^2 n_e \\over 3m_e^*\\nu_e}, \\quad \\nu_e=\\nu_{ei}+\\nu_{ee}. \\label{kappa} \\end{equation} Here, $T$ is the temperature, $k_{\\rm B}$ is the Boltzmann constant, $n_e$ is the number density of electrons, $m_e^*=\\mu/c^2$, $\\mu$ is the electron chemical potential (including the rest-mass term), and $\\nu_e$ is the total effective electron collision frequency. The latter frequency is the sum of the partial collision frequencies $\\nu_{ei}$ and $\\nu_{ee}$. In $\\nu_{ei}$ we include all collisions of electrons mediated by their interactions with ions (direct Coulomb scattering of electrons by ions in an ion gas or liquid; electron-phonon scattering in an ion crystal and electron scattering by impurity ions at low temperatures).  Evidently, Eq.\\ (\\ref{kappa}) can be rewritten as \\begin{eqnarray} && { 1 \\over \\kappa_e} ={ 1 \\over \\kappa_{ei}} +{ 1 \\over \\kappa_{ee}}, \\quad \\kappa_{ei}={\\pi^2 T k_{\\rm B}^2 n_e \\over 3m_e^*\\nu_{ei}}, \\nonumber \\\\ &&  \\kappa_{ee}={\\pi^2 T k_{\\rm B}^2 n_e \\over 3m_e^*\\nu_{ee}}, \\label{kappae+i} \\end{eqnarray} where the partial conductivities $\\kappa_{ei}$ and $\\kappa_{ee}$ are determined, respectively, by the $ei$ and $ee$ collisions.  It is widely believed that the dominant contribution into $\\kappa_e$ comes from the $ei$ collisions. The associated partial conductivity $\\kappa_{ei}$ has been studied in a number of papers (e.g., Potekhin et al.\\ \\cite{pbhy99}, Gnedin et al.\\ \\cite{gyp01} and references therein). The partial thermal conductivity $\\kappa_{ee}$ owing to the $ee$ collisions was calculated by Lampe \\cite{lampe68}, Flowers and Itoh \\cite{fi76}, Urpin and Yakovlev \\cite{uy80}, and Timmes \\cite{timmes92}. These calculations were summarized by Potekhin et al.\\ \\cite{pbhy99}. The main result was that the $ee$ collisions are mainly negligible, except for a hot low-density plasma of light elements (from hydrogen to carbon) with the temperature a few times lower than the electron degeneracy temperature (in that case the $ee$ collisions did affect $\\kappa_{e}$ but never dominated as long as the electrons were degenerate; see Lampe \\cite{lampe68}; Urpin and Yakovlev \\cite{uy80}).  However, all these calculations have neglected an important effect of the Landau damping of the $ee$ interaction owing to the exchange of transverse plasmons. In the context of transport properties of dense matter this effect was studied by Heiselberg and Pethick \\cite{hp93} for degenerate quark plasmas. Similar effects have recently been analyzed by Jaikumar, Gale and Page \\cite{jgp05} for neutrino bremsstrahlung radiation in the $ee$ collisions. Here we reconsider $\\kappa_{ee}$ including the effects of the Landau damping and show that $\\kappa_{ee}$ is actually much more important than thought before. ",
        "conclusions": "\\label{conclusions}  We have reconsidered the electron thermal conductivity $\\kappa_{ee}$ of degenerate electrons produced owing to the $ee$ collisions taking into account the Landau damping due to the exchange of transverse plasmons (following the calculations of kinetic properties of quark plasma by Heiselberg and Pethick \\cite{hp93}). The Landau damping has been neglected in all previous calculations of $\\kappa_{ee}$. We have analyzed the four regimes of the $ee$ collisions in the degenerate electron gas (Section \\ref{regimes}, Table \\ref{tab:regimes}) and obtained analytic expressions for $\\kappa_{ee}$ which accurately approximate the results of numerical calculations of $\\kappa_{ee}$ in wide ranges of the temperature and density of the matter. These results can be applied to study thermal conduction of degenerate electrons in metals, semiconductors, in degenerate cores of evolved stars and white dwarfs, and in envelopes of neutron stars.  Our main conclusions are the following.  (1) The Landau damping strongly modifies $\\kappa_{ee}$ in a relativistic degenerate electron gas, at densities $\\rho \\gtrsim 10^6$~g~cm$^{-3}$, but it is also quite noticeable at lower $\\rho$ (for instance, at $\\rho=10^5$~g~cm$^{-3}$, see Figure \\ref{fig:rho56}). The Landau damping increases the $ee$ collision rate and decreases $\\kappa_{ee}$, increasing the contribution of the $ee$ collisions into the total electron thermal conductivity $\\kappa_{e}$.  (2) The most dramatic effect of the Landau damping on $\\kappa_{ee}$ takes place at $\\rho \\gg 10^6$~g~cm$^{-3}$ for temperatures $T$ much below the electron plasma temperature $T_{pe}$ (the regime IV, Table \\ref{tab:regimes}). In this case $\\kappa_{ee}$ shows a non-Fermi-liquid behavior; it becomes temperature independent and is described by the asymptotic expression (\\ref{kappa4}).  (3) The conductivity $\\kappa_{ee}$ becomes comparable to the electron thermal conductivity $\\kappa_{ei}$ provided by the electron-ion collisions (and gives thus a noticeable contribution into the total conductivity $\\kappa_e$) in a warm plasma of light (low-$Z$) ions at temperatures $T$ a few times lower than $T_{pe}$ (Figure \\ref{fig:rho56}). These conditions are typical for degenerate cores of white dwarfs and giant stars where thermonuclear burning of light elements can occur (particularly, in the vicinity of the carbon ignition curve).  (4) The conductivity $\\kappa_{ee}$ dominates over $\\kappa_{ei}$ and determines the total electron thermal conductivity $\\kappa_e$ of the dense pure cold-catalyzed matter at $T \\sim 10^7$ K and $\\rho \\gtrsim 10^{12}$~g~cm$^{-3}$ (in the inner crust of a cold neutron star, Figure \\ref{fig:econd}). At these $T$ and $\\rho$ the conductivity $\\kappa_{ee}$ is important even for impure cold-catalyzed matter. It can affect the propagation of thermal waves, excited in the inner neutron star crust during pulsar glitches, to the pulsar surface.  The electron conductivity $\\kappa_{e}$ operates also in neutron star cores, at $\\rho \\gtrsim 1.5 \\times 10^{14}$ g~cm$^{-3}$. This conductivity should also be reconsidered taking into account the Landau damping. Our present results cannot be directly applied to this case because in the cores the electrons collide efficiently at least with degenerate electrons, muons, and protons, and these collisions deserve a special study. We intend to analyze them in a subsequent publication."
    },
    "0608/astro-ph0608621_arXiv.txt": {
        "abstract": "We have investigated the relation between the orbital periods ($\\pb$) and the spin periods ($\\ps$) of wind-fed X-ray pulsars in high-mass X-ray binaries (HMXBs), based on population synthesis calculations of the spin evolution of neutron stars during the pre-HMXB stage. We show that most of the neutron stars either have steady accretion or still reside in the radio pulsar phase when the donor star starts evolving off the main sequence. In the former case the values of $\\ps$ can be decelerated to be $\\sim 10^2-10^3$ s depending on $\\pb$. We briefly discuss the possible origin of the $\\ps-\\pb$ correlation in Be/X-ray binaries, and the existence of HMXBs with main sequence donors. We have also investigated the evolution of the inclination angle between the magnetic and spin axes of neutron stars in a massive binary, suggesting secular alignment of the magnetic and spin axes during their evolution. ",
        "introduction": "High-mass X-ray binaries (HMXBs) can be roughly divided into two types: a supergiant star and a Be star are contained respectively. The X-ray source often is a pulsar and powered by accretion of the material offered by the companion star. In the supergiant systems, either Roche-lobe or stellar wind accretion occurs, while in the Be systems commonly only the latter process takes place since the Be star is well inside its Roche lobe \\citep{tau05}.  The relation between the spin periods $\\ps$ and the orbital periods $\\pb$ of HMXBs was first studied by \\citet{cor84,cor85,cor86}, who pointed out that there may be a $\\ps\\propto P_{\\rm orb}^{4/7}$ correlation for neutron stars in systems with radially expanding winds, while $\\ps\\propto P_{\\rm orb}^{2}$ for Be systems. \\citet{van87} noticed that the spin and orbital periods agree reasonably well with the former correlation for the supergiant systems in which the X-ray source is powered by stellar wind accretion. The $\\ps-\\pb$ relation has been studied by \\citet{ste86} in terms of the equilibrium spin period, where the corotation radius equals to the magnetospheric radius defined by equating the ram pressure of accreting matter with the magnetic pressure of the dipole field of the neutron star \\citep[e.g.,][]{dav73,lam73}. However, in order to obtain quantitative agreement between the predicted and observed spin periods for given orbital periods, the assumption of the equilibrium spin periods requires a mass accretion rate two orders of magnitude lower than observed for supergiant systems \\citep{ste86}. \\citet{wat89} presented a comprehensive explanation for the $\\ps-\\pb$ relation for HMXBs. For supergiant systems they suggested that the difference between the required and the observed mass accretion rates can be explained by assuming that the present-day periods are the equilibrium spin periods for the stellar wind of the optical star when it was still on main sequence (MS). The $\\ps-\\pb$ correlation for Be/X-ray binaries was suggested to result from the equilibrium spin realized by interaction between the neutron star magnetic field and the Be star's equatorial wind - the wider orbits expose to the neutron stars to a lower wind density on average and hence lower accretion rate. The latter point was criticized by \\citet{king91}, who proposed that the observed $\\ps-\\pb$ relation was probably the result of an earlier evolutionary stage, as for supergiant binaries. By calculating the angular momentum carried by the Be star's wind material, \\citet{li96} showed that for Be/X-ray binaries in relatively narrow orbits ($\\pb\\lsim$ 100 days), the equilibrium spin could be reached by angular momentum transfer via a disk, formed in the equatorial wind of the Be stars; in wider binaries, the low angular momentum of the wind matter prevents the existence of an accretion disk, and the neutron star's spin remains close to its previous equilibrium value at an earlier evolutionary stage. \\citet{zha04} have calculated the detailed spin evolution of a wind-fed neutron star in OB binaries prior to the HMXB phase, by simulating the time-development of the mass loss rate and radius expansion of a 20 $M_{\\odot}$ donor star.  The above-mentioned investigations are usually either semi-analytic or numerical only for individual cases. To better understand the distribution of HMXBs in the $\\ps-\\pb$ diagram, an evolutionary population synthesis incorporated with neutron star spin evolution is needed. In the present paper, we have conducted a Monte Carlo population synthesis study of the spin evolution of a neutron star in a massive binary. Due to the large theoretical uncertainties on the formation and evolution of HMXBs, our work is limited to the spin evolution of neutron stars before the HMXB phase. The theoretical considerations are described in \\S 2. The calculated results are presented in \\S 3 and their implications for accretion processes in HMXBs are discussed in \\S 4. ",
        "conclusions": "We have calculated the spin evolution of neutron stars in a massive binary. The main ideas in this work are of two folds. First, to appear as X-ray binaries, the spin periods of neutron stars should be longer than the current (instantaneous) equilibrium periods. Second, to satisfy this condition, the neutron stars should be spun-down enough during the previous propeller phase when the companion stars had relatively weak winds. These two requirements set constraints on the spin period distribution in wind-fed X-ray pulsars. Since $\\dot{M}$ depends on the orbital separations, this may result in the particular distributions of the supergiant and Be HMXBs in the $\\ps-\\pb$ diagram. The numerical results presented in Figs.~2 and 3, roughly accounting for the main features of the $\\ps-\\pb$ distribution of HMXBs, lend support to the original suggestions by \\citet{cor84}, \\citet{ste86}, and \\citet{wat89}. Our calculations of the alignment during spin-down imply that most HMXBs may have relatively small inclination angles.  We note that our preliminary results are subject of lots of uncertainties and simplified treatment adopted. In case of isotropic wind accretion, a considerable fraction of neutron stars can reach $P_{\\rm eq}$ when the companion star is still on the MS. The detail number is determined not only by the binary evolutionary processes, but also critically by the propeller mechanism. The spin-down torque (Eq.~[4]) adopted in this work is among the most efficient ones \\citep[][and references therein]{jia05}, and the results should be taken as to be the most optimistic. For less efficient spin-down torques ones, the number of binaries in the propeller phase will obviously increase due to the longer spin-down timescales. The magnitude of $P_{\\rm eq}$ depends on the mass loss rate, the wind velocity, and the magnetic field strength of the neutron star. With the mass loss rates given by Eq.~(9) and typical magnetic fields of $\\sim 3\\times 10^{12}$ G, the neutron star's spin can be decelerated to be $\\sim 10^2-10^3$ s within $t_{\\rm MS}$ of the secondary star. The very long period $\\ps=10^4$ s of 2S 0114$+$650 \\citep{hal00}, however, may be explained by a ultra-strong initial magnetic field ($B\\gsim 10^{14}$ G), in order to be spun-down efficiently by the propeller effect \\citep{li99}.  The structure of the Be star's winds is much more complex and variable than in supergiant systems, and the simple propeller spin-down may not be applicable to this type of systems. It is likely that the $\\ps-\\pb$ distribution and correlation of Be/X-ray binaries result from the balance between the spin-up during outbursts and spin-down during quiescence. Assume that the spin-up and spin-down torques are $\\dot{M}_{\\rm o}(GMR_{\\rm mo})^{1/2}$ and $-\\dot{M}_{\\rm q}R_{\\rm mq}^2\\Omega_{\\rm s}$ respectively, we can derive the following expression for the equilibrium period \\citep[see also][]{men99}, \\begin{equation} P_{\\rm eq}=2\\pi(\\frac{GM}{R_{\\rm mo}^3})^{-1/2}(\\frac{\\dot{M}_{\\rm q}R_{\\rm mq}^2}{\\dot{M}_{\\rm o}R_{\\rm mo}^2})(\\frac{t_{\\rm q}}{t_{\\rm o}})\\propto \\dot{M}_{\\rm o}^{-3/7}(\\frac{\\dot{M}_{\\rm q}}{\\dot{M}_{\\rm o}})^{3/7}(\\frac{t_{\\rm q}}{t_{\\rm o}}). \\end{equation} Here the subscripts o and q denote quantities evaluated during outbursts and quiescence, respectively. All the three terms on the right hand side of Eq.~(13) are likely to increase with $\\pb$, which might account for the observed $\\ps-\\pb$ correlation of Be/X-ray binaries.  Our calculated results also indicate the existence of HMXBs with MS donors in the Galaxy. These systems, as the progenitors of supergiant HMXBs, are less luminous than the latter (but with similar spin periods). Figure 4 suggests that there could be a few hundred of these sources in the Galaxy with X-ray luminosities ranging from $\\sim 10^{33}$ to $10^{35}$ ergs$^{-1}$ (If the propeller effect is not considered, the total number of X-ray binaries, most of which are Be/X-ray binaries, can reach a few thousand). The HMXB 4U2206$+$54 is likely to be the prototype of this kind of sources \\citep{ribo05}. A number of papers have also suggestted that some neutron stars receive low kick speeds of $\\leq50$ kms$^{-1}$ at birth \\citep{pfahl02, pod04, dew05}. If all the neutron stars are born with such small kicks, our calculations show that there would be about $4-5$ times more X-ray sources produced."
    },
    "0608/astro-ph0608598_arXiv.txt": {
        "abstract": "{} {We present optical and near-infrared spectral evolution of the Galactic nova V5114 Sgr (2004) during few months after the outburst.} { We use multi-band photometry and line intensities derived from spectroscopy to put constrains on the distance and the physical conditions of the ejecta of V5114 Sgr. }{The nova showed a fast decline ($t_2 \\simeq 11$~days) and spectral features of \\ion{Fe}{ii} spectroscopic class. It reached $M_V = -8.7\\pm 0.2$~mag at maximum light, from which we derive a distance of $7700 \\pm 700$~kpc and a distance from the galactic plane of about 800 pc. Hydrogen and Oxygen mass of the ejecta are measured from emission lines, leading to $\\sim 10^{-6}$ and $10^{-7}$M$_\\odot$, respectively.  We compute the filling factor of the ejecta to be in the range $0.1$ -- $10^{-3}$. We found the value of the filling factor to decrease with time. The same is also observed in other novae, then giving support to the idea that nova shells are not homogeneously filled in, rather being the material clumped in relatively higher density blobs less affected by the general expanding motion of the ejecta.} {} ",
        "introduction": "Nova Sgr 2004 (V5114 Sgr) was independently discovered  by Nishimura (2004) and Liller (2004) on Mar15.8~UT and Mar17.3~UT, respectively. West (2004) provided precise coordinates R.A. = 18\\hbox{$^{\\rm h}$}19\\hbox{$^{\\rm m}$}32\\fs29, Decl. = $-$28\\degr36\\arcmin35\\farcs7 (gal. coord. $l=3$\\fdg$9$ $b=-6$\\fdg$3$). Early spectroscopy on Mar 18.3 UT by Della Valle et al. (2004) confirmed  this object to be a classical nova caught near maximum light. Here, we present photometric and spectroscopic observations taken at McDonald Observatory, Cerro Tololo Inter-American Observatory (CTIO), ESO-La Silla, and Lick Observatory. The spectra from ESO-LaSilla have been obtained as part of the Target of Opportunity campaign for the observation of Classical Novae in the Galaxy and in the Magellanic Clouds. The paper is organized as follows: in Sect. 2 we present the analysis of the photometric and spectroscopic data, in Sects. 3 and 4 we analyze the evolution of the light curve and spectra of V5114 Sgr, in Sect. 5 we discuss the effects of interstellar absorption. In Sect. 6 we derive the physical parameters of the nova ejecta.  Summary and conclusions are given in Sect. 7.   \\begin{figure*} \\centering \\includegraphics[width=18cm, angle=-90, scale=0.8]{edefig1.eps} \\caption{ \\textbf{Upper left panel:}~ $V$ band light curve. Filled symbols represent our data points, empty symbols represent IAUC data. Vertical dotted lines represent maximum light, $t_2$, and $t_3$. Data points are affected by errors $\\lesssim 0.04$~mag in all bands but in $U$, where the error is 0.1~mag. \\textbf{Lower left panel: } Evolution of different colors: filled squares represent $B-V$, empty triangles  $U-B$, crosses $V-I$ and empty squares $V-R$. \\textbf{Upper right panel: } $J$ band light curve. \\textbf{Lower right panel: } Evolution of near-infrared colors: triangles represent $J-H$ and squares $H-K$. } \\label{Lightcurves} \\end{figure*} ",
        "conclusions": "V5114 Sgr is an \\ion{Fe}{ii} nova that occurred in the bulge of the Milky Way.  The rate of decline characterizes V5114 Sgr as a borderline object between the fast ($t_2 < 12$days) and slow ($t_2 > 12$ days) classes (Della Valle \\& Livio \\cite{MDVL98}).  These authors have shown that He/N and \\ion{Fe}{ii}b (i.e. {\\it \\ion{Fe}{ii} broad}) novae belong to the fast class and are preferentially concentrated towards the galactic disc, i.e. at small $z$ above the galactic plane ($z<200$~pc), while \\ion{Fe}{ii} novae belong to both slow and fast classes and  are observed both in the disk and in the bulge, extending up to $\\sim 1$~kpc. V5114 Sgr located at about 0.8~kpc above the galactic plane in the direction of the galactic bulge is an \\ion{Fe}{ii} nova which does not represent an exception to this scenario (Della Valle et al. \\cite{MDVetal92}).  Spectroscopic observations showed a dramatic change in the overall appearance of the spectrum of the nova at about 30 days after maximum light (in coincidence with the entrance in the ``auroral phase''): permitted lines start fading at a different rate, P-Cyg profiles disappear, and velocities reach a ``plateau'' phase and $U-V$ and $J-K$ colors increase. The lack of detection of [\\ion{Ne}{iii}] line hints for a ``standard'' evolution of the nebular spectrum (see Williams et al. 1994).  The very high values of optical depth in the \\ion{O}{i}$\\lambda6300$ suggest very high densities for the zones of the ejecta where these lines are formed.  We have derived the filling factor in the range $\\sim 7.1 \\times 10^{-2}$ to $7.9 \\times 10^{-4}$.  Comparing these values with other values reported in the literature (see Table \\ref{filfactab}), two facts emerge: a) the filling factors in nova ejecta are definitely smaller than 1, likely close to 0.1, during the early stages; b) these values decrease by 1--3 orders of magnitude with time. This fact indicates that the volume of the expanding shell (computed with Eq.\\ref{volume}) increases with time more rapidly than the volume actually occupied by most of the ejected material. In other words, the decreasing trend exhibited by the filling factor suggests that the ejected matter tends to remain clumped in sub-structures having higher density than the average density characterizing the expanding shell. This is consistent with the $[$\\ion{O}{i}$]$ $\\lambda\\lambda$6300, 6364 ratio $< 2$ (Table \\ref{Physpars1}), which suggests that $[$\\ion{O}{i}$]$ lines should be formed in very dense small blobs of neutral material embedded within the ionized shell (see paragraph 6). However, recently Williams \\& Mason (2006) proposed an alternative interpretation of this behavior, assuming that the $[$\\ion{O}{i}$]$ lines arise in regions of high magnetic field, and their intensity and profile are modified by Quadratic Zeeman Effect. \\begin{table} \\caption{Physical parameters for V5114 Sgr: $\\epsilon$ and the hydrogen mass.} \\label{Physpars2} \\centering \\begin{tabular}{c c c} \\hline\\hline Date\t& $\\epsilon$\t& $M_H$ (M$_\\odot$)\\\\ \\hline Jun26  & 7.08E-2 & 3.02E-5\t \\\\ Jun28  & 7.04E-2 & 2.91E-5\t \\\\ Aug14  & 3.97E-3 & 6.49E-6\t \\\\ Sep26  & 7.92E-4 & 1.10E-6\t \\\\ \\hline \\end{tabular} \\end{table}  Computed oxygen and hydrogen masses are in the ranges $ 1.9 \\times 10^{-5} - 2.4 \\times 10^{-7} $M$_\\odot$ and $ 3.0 \\times 10^{-5} - 1.1 \\times 10^{-6} $M$_\\odot$. This high mass ratio is close to the upper limit for classical novae shown in Warner (\\cite{warner}).  \\begin{table}[h] \\caption[]{Comparison between filling factors in literature.} \\label{filfactab} \\begin{tabular}{l l l} \\hline Nova\t&\tFilling factor value\t&\tReference\t\\\\ \\hline T Pyx\t\t&\t$10^{-2}$--$10^{-5}$ & Shara et al. (\\cite{Shara}) \\\\ Nova Vel 1999\t&\t0.1\t\t& \tDella Valle et al.(\\cite{MDVPDW}) \\\\ Nova SMC 2001\t&\t0.1 -- $10^{-3}$&\tMason et al. (\\cite{Mason2005}) \\\\ Nova LMC 2002\t&\t$10^{-2}$ -- $10^{-4}$&\tMason et al. (\\cite{Mason2005}) \\\\ \\hline {\\bf V5114 Sgr } & {\\bf $0.1$ -- $10^{-3}$}\t& {\\bf \\textit{this paper} }\\\\ \\hline \\end{tabular} \\end{table}"
    },
    "0608/astro-ph0608551_arXiv.txt": {
        "abstract": "{} {To determine whether or not gas-phase chemical models with homogeneous and time-independent physical conditions explain the many observed molecular abundances in astrophysical sources, it is crucial to estimate the uncertainties in the calculated abundances and compare them with the observed abundances and their uncertainties. Non linear amplification of the error and bifurcation may limit the applicability of chemical models. Here we study such effects on dense cloud chemistry. } { Using a previously studied approach to uncertainties based on the representation of rate coefficient errors as log normal distributions, we attempted to apply our approach using as input a variety of different elemental abundances from those studied previously. In this approach, all rate coefficients are varied randomly within their log normal (Gaussian) distribution, and the time-dependent chemistry calculated anew many times so as to obtain good statistics for the uncertainties in the calculated abundances.} {Starting with so-called ``high-metal'' elemental abundances, we found bimodal rather than Gaussian like distributions for the abundances of many species and traced these strange distributions to an extreme sensitivity of the system to changes in the ratio of the cosmic ray ionization rate $\\zeta_{\\rm He}$ for He and that for molecular hydrogen $\\zeta_{\\rm H_2}$.  The sensitivity can be so extreme as to cause a region of bistability, which was subsequently found to be more extensive for another choice of elemental abundances.  To the best of our knowledge, the bistable solutions found in this way are the same as found previously by other authors, but it is best to think of the ratio $\\zeta_{He}/\\zeta_{H_{2}}$ as a control parameter perpendicular to the ''standard'' control parameter $\\zeta/n_{\\rm H}$. } {} ",
        "introduction": "The equations describing the chemical evolution of quiescent cores (also known as molecular clouds) are highly non-linear and can result in extreme sensitivity to the initial conditions or to some of the parameters. For example, the elemental ratio of carbon to oxygen is well known to be an important parameter for dense cloud chemistry \\citep{1998ApJ...501..207T}. Another well-known consequence of this non-linearity is the presence of bistable regions for  some ranges of model parameters (temperature, density, comic-ray ionization rate, rate coefficients and elemental abundances).  The phenomenon of bistability  is characterized by two stable solutions for chemical abundances at steady-state for the same set of parameters. Bistability in dark cloud conditions, i.e. at low temperature, was first discovered by \\citet{1993ApJ...416L..87L} and subsequently studied by a large number of authors \\citep{1995A&A...297..251L,1995A&A...296..779S,1998A&A...334.1047L,2000RSPTA.358.2549P,2003A&A...399..583C,2006ApJ...645..314B}. The two solutions of this bistability are characterized by a large difference in the C/CO and O/O$_2$ ratios and in the ionization fraction. For this last reason, they were named the High and Low Ionization Phases (HIP and LIP). \\citet{1995A&A...297..251L} showed that the region of bistability can be mapped out by variations in density, temperature, cosmic-ray ionization rate, and elemental depletions within specific ranges.   \\citet{2000RSPTA.358.2549P} later showed that even variations in rate coefficients can lead to the two solutions. Recently, \\citet{2006ApJ...645..314B} identified the chemical mechanisms of this phenomenon as a cycle involving the H$^+_3$, O$_2$, and S$^+$ species.  In two previous papers, we presented a Monte Carlo method to compute the theoretical error of abundances in chemical models due to uncertainties in rate coefficients and physical conditions \\citep{2005A&A...444..883W,2006A&A...451..551W}. The errors in the rate coefficients are defined by a log-normal distribution. As a consequence, the resulting abundances  at steady-state usually follow a Gaussian distribution \\citep{2005A&A...444..883W}. When applying this method, we found that the chemical abundances are very sensitive to the ratio between the ionization rates of helium $\\zeta_{\\rm He}$ and molecular hydrogen $\\zeta_{\\rm H_2}$ caused by cosmic-ray particles. Characterized by a variation of several orders of magnitude in some abundances for small variations of $\\zeta_{\\rm He} / \\zeta_{\\rm H_2}$, this sensitivity can lead to bimodal rather than Gaussian distributions for the abundances. For certain ranges of the parameters, it can even lead to bistability, with the ratio $\\zeta_{\\rm He} / \\zeta_{\\rm H_2}$ appearing to be a control parameter.  This paper, which describes the sensitivity and some of its consequences, is organized as follows. In section 2, we briefly present the chemical model and the uncertainty method used for this analysis. In a somewhat unorthodox way, we have decided to present, in section 3, the manner in which the sensitivity was found while studying the use of so-called ``high-metal'' elemental abundances. Section 4 shows the influence of the elemental abundances and the chemical network used  on the sensitivity. In section 5, we present and discuss the phenomenon of bistability as obtained with the variation of $\\zeta_{\\rm He} / \\zeta_{\\rm H_2}$. In section 6,   the chemical reactions involved in this sensitivity/bistability are elucidated. Finally, the last section presents our conclusions. ",
        "conclusions": "Using a Monte Carlo uncertainty method to compute the theoretical error of molecular abundances in chemical models, we found a strong sensitivity of the steady-state abundances to the ratio between the ionization rate coefficient of helium $\\zeta_{\\rm He}$ and that of molecular hydrogen $\\zeta_{\\rm H_2}$ by cosmic rays. This sensitivity to the ratio has the consequence of changing the abundance of some species by several orders of magnitude for small variations of $\\zeta_{\\rm He}$/$\\zeta_{\\rm H_2}$. Lower values of $\\zeta_{\\rm He}$/$\\zeta_{\\rm H_2}$ are characterized by large abundances of smaller molecules such as O$_2$, H$_2$O, SO, etc. whereas high values of $\\zeta_{\\rm He}$/$\\zeta_{\\rm H_2}$ lead to large abundances of atoms. If the sensitivity is strong enough, it can result in the phenomenon of bistability, which we located at low densities for two sets of elemental abundances - the low- and high-metal cases, and over a wide range of densities for the case of the so-called intermediate-metal abundances, in which sulfur plays a particularly prominent role.  The bistability has been determined to be the same as studied earlier by a variety of authors with other control parameters; it is perhaps simplest to consider our control parameter --  $\\zeta_{\\rm He}$/$\\zeta_{\\rm H_2}$-- as orthogonal to one of the others - the ratio of the overall ionization rate parameter $\\zeta$ to the density $n$ --  so that the bistability can be imagined to exist in a region in three-dimensional space rather than in the customary two-dimensional hysteresis picture.  It is interesting to speculate that the sensitivity detected here via our uncertainty analysis is a necessary but not sufficient condition for bistability because only the most extreme sensitivity, in which the bimodal distributions detected have no solutions in between them, leads to bistability. Unfortunately, it is not guaranteed that our type of uncertainty analysis can always locate a bistable region since finding an extreme sensitivity may require large uncertainties in some parameters, such as the gas density, to take us from our standard parametric values to the bistable region.  Several questions arise when considering the dramatic effects of small $\\zeta_{\\rm He}$/$\\zeta_{\\rm H_2}$ variations on the gas-phase chemistry. First, even if one assumes a given flux and energy distribution of the cosmic rays, the determination of $\\zeta_{\\rm He}$/$\\zeta_{\\rm H_2}$ suffers some significant uncertainties, with the role of secondary electrons particularly important.  Besides these  difficulties, one might also have to consider variations of the ratio $\\zeta_{\\rm He}$/$\\zeta_{\\rm H_2}$ in space, such as between one cloud and another, or between the center and the edges of a same cloud,  and in time. Changes in the ratio imply important variations in the spectral distribution of the cosmic rays \\citep{1973ApJ...186..859G}. Such variations might be produced by the attenuation of the cosmic rays inside a cloud or  the proximity of supernovae in star forming clouds. Whether such variations are expected or not and how they could affect the chemistry should be considered in the future. Finally, a knowledge of the helium elemental abundance appears to be crucial since the sensitivity region depends on this parameter."
    },
    "0608/astro-ph0608284_arXiv.txt": {
        "abstract": "{} % {We investigate the non--Local Thermodynamic Equilibrium (non--LTE) line formation of neutral carbon in late-type stars in order to remove some of the potential systematic errors in stellar abundance analyses employing C\\,{\\sc i} features.} {The statistical equilibrium code {\\small MULTI} was used on a grid of plane-parallel 1D {\\small MARCS} atmospheric models.} {Within the parameter space explored, the high-excitation C\\,{\\sc i} lines are stronger in non--LTE due to the combined effect of line source function drop and increased lower level population; the relative importance of the two effects depends on the particular combination of \\mbox{$T_{\\rm{eff}}$}, \\mbox{log $g$}, \\mbox{[Fe/H]}\\, and \\mbox{[C/Fe]} and on the analysed C\\,{\\sc i} line. As a consequence, the non--LTE abundance corrections are negative and can be substantially so, for example $\\sim -0.4$~dex in halo turn-off stars at \\mbox{[Fe/H]}$\\sim -3$. The magnitude of the non--LTE corrections is rather insensitive to whether inelastic H collisions are included or not.} {Our results have implications on studies of nucleosynthetic processes and on Galactic chemical evolution models. When applying our calculated corrections to recent observational data, the upturn in [C/O] at low metallicity might still be present (thus apparently still necessitating contributions from massive Pop. III stars for the carbon production), but at a lower level and possibly with a rather shallow trend of $\\sim -0.2$~dex/dex below [O/H]$\\sim -1$.} ",
        "introduction": "Carbon is synthesised through the $3\\alpha$ nuclear reaction in stellar interiors (Burbidge et al. 1957; Arnett 1996; Pagel 1997). It is one of the most abundant elements in the Universe, together with hydrogen, helium, oxygen and neon. It can be found in a variety of compounds, so that it forms the basis of organic chemistry and of life (at least as we know it). Carbon is also important for dust formation processes in the interstellar medium (e.g.  Dopita \\& Sutherland 2003). Since it has a quite high ionization potential (11.26~eV), it remains mostly neutral in the photosphere of solar-type stars. Stellar spectroscopists thus rely on high-excitation permitted C\\,{\\sc i} features and on the forbidden [C\\,{\\sc i}] line at $872.7$~nm, as well as on various molecular species (mainly CH and C$_2$), to derive solar and stellar C abundances (e.g. Asplund et al. 2005a).  A number of studies have investigated the carbon content of disk stars in the Galaxy (e.g. Andersson \\& Edvarsson 1994; Gustafsson et al. 1999), arguing that [C/Fe]\\footnote{The usual spectroscopic notation [A/X]$=\\log({\\rm N_{A}/N_{X} })_{*}-\\log({\\rm N_{A}/N_{X} })_{\\odot}$ is used throughout} is slowly decreasing with time (i.e. with increasing metallicity) and probably mostly contributed by massive stars (MSs). In particular, Gustafsson et al. rule out a {\\it main} origin of carbon in low-mass stars while they argue that the contribution of intermediate-mass stars to the production of carbon in the Galaxy is unclear. Takeda \\& Honda (2005) have recently studied the CNO abundances of 160 disk F, G, and K dwarfs and subgiants with $-1<$\\mbox{[Fe/H]}$<+0.4$ to look for any differences between the abundances of stars with and without planets. In particular, when taking their calculated non--LTE corrections into account, they also found that [C/Fe] progressively decreases with increasing metallicity for these galactic disk stars, by around $\\sim 0.2$~dex per dex. A similar slope has been found for [C/Fe] by Reddy et al. (2003, 2006) in their analysis of a number of elements in a large sample of F and G dwarfs in the thin- and thick-disk. In contrast with the above authors, Bensby \\& Feltzing (2006) find essentially no trend of [C/Fe] versus \\mbox{[Fe/H]}\\, for metal-rich thick- and thin-disk stars and conclude that the carbon enrichment at such metallicities is mainly due to low- and intermediate-mass stars (LIMSs), with massive stars still playing the role of main contributors at lower metallicities in the thick disk/halo.  The halo stellar population is less well studied. It is still currently a matter of debate whether for halo stars carbon is produced primarily in LIMSs (e.g. Chiappini et al. 2003), or in massive stars (e.g. Akerman et al. 2004). It could well be that contributions from both of these sources operating on different timescales are important. Akerman et al. were the first to derive a previously unrecognized [C/O] upturn at low metallicities. However, they warned about possible non--LTE effects at low metallicity for the high-excitation C\\,{\\sc i}\\, lines they employed and cautioned that their results might be prone to uncertainties due to their assumption that C\\,{\\sc i}-based abundances would be affected by the same non--LTE corrections as for O\\,{\\sc i}\\, (which they estimated for their sample by interpolating between the values given in Nissen et al. 2002) and thus that the [C/O] ratio would be unchanged with respect to LTE. They mentioned the possibility that non--LTE effects for the carbon and oxygen high-excitation lines might differ by an amount increasing with decreasing metallicity. At that stage, they were not able to exclude that the [C/O] values were being overestimated by $\\sim 0.2$~dex under the assumption of LTE. The idea that the upturn could thus be removed is in fact presented in some recent work (e.g. Takeda \\& Honda 2005; Fabbian et al. 2005; Asplund 2005; Bensby \\& Feltzing 2006). Spite et al. (2005) study a sample of extremely metal-poor giants to investigate the abundances of C, N, O and Li. They find that in their ``unmixed'' stars (which should reflect the abundances in the early Galaxy) the [C/Fe] ratio is remarkably constant (at a level of $\\sim +0.2-0.3$~dex) over the entire range $-4.0<$\\mbox{[Fe/H]}$<-2.5$ and that the [C/O] ratio is close to solar at the lowest metallicities.  According to the studies of Akerman et al. (2004) and Spite et al. (2005), the upturn in the [C/O] abundances of halo stars present at low metallicity could signal substantial early nucleosynthetic production of carbon in massive Pop. III stars. However it is worth noting that the inclusion of non--LTE/3D effects plays a substantial role in testing such hypothesis. For example, 3D effects affecting the C and O abundance indicators used in the study by Spite et al. (2005) might reach $\\sim -0.6$~dex or more for CH, while they likely remain around $-0.2$~dex for the [O\\,{\\sc i}] forbidden oxygen lines (Asplund 2005; see also Collet et al. 2006). This would then make their resulting [C/O] trend more or less flat at a level of $\\sim -0.6$~dex down to very low [O/H] ratios, implying that resorting to high yields from massive zero-metallicity stars might be premature.  On the theoretical side, low-metallicity fast-rotating models of massive stars (Meynet et al. 2006; see also Chiappini et al. 2006) predict important CNO yields through large mass loss during the red supergiant phase via stellar winds caused by efficient mixing and dredge-up enriching the stellar surface in heavy elements. These yields, especially if used in conjuction with a top-heavy IMF, would thus give large carbon abundances at low metallicity without introducing Pop. III stars. Carigi et al. (2005) have constructed various Galactic chemical evolution models with different C, N and O yields. The only models that are able to fit the carbon gradient are those including C yields that increase with metallicity due to stellar winds in massive stars and decrease with metallicity due to stellar winds in LIMSs. They argue that the fraction of carbon in the interstellar medium due to MSs and LIMSs would thus be dependent on time (with MSs dominating the C enrichment at early times, while around $12+\\log$(O/H)$\\sim 8$ the LIMSs formed in the halo should be at the end of their evolution and they would start contributing with C ejecta in a comparable amount to that of MSs, see also Bensby \\& Feltzing 2006) and on galactocentric distance (LIMSs giving a larger contribution at large r$_{\\rm GC}$), with similar contribution from both sources to the present carbon abundance in the interstellar medium of the solar vicinity. However, no attempt was made in that work to homogenize the sample of available observational data. Moreover, their ``best'' models are still not able to properly fit the N abundance values for metal rich stars and the C/O values at intermediate or low metallicity obtained from stellar observations. Gavil\\'{a}n et al. (2005) have recently suggested that a Galactic Chemical Evolution (GCE) model including yields from Woosley \\& Weaver (1995) for massive stars needs to be integrated with yields from LIMSs to properly reproduce the wealth of observational data on the [C/Fe] trend down to \\mbox{[Fe/H]}$\\sim -2.4$, suggesting that a downward revision of mass loss by stellar winds from massive stars is necessary and adding to the increasing evidence that LIMSs are significant contributors to the Inter Stellar Medium (ISM) carbon enrichment. Interestingly, Ballero et al. (2006) test several sets of yields for very massive zero-metallicity stars, using the two-infall chemical evolution model by Chiappini et al. (1997). They argue that the contribution of massive Population III stars to the Galactic evolution of carbon, nitrogen and iron would be negligible even at the lowest metallicities currently probed and that the existence of such objects could in any case not be proven or excluded on the basis of available data on known metal-poor halo stars, which likely formed as second-generation objects after the explosion of primordial supernovae.  Absorption systems give the opportunity to measure abundances in protogalactic structures at high redshift. For example, Becker et al. (2006) recently detected an overabundance of O\\,{\\sc i}\\, absorption features in spectra of background high-redshift quasars (QSOs) towards the line-of-sight of Lyman-$\\alpha$ forest clouds, deriving a mean $<[{\\rm C/O}]>= -0.31 \\pm 0.09$. Moreover, Pettini (private communication) and Erni et al. (2006) also find [C/O]$\\sim -0.3$ in metal-poor Damped Lyman-$\\alpha$ systems (DLAs) with [O/H]$\\sim -2.5$, in reasonable agreement with the values for Galactic stars (Akerman et al. 2004). These results are obviously precious in complementing Galactic studies at low \\mbox{[Fe/H]}. In particular, Erni et al. argue that enrichment from a population of massive (10-50 $M_{\\odot}$) Pop. III stars exploding as core-collapse supernovae (SNeII) and possibly as hyper-energetic (E$>10^{51}$~erg) so-called hypernovae (HNe) is the most likely scenario, in contrast to models using non-zero metallicity progenitors or other explosion mechanisms, such as super-massive (140-260 $M_{\\odot}$) pair-instability supernovae (PISNe) or Type Ia supernovae (SNeIa). Levshakov et al. (2006) derive a very low abundance of $^{13}$C towards a quasar at ${\\rm z}=1.15$. This does not support the enrichment of gas by fast-rotating, very low-Z massive stars (e.g. Meynet et al. 2006) and puts a tight bound to the possible amount of contributions from intermediate-mass AGB stars to the C nucleosynthesis in matter sampled by the line of sight to this system.  The still open problem of the origin of carbon in the early Universe constitutes the main motivation for the present non--LTE study. The quality of stellar spectroscopic observations possible at present, require us to make corresponding progress in models of stellar atmospheres and line formation processes in order to remove systematic errors. The impact on Galactic chemical evolution studies of such non--LTE calculations promises to be rewarding. ",
        "conclusions": "We have carried out non--LTE spectral-line formation calculations for C\\,{\\sc i} absorption features. The computations have been performed on a grid of 168 1D {\\small MARCS} stellar model atmospheres with varying parameters ($4500\\le$\\mbox{$T_{\\rm{eff}}$}$\\le7000$~K, $2\\le$\\mbox{log $g$}$\\le5$, $-3\\le$\\mbox{[Fe/H]}$\\le0$) and $-0.60\\le$[C/Fe]$\\le+0.60$). The main results are:  \\begin{itemize}  \\item We generally find large, negative non--LTE abundance corrections for high-excitation neutral carbon permitted lines, i.e. LTE abundances are too large. At solar metallicity, the dominant non--LTE effect is that of the line source function dropping below the Planckian mostly due to under-populated upper atomic levels in the transitions, but at low metallicity it is that of line opacity due to overpopulated lower levels  \\item We find substantial non--LTE abundance corrections for metal-poor halo turn-off stars: typically $\\sim -0.4$~dex at \\mbox{[Fe/H]}$\\sim -3$; these results are relatively immune to the adopted H collisional efficiency  \\item Using our resulting non--LTE corrections, we revisit the observational study of Akerman et al. (2004) and find that the [C/O] trend at low metallicity might still be suggesting an upturn indicating possible production of carbon through nucleosynthesis in Pop. III stars. We however warn that more observational data is needed - together with detailed non--LTE calculations for oxygen, to understand the magnitude of the abundance corrections applicable for this element at metallicities \\mbox{[Fe/H]}$\\apprle-2.5$ - in order to reduce the still large uncertainties and finally settle this issue  \\end{itemize}  Without doubt one of the necessary steps for obtaining reliable stellar abundances in future large scale surveys will be to perform a vast quantity of detailed non--LTE studies for an increased number of different chemical elements and stellar parameters. Very many precise atomic input data for collisional and radiative processes have to be known for realistic results. The use of sophisticated techniques for quantum-mechanical computations have led to accurate atomic radiative data fortunately becoming available in recent years to the astrophysical community, thanks to such large scale efforts as e.g. the Opacity and IRON Projects. We underline the necessity for more theoretical and laboratory studies aimed at obtaining such accurate atomic data and in particular at improving our knowledge on the correct treatment of the still poorly understood collisions with electrons and with hydrogen atoms. The effect of using 3D hydrodynamical model atmospheres should be investigated as the 3D non--LTE effects could prove to be even stronger than in 1D, due to the presence of large temperature inhomogeneities and to the different temperature structure in the optically thin atmospheric layers with respect to 1D hydrostatic models, especially for the low-metallicity regime (Asplund et al. 1999).  \\bigskip"
    },
    "0608/astro-ph0608259_arXiv.txt": {
        "abstract": "We have assembled a sample of \\burstnum\\ thermonuclear (type-I) X-ray bursts from observations of \\numbursters\\ accreting neutron stars by the {\\it Rossi X-ray Timing Explorer}, spanning more than ten years. The sample contains examples of two of the three theoretical ignition regimes (confirmed via comparisons with numerical models) and likely examples of the third. We present a detailed analysis of the variation of the burst profiles, energetics, recurrence times, presence of photospheric radius expansion, and presence of burst oscillations, as a function of accretion rate.  We estimated the distance for \\numsrc\\ sources exhibiting radius-expansion bursts, and found that in general the peak flux of such bursts varies typically by \\meanonly, We classified sources into two main groups based on the burst properties: both long and short bursts (indicating mixed H/He accretion), and consistently short bursts (primarily He accretion), and calculated the mean burst rate as a function of accretion rate for the two groups. The decrease in burst rate observed at $>0.06\\dot{M}_{\\rm Edd}$ ($\\ga2\\times10^{37}\\ \\eps$) is associated with a transition in the persistent spectral state and (as has been suggested previously) may be related to the increasing role of steady He-burning. We found many examples of bursts with recurrence times $<30$~min, including burst triplets and even quadruplets.  We describe the oscillation amplitudes for \\nsrcsearched\\ of the \\noscsrc\\ burst oscillation sources, as well as the stages and properties of the bursts in which the oscillations are detected. The burst properties are correlated with the burst oscillation frequency; sources spinning at $<400$~Hz generally have consistently short bursts, while the more rapidly-spinning systems have both long and short bursts. This correlation suggests either that shear-mediated mixing dominates the burst properties, or alternatively that the nature of the mass donor (and hence the evolutionary history) has an influence on the long-term spin evolution. ",
        "introduction": "\\label{intro}  Thermonuclear (type-I) X-ray bursts manifest as a sudden increase in the X-ray intensity of accreting neutron stars (NSs), to many  times brighter than the persistent level. Typical bursts exhibit rise times of between $\\la 1$ and $10$~s, and last from tens to hundreds of seconds (Fig. \\ref{examples}). These events are caused by unstable burning of accreted H/He on the surface of neutron stars in low-mass X-ray binary (LMXB) systems \\cite[e.g.][]{sb03}, in contrast to type-II bursts, which are thought to be caused by sudden accretion events \\cite[e.g.][]{lew93}. The H/He fuel for type-I bursts is accreted from the binary companion and accumulates on the surface of the neutron star, forming a layer several meters thick. The accreted material is compressed and heated hydrostatically, and if the temperature is sufficiently high any hydrogen present burns steadily into helium via the ``hot'' ($\\beta$-limited) carbon-nitrogen-oxygen (CNO) process.  After between $\\sim$1 and several tens of hours, the temperature and density at the base of the layer become high enough that the fuel ignites, burning unstably and spreading rapidly to consume all the available fuel on the  star in a matter of seconds. Such bursts have been observed to date from more than 70 sources \\cite[e.g.][]{zand04b}.  The burst X-ray spectrum is generally consistent with a blackbody of color temperature $T_{\\rm bb}=2$--3~keV.  Time-resolved spectral fits give evidence for an initial rise in $T_{\\rm bb}$ followed by a more gradual decrease following the burst peak, giving an approximately exponential decay in X-ray brightness back to the persistent level.  This is naturally interpreted as heating resulting from the initial fuel ignition, followed by cooling of the ashes once the available fuel is exhausted. The primary evidence that the energy source for type-I bursts is thermonuclear comes from comparisons of the time-integrated persistent and burst flux.  The ratio $\\alpha$ of the integrated persistent flux to the burst fluence is the standard measure of the relative efficiency of the two processes. Early in the study of type-I bursts it was determined that the energy derived from accretion was between 40 and a few hundred times greater than the energy liberated during the bursts. These values are comparable to those predicted assuming that the burst energy arises from nuclear burning.  \\begin{figure} \\epsscale{1.2} \\plotone{f1.eps} \\caption{Example light curves of bursts observed by \\xte. The top panel shows a long burst from GS~1826$-$24 on 1998 June  8 04:11:45 UT. The lower left panel shows a burst observed from 4U~1728$-$34 on 1999 June 30 19:50:14 UT, while the burst at lower right was observed from 4U~2129$+$12 in the globular cluster M15 on 2000 September 22 13:47:41 UT. The persistent (pre-burst) level has been subtracted (dotted line). Note the diversity of burst profiles, which arises in part from variations in the fuel composition; bursts with a slow rise and decay are characteristic of mixed H/He fuel, while bursts with much faster rises likely burn primarily He.  Both bursts in the lower panels exhibited photospheric radius-expansion. \\label{examples} } \\end{figure}  Numerical models of unstable nuclear burning on the surface of a neutron star reproduce the observed rise times (seconds), durations (minutes), recurrence times (hours), and total energies of the bursts \\cite[$10^{39}$--$10^{40}$~ergs; e.g.][]{fhm81, aj82, fl87, fuji87, bil98a,cb00,ramesh03,woos03}. The frequency, strength, and time  scales of thermonuclear bursts all depend on the composition of the burning material, as well as the metallicity (here referring to the CNO mass fraction, $Z_{\\rm CNO}$) of the matter accreted onto the neutron star; the amount of hydrogen burned between bursts; and the amount of fuel left-over from the previous burst. Variations from source to source are also expected because of differences in the core temperatures of the neutron stars and the  average accretion rate onto the surface \\citep{aj82, fl87,ramesh03}.  The recently discovered class of extremely long-duration bursts or ``super'' bursts are also thought to arise from thermonuclear processes. The fuel for these bursts is probably carbon rather than H/He, giving distinctly different time scales, recurrence times and energetics \\cite[][]{stroh02,cumming04}. However, the predicted temperatures in the fuel layer are too low to give carbon bursts with the observed fluences, suggesting that the cooling in the crust may be less efficient than previously thought \\cite[]{cumming06}. Superbursts have been detected from around 10\\% of the Galactic X-ray burst population, with recurrence times estimated at 1.5~yr \\cite[]{zand04c}, and tend to quench the regular type-I bursts for weeks--months afterwards. The connection with intermediate-duration ($\\sim30$~min) events observed from a few systems is not clear \\cite[]{zand04a}.  \\subsection{Bursts as a Function of Accretion Rate}  Theoretical ignition models for H- and He-burning thermonuclear bursts predict how burst properties in an individual system change as the accretion rate onto the neutron star varies \\citep[e.g.,][]{fhm81,bil98a,ramesh03}. Several regimes of thermonuclear igntion may be identified, depending upon the local accretion rate ($\\dot m$\\footnote{Following the usual convention, we refer to the accretion rate per unit area as $\\dot{m}$, and the total accretion rate integrated over the neutron star as $\\dot{M}$.}), which is usually expressed as a fraction of the local Eddington rate $\\dot{m}_{\\rm Edd}$ ( $8.8\\times10^4\\ {\\rm g\\,cm^{-2}\\,s^{-1}}$, or $\\equiv1.3\\times10^{-8}\\ M_\\odot\\,{\\rm yr^{-1}}$ averaged over the surface of a 10~km NS). However, the quantitative values of the accretion rates separating these regimes of burning are a matter of some debate, as we outline below, and summarize in Table~\\ref{burstregimes}.  At the lowest accretion rates \\citep[$\\lesssim 0.01\\dot{m}_{\\rm Edd}$, referred to as case~3 by][]{fhm81},\\footnote{See, e..g, \\citet{bil98a} for dependences of these critical accretion rates on the metallicities, and \\citet{ramesh03} for dependences on neutron star compactness.} the temperature in the burning layer is too low for stable hydrogen burning; the hydrogen ignites unstably, in turn triggering helium burning, which produces a type~I X-ray burst in a hydrogen-rich environment.  At higher accretion rates (case~2; $0.01 \\lesssim \\dot{m}_{\\rm Edd}\\lesssim 0.1$), hydrogen burns stably into helium between bursts, leading to a growing pure helium layer at the base of the accreted material. The fuel layer heats steadily until He ignition occurs and the He burns via the triple-$\\alpha$ process. At these temperatures and pressures, helium burning is extremely unstable, and a rapid and intense helium burst follows.  At yet higher accretion rates, hydrogen is accreted faster than it can be consumed by steady burning (limited by the rate of $\\beta$-decays in the CNO cycle), so that the helium ignites unstably in a H-rich environment (case 1; $0.1 \\lesssim\\dot{m}_{\\rm Edd} \\lesssim 1$). Finally, at the highest accretion rates, helium also begins to burn steadily between bursts. At just below $\\dot{m}_{\\rm Edd} (\\approx$0.9 in \\citealt{hcw07}; see also \\citealt{ramesh03}), an over-stability may arise that leads to oscillatory H and He burning, and in turn to intermittent bursts. Once the accretion rate exceeds $\\dot{m}_{\\rm Edd}$ (``case 0'' in \\citealt{fhm81}), stable helium burning depletes the fuel reserves and causes bursts to cease altogether.  Although the gross features of these regimes should be robust, several factors (some poorly-understood) could significantly change the accretion rates at which the bursting regimes occur. First, the burst behavior may be sensitive to certain individual thermonuclear reaction rates. % One reaction that has drawn particular attention is the ``break-out'' reaction \\OtoNe, which removes a catalyst from the CNO cycle \\citep[e.g.,][]{fis06}. A low rate for this reaction causes more H burning to occur, which produces a hotter burning layer in which steady helium burning also occurs \\cite[]{cn06}. As a result, unstable burning will cease altogether at $\\dot{m}_{\\rm Edd}$$\\ga$0.3, in contrast to predictions using higher rates \\cite[]{fhm81,bil98a,hcw07} but in partial agreement with observations. Recent experimental measurements however favor the original, higher \\OtoNe\\ rate \\citep{fis07}, and the reduction in the uncertainty means that this reaction cannot explain the cessation of bursts for most sources around $\\dot{m}_{\\rm Edd}$$\\ga$0.3.   Second, sedimentation or mixing in the burning layer could change the composition at the base of the burning layer. At low accretion rates ($\\dot{m}_{\\rm Edd} \\la 0.01$), CNO elements may settle to the bottom of the burning layer, which may prevent unstable H burning from inducing unstable He burning. Therefore, most case 3 bursts would be pure hydrogen. However, a large He layer would also build up that would eventually produce a very energetic burst \\citep{pbt07}, such as those seen from ``burst-only'' sources \\citep{corn04}. On the other hand, at high accretion rates ($\\dot{m}_{\\rm Edd} \\ga 0.1$), turbulent mixing of accreted fuel into deeper layers could increase the amount of steady burning in between bursts \\citep{pb07}.     Third, the accreted nuclear fuel may not be distributed evenly on the neutron star \\citep[e.g.,][]{is99,bil00,ps01}. If the material is deposited at the equator, a latitudinal gradient could develop in the amount of fuel burned steadily between bursts. Whether such inhomogeneous distribution of fuel will produce bursts that are confined to one hemisphere \\citep{bhatt06a}, or ignite slowly-propagating fires that burn fuel over limited regions of the neutron star \\citep[e.g.,][]{bil95} is uncertain. The variation in the effective gravity between the equator and higher latitudes could also lead to different ignition regimes, depending upon the spin rate of the neutron star \\cite[]{cn07b}. The spin rate may also affect the spreading via Coriolis forces, which may give rise to vortices that drift relative to the star as the burning spreads \\cite[]{slu02}.    Understanding these mechanisms is important, because current models for X-ray bursts have met with only partial success in explaining how their rates, energetics, and time scales vary with accretion rates. The basic predictions of Fujimoto et al. (1981; see also Cumming \\& Bildsten 2000) have found validation with the success of models in reproducing the energetics of case 3 mixed H/He bursts from EXO 0748--676 at an accretion rate of $\\dot{M} = 0.01\\dot{M}_{\\rm Edd}$ \\citep{boirin07a}, case 2 He bursts bursts from SAX J1808.4--3658 at an accretion rate of $\\dot{M} = 0.06\\dot{M}_{\\rm Edd}$ \\citep{gal06c}, and regularly-recurring case 1 mixed H/He bursts from GS 1826--24 at $\\dot{M} = 0.1\\dot{M}_{\\rm Edd}$ \\citep{gal03d}.     On the other hand, for several sources the burst rate {\\it decreases}\\/ as the accretion rate increases.  This decrease typically begins at $\\dot{M}_{\\rm Edd} \\sim 0.3$, well below the rate at which He-burning is expected to stabilize. These sources include 4U~1636$-$536 \\citep{lew87} and 4U~1705$-$44 \\citep{langmeier87}, as well as most of the sources in the sample assembled by \\cite{corn03a}. Furthermore, no correlation was found between persistent flux and burst recurrence times in Ser~X-1 \\citep{szt83} or 4U~1735$-$44 \\citep{lew80,vp88}. These observations may be evidence for ``delayed mixed bursts'' \\cite[between cases 1 and 0;][]{ramesh03}, in which helium begins to burn between bursts \\cite[see also][]{bil95,hcw07}. % A drop in burst frequency at comparable accretion rates has also been observed from 3A~1820$-$303 \\citep{clark77b}, which, with it's evolved donor, likely does not accrete any hydrogen. Alternatively, these observations may indicate that the accretion rate per unit area (which sets the burst ignition conditions) is decreasing even though the total accretion rate is increasing \\cite[e.g.][]{bil00}, or perhaps that the persistent fluxes are not a good measures of the accretion rates in these sources.    The change in the composition of the fuel layer as $\\dot{M}$ increases also affects the properties of the bursts. Helium burning occurs via the triple-$\\alpha$ process, which is moderated by the strong nuclear force and proceeds  very quickly at the temperatures and densities of the burning layer. Hydrogen burning proceeds more slowly, because it is limited by $\\beta$-decays moderated by the weak force. Therefore, faster, more intense bursts characteristic of a helium-rich burning layer should occur at relatively low accretion rates (case~2), while hydrogen-rich bursts with slower rise and decay times should occur at higher rates (case~1). Surprisingly, most sources behave in the opposite manner.  The decay  time scales of bursts has been observed to {\\it decrease} as the apparent $\\dot{M}$ increases from 0.01--0.1~\\mdotedd for 4U~1608$-$52 \\citep{murakami80b}, 4U~1636$-$536 \\citep{lew87}, 4U~1705$-$44 \\citep{langmeier87}, KS 1731$-$260 \\citep{muno00}, and in the sample of \\citet[][which includes several of the above LMXBs]{corn03a}. This discrepancy has been taken as evidence that steady helium burning is more prolific than expected at $\\sim$0.3$\\dot{M}_{\\rm Edd}$ \\citep[e.g.][]{ramesh03,cn06}. Some support for this hypothesis has been found in the appearance of low-frequency noise and mHz QPOs at these same accretion rates, that has been attributed to marginally unstable helium burning \\citep{bil95,rev01,ramesh03,hcw07}. %  \\begin{deluxetable}{llllc} \\tablewidth{0pt} \\tablecaption{Bursting Regimes \\label{burstregimes} } \\tablehead{ & & & \\colhead{Burst} & \\\\ \\colhead{Case} & \\colhead{$\\dot{m}_{\\rm Edd}$} & \\colhead{Steady Burning} & \\colhead{Composition} & \\colhead{Ref.} } \\startdata 3/V & $\\la$0.01 & none & mixed H/He & [1,2] \\\\ 2/IV & 0.01 -- 0.1 & stable H & pure He & [3,4] \\\\ 1/III & 0.1 -- 1.0 & stable H & mixed H/He & [4,5,6] \\\\ -/II & $\\sim$1.0 & over-stable H/He & mixed H/He? & [7,8] \\\\ 0/I  & $\\ga$1.0 & stable H/He & none & [9] \\\\ \\enddata \\tablecomments{The names for each of the burst cases are taken from \\cite[][Arabic numerals]{fhm81} and \\cite[][Roman numerals]{ramesh03}. The ranges in accretion rate ($\\dot{m}_{\\rm Edd}$; normalized to the Eddington rate) for each case are taken from \\citet{fhm81}; note that \\citet{ramesh03} predict lower rates for their cases II and I ($\\sim0.3\\dot{m}_{\\rm Edd}$). The references represent recent examples of calculations and/or comparisons to observations in each regime.} \\tablerefs{1. \\cite{pbt07}; 2. \\cite{cn07a}; 3. \\cite{gal06c}; 4. \\cite{woos03}; 5. \\cite{gal03d}; 6. \\cite{fis08}; 7. \\cite{ramesh03}; 8. \\cite{hcw07}; 9.  \\cite{schatz01}} \\end{deluxetable}    \\begin{figure} \\epsscale{1.2} \\plotone{f2.eps} \\caption{Spectral evolution in a thermonuclear burst exhibiting photospheric radius-expansion, from KS~1731$-$26. {\\it Top panel}\\/ Burst luminosity $L_X$, in units of $\\eps$; % {\\it middle panel}\\/ blackbody (color) temperature $kT_{\\rm bb}$; and {\\it bottom panel}\\/ blackbody radius $R_{\\rm bb}$. $L_X$ and $R_{\\rm Bbb}$ are calculated at an assumed distance of 7.2~kpc (Table \\ref{dist}). \t% Note the anticorrelation between $kT_{\\rm bb}$ and $R_{\\rm bb}$ in the first few seconds, indicative of the expanding photosphere, and the approximately constant flux throughout the expansion. The time at which the flux reaches a maximum is indicated by the open circle; by then the radius has declined to the asymptotic value in the burst tail, suggesting that the photosphere has settled (``touched down'') on the NS surface. \\label{pre_example} } \\end{figure}  \\subsection{Bursts as Standard Candles}  The peak flux for very bright bursts can reach % the Eddington luminosity at the surface of the NS, at which point the (outward) radiation pressure equals (or exceeds) the gravitational force binding the outer layers of accreted material to the star. Such bursts frequently exhibit a characteristic spectral evolution in the first few seconds, with a local peak in blackbody radius and at the same time a dip in color temperature, while the flux remains approximately constant (Fig. \\ref{pre_example}). This pattern is thought to result from expansion of the X-ray emitting photosphere once the burst flux reaches the Eddington luminosity; the effective temperature must decrease in order to maintain the luminosity at the Eddington limit, and excess burst flux is converted into kinetic and gravitational potential energy in the expanded atmosphere.  The largest uncertainty in the theoretical Eddington luminosity arises from possible variations in the photospheric composition. The limiting flux for a composition with hydrogen present at solar mass fraction will be a factor of 1.7 below that of a pure helium atmosphere. Nevertheless, the Eddington luminosities \\ledd\\ measured for LMXBs with independently-known distances are generally consistent to within the uncertainties, at a value estimated as $(3.0\\pm0.6)\\times10^{38}\\ {\\rm ergs\\,s^{-1}}$ by \\cite{lew93}, or, more recently, $(3.79\\pm0.15)\\times10^{38}\\ {\\rm ergs\\,s^{-1}}$ \\cite[cf. with equation \\ref{ledd};][]{kuul03a}.  This result is consistent with the narrow ranges for masses and surface redshifts expected for the neutron stars in these bursters. % Consequently, these photospheric radius-expansion (PRE) bursts can be used as distance indicators \\cite[]{bas84}. Time-resolved spectroscopy of radius-expansion bursts also allow in principle measurement of the surface gravitational redshift \\cite[e.g.][]{damen90,smale01}, although this has proved a considerable challenge \\cite[see e.g.][for a more recent study]{ozel06}.  \\subsection{A New Diagnostic of Nuclear Burning}  \\begin{figure} \\epsscale{1.2} \\plotone{f3.ps} \\caption{A dynamic power spectrum illustrating the typical frequency evolution of a burst oscillation. Contours of power as a function of frequency and time were generated from power spectra of 2 s intervals computed every 0.25 s. A Welch function was used to taper the data to reduce sidebands in the power spectrum due to its finite length \\citep{nr}. The $n$th contour level has a single-trial probability of $0.02^n$ of occurring randomly due to noise. The PCA count rate (histogram) is plotted referenced to the right axis. The time intervals defined in \\S\\ref{ibo} as the rise, peak, and decay are indicated with the vertical dashed lines. \\label{fig:freqev}} \\end{figure}  One of the key capabilities of \\xte\\/ is for high temporal resolution X-ray timing studies. Since 1996, this capability has led to the discovery of several distinct types of kHz variability in LMXBs \\cite[for a recent review, see ][]{gal08b}. Highly coherent burst oscillations with fractional amplitudes in the range 5--20\\% rms have been detected in thermonuclear bursts from \\noscsrc\\ sources to date (\\citealt{stroh96}; see also \\citealt{sb03}). As the burst evolves, these ``nuclear-powered pulsations'' typically increase in frequency by a few Hz, most approaching an asymptotic value which is stable for a given source from burst to burst \\cite[e.g. Fig. \\ref{fig:freqev};][]{muno02b}.  That the asymptotic frequency traces the NS spin has been confirmed by the detection of burst oscillations at the spin frequency in the millisecond pulsars SAX~J1808.4$-$3658 \\cite[]{chak03a} and XTE~J1814$-$338 \\cite[]{stroh03a}, as well as the prolonged oscillation detected during a superburst in 4U~1636$-$536 \\cite[]{stroh02b}.  Substantial questions remain regarding the mechanism of burst oscillations, as well as what conditions determine whether or not the oscillations will be detectable in a given burst, or a given source. The oscillations have been suggested to result from initially localized nuclear burning, which spreads over the surface of the neutron star during the early stages of the burst \\cite[]{stroh96}. However, this explanation does not account for the oscillations which persist as long as 5--10~s after the burst rise. The frequency drift is likely too large to be explained by angular momentum conservation in a decoupled expanded burning layer \\cite[]{cum02}, and may instead result from changes in the  velocity of a pattern in the surface brightness. % Slow-moving (in the rotating neutron star frame) anisotropies in the surface brightness may originate from hydrodynamic instabilities \\citep{slu02} or modes excited in the neutron star ocean (\\citealt{cb00}; see also \\citealt{heyl04,pb05} and references therein). Recently, the phenomenology has become even more complex with the observation of intermittent persistent pulsations which appear to be in some cases related to the occurence of bursts \\cite[]{gal07a,altamirano07} but in other cases not \\cite[]{casella07}.   \\subsection{The Need for a Global Study}  It is generally unfeasible to to study the variation in burst properties over a wide range of accretion rate using data from a single source. Typically only a narrow range of accretion rates is observed, and insufficient exposure time is available, leading to only a small number of bursts in total. A few previous observational studies have focussed on the properties of bursts from more than one source. A compilation of  45 measurements in the literature from ten LMXBs by \\cite{vppl88} revealed a global decrease in burst duration with increasing persistent flux, similar to that seen individually for several sources. They also found that $\\alpha$ was correlated with the (normalized) persistent flux, more strongly than was predicted by numerical models \\cite[e.g.][]{fhm81}. The normalized fluence depended principally on the burst interval $t_{\\rm rec}$, which suggested that continuous stable burning between bursts is a general phenomenon. % More recently, \\cite{corn03a} analysed six years of \\sax\\/ observations of 37 LMXBs, with a combined sample of 1823 bursts, and identified a transition between long, H-rich bursts (assumed to result from case~3 ignition) to short, pure He bursts (case~2), inferring the onset of steady H burning at a persistent luminosity of $2\\times10^{37}\\ \\eps$ (equivalent to $0.1\\dot{M}_{\\rm Edd}$), a factor of 10 higher than predicted by theory. Below this level, bursts were long, frequent and occurred quasi-periodically, typified by GS~1826$-$24 and KS~1731$-$26. Above $2\\times10^{37}\\ \\eps$ the burst rate dropped by a factor of five, and the bursts were short and occurred irregularly (although short bursts were also observed in the low accretion rate regime). At even higher luminosities, bursts ceased altogether in these sources (although are subsequently observed at $\\approx\\dot{M}_{\\rm Edd}$ in two sources, GX~17+2 and Cyg~X-2).   To date, the {\\it Rossi X-ray Timing Explorer}\\/ (\\xte) has observed 66\t% of the known thermonuclear burster sources, and discovered several new ones. The \\xte\\/ data are unparalleled for studies of bursts and bursters, thanks to the large instrumental effective area and high timing resolution. New data enter the public archive continually, and published analyses rarely take advantage of all the available bursts in all the public observations, let alone all the bursts from all the known bursters. To date, no global comparisons of theory with these data have been made. The wealth of observational data motivate the present work, which seeks to present a uniform analysis of all thermonuclear bursts from the bursters observed by \\xte\\/ through \\burstpubdate. By combining the bursts from different sources we achieve much larger burst numbers and a larger range of $\\dot{M}$ for global characterization of burst behaviour than is possible for any individual source. We also include information on the presence of burst oscillations, which is only available in the \\xte\\/ data. While \\xte\\/ has also observed several superbursts, we do not analyse these events in this paper.  We present the contents of the catalog, and the results of our studies, as follows. In \\S\\ref{analysis} we describe the analysis methods and products, and relate to the physical properties of the bursts. We summarise the properties of the catalog in \\S\\ref{results}, and present our detailed analysis in the subsequent sections. In \\S\\ref{pflux} we measure the mean peak flux of radius-expansion bursts, and determine the source distances. We analyse the properties of the individual bursts in \\S\\ref{ts}, and explore the consistency of the burst behaviour of different sources as a function of accretion rate in \\S\\ref{diversity}. In \\S\\ref{global} we combine the bursts from various sources in an attempt to quantify the global burst properties as a function of accretion rate and compare these properties to predictions from burst theory. We further discuss the global behaviour of the burst energetics in \\S\\ref{secalpha}. We attempt to place observational limits on the boundaries of the theoretical ignition regimes in \\S\\ref{boundary}. We discuss the properties of the millisecond oscillations in \\S\\ref{milosc}. Finally, we present a number of outstanding theoretical challenges in \\S\\ref{challenges}, and summarise our results in \\S\\ref{summary}. In appendix \\ref{sources} we present the results for individual bursters on a case-by-case basis; we constrain the origin for bursts in crowded fields in appendix \\ref{localize}. ",
        "conclusions": ""
    },
    "0608/astro-ph0608545_arXiv.txt": {
        "abstract": "% {} {The purpose of this study is to explore the relationship between galaxy stellar masses, based on multiwavelength photometry spectral template fitting and dynamical masses based on published velocity dispersion measurements, for a sample of 48 early-type galaxies at $z \\sim 1$ with \\textit{HST}/ACS morphological information.} {We determine photometric-stellar masses and perform a quantitative morphological analysis of cluster and field galaxies at redshift $0.6< z < 1.2$, using ground- and space-based multiwavelegth data available on the GOODS-S field and on the field around the X-ray luminous cluster RDCS1252.9-2927 at $z=1.24$. We use multi-band photometry over 0.4-8$\\mu$m from \\textit{HST}/ACS, \\textit{VLT}/ISAAC and \\textit{Spitzer}/IRAC to estimate photometric-stellar masses using Composite Stellar Population (CSP) templates computed with PEGASE.2 (Fioc \\& Rocca-Volmerange, 1997) models. We compare stellar masses with those obtained using CSPs built with Bruzual \\& Charlot (2003; BC03) and Maraston (2005; M05) models.  We then compare photometric-stellar mass and dynamical mass estimates as a function of morphological parameters obtained from \\textit{HST}/ACS imaging.} {Based on our sample, which spans the mass range $\\log M_{phot} \\simeq [10,11.5]$, we find that 1) PEGASE.2, BC03, M05 yield consistent photometric-stellar masses for early-type galaxies at $z \\sim 1$ with a small scatter (0.15 dex rms); 2) adopting a Kroupa IMF, photometric-stellar masses match dynamical mass estimates for early-type galaxies with an average offset of 0.27 dex; 3) assuming a costant IMF, increasing dark matter fraction with the increasing galaxy mass can explain the observed trend; 4) we observe that early-type galaxies with significant disk components (Sa/Sab) or with signs of dynamical interaction tend to have the largest deviation from a one-to-one $M_{dyn} vs M_{phot}$ relation.} {} ",
        "introduction": "Galaxy surveys have traditionally built samples based on the luminosities of galaxies, and have used luminosity functions to study the evolution of the galaxy population. The need to select galaxies on the basis of their (stellar or possibly total) mass and to trace the evolution of the space density of galaxies in mass bins has long been advocated. In such a way, one can reconstruct the build-up of galaxy mass across cosmic time and directly compare observations to predictions of currently favored hierarchical structure formation models. The evolution of the galaxy luminosity function is generally affected in a complex fashion by all the physical processes which modulate the star formation history of galaxies, whereas the mass function evolves in a smoother way and is a direct probe of galaxy formation models (e.g., \\citet{Baugh03}, \\citet{Hernquist03}, \\citet{Somerville04}, \\citet{Nagamine04}, \\citet{DeLucia05}). This has motivated many authors during the last decade to carry out surveys by selecting galaxies at near-infrared rest-frame wavelengths (e.g., \\citet{Songaila94}, \\citet{Cohen99}, \\citet{Drory01}, \\citet{Cimatti02}), since the rest-frame near-IR light is closely related to the total mass (e.g., \\citet{Gavazzi96}). \\\\More recently, the efficiency with which multi-wavelength photometry can be collected from ground-based and space-based observatories has enabled the measurement of stellar masses by fitting widely-sampled spectral energy distributions (SED) of galaxies with spectrum synthesis models (e.g., \\citet{Brinchmann00}, \\citet{Cole01}, \\citet{Papovich01}, \\citet{Shapley01}, \\citet{Dickinson03}, \\citet{Fontana03}, \\citet{Rudnick03}, \\citet{Fontana04}, \\citet{Drory04a}, \\citet{Saracco04}, \\citet{Rocca04}).  In parallel, considerable efforts have been devoted to measuring the evolution of the mass-to-light ratio of galaxies out to $z\\sim 1$, via Fundamental Plane studies (FP, \\citet{Djorgo87}, \\citet{Dressler87}) at different environmental densities (e.g.  \\citet{Franx93}, \\citet{vandokkum96}, \\citet{Busarello97}, \\citet{vandokkum98}, \\citet{Treu01}, \\citet{Gebhardt03}, \\citet{Treu05}, \\citet{vdw05a}, \\citet{dSA05}). These studies have provided increasing evidence that galaxy mass plays a critical role in driving galaxy evolution.  While it has become common to derive stellar masses from multi-wavelength surveys via SED fitting methods (the so called photometric-stellar masses), it remains important to test the reliability of this methodology which relies heavily on stellar population spectrum synthesis models.\\\\ Moreover, by comparing dynamical and photometric-stellar masses one can investigate the relevance of the dark matter component of early-type galaxies as a function of the total mass (e.g., \\citet{Napolitano05}).\\\\ \\citet{Drory04b} have performed such a comparison for a large sample of local galaxies in the SDSS \\citep{York00} spanning a large mass range, $8 < log(M_{phot}/M_{\\odot}) < 12$.  In this paper, we extend their study to higher redshifts, inevitably focusing on massive systems with $M_{phot} \\ge 10^{10} M_{\\odot}$. For these galaxies, \\citet{Drory04b} find a good match between the photometric-stellar mass, $M_{phot}$, and a quantity $M_{dyn }\\propto \\frac{\\sigma_{0}^{2}R_e}{G}$, although with a substantial scatter around the line of equality.\\\\ In a similar study, \\citet{Lintott05} have compared dynamical and photometric-stellar masses within the effective radii of a sample of local massive ($M_{phot} > 8 \\cdot 10^{10} M_{\\odot}$) early-type galaxies also selected from the SDSS. They find $M_{dyn} \\propto M_{phot}^{1.3}$, suggesting that dark matter might become increasingly important in massive galaxies. \\\\ \\citet{Cappellari05} have studied the correlation between the dynamical $M/L$ from virial and Schwarzschild modeling and the photometric-stellar mass (all of them within the effective radius) for a sample of early-type galaxies at $z \\sim 0$, with high $S/N$ integral field spectroscopy.  They use \\citet{VZ96,VZ99} stellar synthesis models and assume a \\citet{Kroupa01} IMF to compute photometric-stellar masses. They find the adoption of a Salpeter IMF to provide unphysical results, as a number of galaxies would have $M_{phot} > M_{dyn}$ and note that $\\sim 30 \\%$ dark matter within the effective radius and/or IMF variations are needed to explain the discrepancies between the two mass estimators.  This is consistent with earlier findings at $z \\sim 0$ from dynamical studies (e.g., \\citet{Gerhard01}, \\citet{Thomas05}) and at $z \\sim 1$ from gravitational lensing (e.g., \\citet{Treu04}, \\citet{Rusin05}).\\\\ Using total masses derived from gravitational lensing study of a sample of $0.3 < z < 1.0$ early-type galaxies, \\citet{Ferreras05} note a transition from little or no dark matter in the inner regions (within the effective radius) to dark matter dominating in the outer regions in the massive early-type galaxies, whereas no such a trend is observed in lower mass galaxies. \\\\ On a sample of 17 early-type galaxies at $z \\sim 1$ \\citet{dSA05} find a good agreement between the two mass estimators for the high mass galaxies, but most of their lower mass galaxies have stellar masses larger than the dynamical ones.\\\\ Here, we significantly extend their work {\\it i}) by using a sample almost a factor of 3 bigger, {\\it ii}) by measuring in a homogeneous fashion structural parameters for the entire galaxy sample using deep \\textit{HST}/ACS data, and {\\it iii}) by including rest-frame near-infrared photometry to estimate photometric-stellar masses.  With the availability of \\textit{Spitzer}/IRAC (InfraRed Array Camera) data \\citep{Fazio04}, one can access the critical restframe near IR-domain at $z\\simeq 1-2$, thus enabling a more accurate determination of photometric-stellar masses at redshifts where a significant fraction of the stellar mass is being assembled (e.g., \\citet{Dickinson03}).\\\\ This paper is organized as follows. In Sect. 2 we describe the sample selection; in Sect. 3 we present the dataset and describe the data analysis; in Sect. 4 we describe the morphological analysis based on \\textit{HST}/ACS images; in Sect. 5 we account for our method of deriving photometric-stellar masses for different spectrum synthesis models and discuss the results of their comparison; in Sect.6 we derive dynamical masses and discuss their comparison with photometric-stellar estimates in sect. 7. And finally in Sect. 8 we summarize the results.\\\\  We assume a $\\Omega_{\\Lambda} = 0.73$, $\\Omega_{m} = 0.27$ and $H_{0} = 71\\ \\rm{km} \\cdot \\rm{s}^{-1} \\cdot \\rm{Mpc}^{-1}$ flat universe \\citep{Spergel03}, and use magnitudes in the AB system throughout this work. \\begin{figure}[ht!] \\centering \\includegraphics[width=8cm]{5273F1.ps} \\caption{Top panel: The galaxy redshift distribution of the whole sample, consisting of 48 galaxies (black line), is compared with the distribution of each sub-sample over-plotted in red, blue and green for vdW05, dSA05 and Hol05, respectively.  Bottom panel: galaxy redshift distribution divided by morphological classes, 42 early-types (E/S0) (red) and 6 bulge-dominated spirals (Sa/Sab)(blue). Images of these 6 early-type spiral galaxies are shown in Fig. ~\\ref{spirals}.} \\label{zspec} \\end{figure} ",
        "conclusions": "In this work, we have used photometry from 9-10 passbands to build accurate SEDs, covering the 0.2-4$\\mu$m rest-frame wavelengths, of a sample of 48 early-type galaxies at $z\\sim\\! 1$ with published velocity dispersions.  The large wavelength baseline and accuracy of our photometric measurements allows us to compare measured stellar masses with different spectrum synthesis models.\\\\ Based on our sample, which spans a limited mass range $\\log M_{phot} \\simeq [10,11.5]$, corresponding to the bright end of the stellar mass function, $M \\gtrsim M\\ast$, we find that photometric-stellar mass estimates are not strongly dependent on the choice of the stellar population model. Regardless of the actual implementation of the TP-AGB phase in the different codes, we find the overall difference in photometric-stellar masses of early-type galaxies at $z \\sim 1$ from PEGASE.2, BC03, M05 not to be statistically significant (below 0.1 dex). \\\\  We have also investigated the other inherent systematic uncertainties on stellar masses, such as those due to reddening and adopted IMF, and found them of the order 0.2-0.3 dex.\\\\ We have then compared our photometric-stellar masses to the total dynamical masses as inferred from velocity dispersion measurements and half-light radii measured in deep \\textit{HST}/ACS images. Strong deviations from $M_{dyn}=M_{phot}$ may be ascribed to possible biases in dynamical mass measurements, as suggested by the evidence that deviations increase for early-type galaxies with small disk components and/or complex morphologies, as well as in galaxies showing signs of dynamical interaction with close-by companions. In the other hand, photometric-stellar masses depend on the assumed model of dust extinction and metallicity evolution as well as from the assumed IMF. \\\\ Assuming \\citet{Kroupa01} IMF, we find the photometric-stellar masses to reproduce the dynamical mass estimates with an an average offset of 0.27 dex. We note that the average offset depends on the assumed IMF, although a \\citet{Salpeter55} IMF would produce unphysical results by implying that a large number of galaxies would have photometric-stellar masses greater than dynamical estimates.  \\\\ We also find that an increasing dark matter fraction with the increasing total galaxy mass may be needed to explain the observed trend in $M_{dyn}(K_V) \\propto M_{phot}^{(1.25 \\pm 0.08)}$.  It is reassuring that a similar relation and slope was found in studies where the galaxy mass was directly derived from strong lensing models. The increase in $M_{dyn}/M_{phot}$ with increasing dynamical mass that we find is consistent with the increase of the mass-to-light ratio ($M_{dyn}/L$) of early type galaxies implied by FP studies. \\\\  We conclude that the determination of photometric stellar-masses of massive early-type galaxies at $z\\sim\\! 1$ is robust against stellar population models, when a large wavelength baseline is available, and its accuracy hinges primarily on the adopted IMF. \\\\ Although a clear correlation $M_{dyn}-M_{phot}$ is found, it remains difficult to use photometric masses as reliable surrogates of total galaxy masses, or mass-to-light ratios, over a wide mass range. Specifically, it would be interesting to extend our study to less massive galaxies ($ M_{phot} < 10^{10} M\\odot$) with possibly younger ages, not probed by our sample. Further studies on galaxy-scale lensing systems (e.g. , \\citep{Koopmans03}), with masses accurately determined by parameter free strong lensing models, will surely stimulate significant progress in the years to come."
    },
    "0608/astro-ph0608635_arXiv.txt": {
        "abstract": "We measure the cosmological matter density by observing the positions of baryon acoustic oscillations in the clustering of galaxies in the Sloan Digital Sky Survey (SDSS). We jointly analyse the main galaxies and LRGs in the SDSS DR5 sample, using over half a million galaxies in total. The oscillations are detected with 99.74\\% confidence (3.0$\\sigma$ assuming Gaussianity) compared to a smooth power spectrum. When combined with the observed scale of the peaks within the CMB, we find a best-fit value of $\\Omega_M=0.256^{+0.029}_{-0.024}$ (68\\% confidence interval), for a flat $\\Lambda$ cosmology when marginalising over the Hubble parameter and the baryon density. This value of the matter density is derived from the locations of the baryon oscillations in the galaxy power spectrum and in the CMB, and does not include any information from the overall shape of the power spectra. This is an extremely clean cosmological measurement as the physics of the baryon acoustic oscillation production is well understood, and the positions of the oscillations are expected to be independent of systematics such as galaxy bias. ",
        "introduction": "Baryonic Acoustic Oscillations (BAO) are predicted in the matter distribution with a calibration that depends on $\\Omega_M h^2$ \\citep{silk68,peebles70,sunyaev70,bond84,bond87,holtzman89}.  The oscillations arise because sound waves in the coupled baryon-photon plasma after an inflationary epoch will lead to the expansion of the baryonic material in a spherical shell around a small perturbation, reaching a radius $r_S(z_*)$, the comoving sound horizon size at recombination, before sound waves are no longer supported within the plasma \\citep{bashinsky01,bashinsky02}. At the high redshifts of interest the vacuum energy can be neglected, and $r_S(z_*)$ can be simply written \\citep{hu95} \\be r_S(z_*) / \\mpcoh \\equiv \\frac{1}{100 \\Omega_m^{1/2}} \\int_0^{a_*} \\frac{c_S}{(a + a_{\\rm eq})^{1/2}} \\, da. \\label{eq:rs} \\ee The expansion factor $a \\equiv (1+z)^{-1}$ and $a_*, a_{\\rm eq}$ are the values at recombination and matter-radiation equality respectively. Thus $r_S(z_*)$ depends on the matter density $\\Omega_M$ through the expansion rate and the recombination redshift. For a baryon density $\\Omega_bh^2 \\simeq 0.02$, we can approximate $c_S \\simeq 0.90 \\, c / \\sqrt{3}$.  Inserting $z_* = 1100$ and $a_{\\rm eq} = (23900 \\,\\Omega_Mh^2)^{-1}$ gives $r_S(z_*)=109\\mpcoh$ for $\\Omega_M=0.24$ and $h=0.73$.  In real space this leads to a peak in the correlation function at $r_S(z_*)$. In Fourier space, this process leads to oscillations in the power spectrum in the same way that the transform of a top-hat function yields a sinc function. The wavelength of these oscillations for $\\Omega_M=0.24$ and $h=0.73$ is $k_S=2\\pi/109=0.06\\hompc$. Numerical simulations have shown that a number of subtle corrections are required to this simple picture, although these corrections do not significantly affect the underlying important physics. On large scales, there is a phase shift in the position of the oscillations due to the contribution of the baryons to the drag epoch \\citep{eisenstein98}. We must also consider the damping of the oscillations on small scales at high redshift \\citep{silk68}, and, when modelling the oscillations at low redshifts, due to structure formation \\citep{eisenstein06}.  The resulting acoustic peaks in the cosmic microwave background (CMB) have now been observed with extreme precision \\citep{hinshaw06,spergel06}. However, interpreting the observed angular separation of these peaks in terms of the physics of the early universe requires knowledge of the angular diameter distance to the last scattering surface -- in the context of flat $\\Lambda$ cosmological models this leads to a parameter degeneracy between the matter density $\\Omega_M$ and the Hubble constant $h$, and models with the same value of $\\Omega_M^{0.275}h$ have the same projected acoustic horizon scale \\citep{percival02,page03,spergel06} (throughout this paper, a ``flat $\\Lambda$'' cosmological model implies a universe that is spatially flat with a time-independent ``dark energy'' component, i.e., a cosmological constant). Theoretically, we should expect the oscillations to survive in the galaxy power spectrum \\citep{meiksin99,springel05,seo05,white05,eisenstein06}, and the effects of baryons have been previously detected on large scales in the clustering of galaxies \\citep{percival01,miller01,cole05,eisenstein05,huetsi06}. These low redshift observations have a different dependence on the cosmological distance-redshift relation compared with the CMB because of the different angular and radial projection of these features. By comparing the two observations we can therefore probe the cosmological expansion history in addition to the physics of the BAO production. In this paper we use the BAO measured in the SDSS DR5 galaxy sample to set tight, clean constraints on the cosmological matter density. ",
        "conclusions": "In a separate paper, \\citet{tegmark06} carry out full multi-parameter fits to the LRG power spectrum and WMAP data over a broader CDM model space, drawing on both the BAO measurement and the broadband shape of the power spectrum.  The more focused analysis presented in this paper is complementary, obtaining constraints with minimal dependence on detailed cosmological assumptions.  Within the context of flat $\\Lambda$ models, the measurement of the matter density, $\\Omega_M = 0.256^{+0.029}_{-0.024}$, from the locations of the acoustic oscillations (the second combination above) is especially ``clean,'' relying on a single piece of well understood physics. In particular, we have decoupled constraints from the BAO with constraints from the overall shape of the power spectrum on the same scales. As demonstrated in \\citet{percival06}, the overall shape of the power spectrum, even on scales where the matter clustering has not deviated strongly from a linear CDM model is dependent on the galaxy population studied. It is therefore imperative to accurately model galaxy bias before robust cosmological constraints can be derived from such observations. By simply considering the BAO in this paper, we avoid this complexity.  Given the different physics probed, the consistency of the three $\\Omega_M$ estimates calculated with different additional information, itself provides support for the assumptions of a flat universe with a cosmological constant. With larger samples and a wider redshift range, the BAO ``standard ruler'' can be used to test these assumptions at high precision through the comoving distance-redshift relation \\citep{blake03,seo03}.  Forthcoming surveys have been designed exploit this effect.  It is clear that we are entering an era where BAO offer an extremely attractive route to strong, robust cosmological constraints."
    },
    "0608/astro-ph0608129_arXiv.txt": {
        "abstract": "We have produced a cleaned map of the \\emph{Wilkinson Microwave Anisotropy Probe} (WMAP) three-year data using an improved foreground subtraction technique. We perform an internal linear combination (ILC) to subtract the Galactic foreground emission from the temperature fluctuations observed by the WMAP. We divide the whole sky into hundreds of pixel groups with similar foreground spectral indices over a range of WMAP frequencies, apply the ILC for each group, and obtain a CMB map with foreground emission effectively reduced. With the resulting foreground-reduced ILC map (Figure \\ref{fig:SILC400}$b$, available on-line), we have investigated the known anomalies in CMB maps at large scales, namely the low quadrupole ($l=2$) power, and the strong alignment and planarity of the quadrupole and the octopole ($l=3$). Our estimates are consistent with the previous measurements. The quadrupole and the octopole powers measured from our ILC map are $\\delta T_2^2 = 276_{-126}^{+94}~\\mu\\textrm{K}^2$ and $\\delta T_3^2 = 952_{-83}^{+64}~\\mu\\textrm{K}^2$, respectively. The 68\\% confidence limits are estimated from the ILC simulations and include the cosmic variance. The measured quadrupole power is lower than the value expected in the concordance $\\Lambda$CDM model ($1250~\\mu\\textrm{K}^2$), in which the probability of finding a quadrupole power lower than the measured value is $5.7$\\%. We have confirmed that the quadrupole and the octopole are strongly aligned with angle $\\theta_{23} = 11\\fdg8_{-8\\fdg0}^{+6\\fdg4}$, and are planar with high planarity parameters $t=0.98_{-0.02}^{+0.02}$ for $l=2$ and $t=0.91_{-0.03}^{+0.02}$ for $l=3$. The observed angular separation $\\theta_{23}$ is marginally statistically significant because the probability of finding the angular separation as low as the observed value is 4.3\\%. However, the observed planarity is not statistically significant. The probability of observing such a planarity as high as the measured $t$ values is over 18\\%. The ILC simulations show that the residual foreground emission in the ILC map does not affect the estimated values significantly. The large scale modes ($l=2$--$8$) of SILC400 shows anti-correlation with the Galactic foreground emission on the southern hemisphere. It is not clear whether such anti-correlation occurs due to the residual Galactic emission or by chance. ",
        "introduction": "\\label{sec_introduction}  The \\emph{Wilkinson Microwave Anisotropy Probe} (WMAP; \\citealt{ben03:basic}) has measured the cosmic microwave background (CMB) temperature anisotropy and polarization with high resolution and sensitivity, and opened a new window to precision cosmology. The WMAP one-year data imply that the observed CMB fluctuations are consistent with predictions of the concordance $\\Lambda\\textrm{CDM}$ model with scale-invariant adiabatic fluctuations generated during the inflationary epoch \\citep{hin03,kog03,spe03,pag03,pei03}. The recent release of the WMAP three-year data has confirmed the primary results from the one-year data, giving more accurate determination of cosmological parameters \\citep{hin06,spe06,pag06}.  However, some peculiar aspects of the CMB maps have been noticed, leading to many controversies. First, there have been many reports of detection of non-Gaussian signatures (\\citealt{chi03,par04,eri04:asym,eri04:genus, col04,cop04,vie04,cru05,cru06,toj06}), as opposed to the WMAP team's result \\citep{kom03,spe06}. In particular, the origin of the asymmetry in statistical properties of the CMB between the Galactic northern and southern hemispheres (\\citealt{par04,eri04:asym}) has not been fully explained. \\citet{fre06} have studied the effects of the map-making algorithm on the observed asymmetry in CMB temperatures between the Galactic northern and southern hemispheres in the WMAP data. Another interesting issue is the low CMB quadrupole power. Cut-sky analysis of WMAP three-year data gives a quadrupole power of $211~\\mu\\textrm{K}^2$, which is quite a bit lower than the expected value in the best-fit $\\Lambda\\textrm{CDM}$ model ($1250~\\mu\\textrm{K}^2$; \\citealt{hin06,spe06}). Through subsequent analyses of the WMAP data, such a low quadrupole power has been confirmed and the effect of the residual foreground emission on the statistical properties of CMB has been extensively studied (e.g., \\citealt{teg03,efs04,eri04:ilc,bie04,bie05,slo04a,slo04b,nas06,deo06}). The alignment and planarity of $l=2$ and $3$ modes on the sky are also interesting features that have been seen in the WMAP data \\citep{deo04,sch04,lan05,cop06,abr06}.  Although each of the anomalies in the CMB anisotropy may have its own cosmological origin, there is a possibility that the residual Galactic foreground emission has affected the nature of the observed temperature fluctuations. Therefore, foreground subtraction is very important as a starting point of all CMB-related analyses. The most popular method of foreground removal is to model the Galactic emission as the weighted sum of foreground templates such as synchrotron, free-free, and dust emission maps (e.g., \\citealt{ben03:galaxy}). Another method of foreground removal is the internal linear combination (ILC) method using multi-frequency data \\citep{bra94,teg96,ben03:galaxy}.  The WMAP team has created the ILC map (hereafter WILC1YR) by computing a weighted combination of the WMAP maps that have been band averaged within each of the five WMAP frequency bands, all smoothed to $1\\degr$ resolution \\citep{ben03:galaxy}. The weights are determined by minimizing the variance of temperatures in the combined map, with a constraint that the sum of the weights is one. They have defined twelve disjoint regions on the sky, within which weights are determined independently. Except for the biggest Kp2 region, the other 11 regions are defined by subdividing the inner Galactic plane with lines of constant Galactic longitudes. However, the combination weights found by the WMAP team do not give the minimum variance, and their method of `non-linear searching' for the set of weights has not been described.  \\citet{eri04:ilc} have improved the WMAP team's ILC method by applying a Lagrange multiplier to produce a variance-minimized ILC map (hereafter LILC). They use the same 12 disjoint regions and the smoothing scheme to remove discontinuous boundaries as used in \\citet{ben03:galaxy}. \\citeauthor{eri04:ilc} has emphasized that the effects of noise on the performance of the ILC method is very important. If noise is high, the ILC method finds the best combination weights that minimize the instrumental noise rather than foregrounds.  \\citet{teg03} have applied a variant of the ILC method to make a foreground-cleaned CMB map (hereafter TCM) with high resolution. By noticing that the Galactic foregrounds appear dominantly on larger angular scales while the instrument noise dominates only at smaller scales, they have applied linear combinations in harmonic space to remove foreground emission at different angular scales separately. Based on levels of the Galactic foreground intensity, \\citet{teg03} have defined nine disjoint regions, where foreground removal has been done independently. Recently, \\citet{deo06} have made a new foreground-cleaned CMB map (hereafter TCM3YR) by applying the same technique to the WMAP three-year data.  In the three-year data analysis, the WMAP team has made a new ILC map (hereafter WILC3YR) by combining the WMAP data at five different bands \\citep{hin06}. From the one-year version of region definition map, they have eliminated the Taurus A region that is too small to give a reliable foreground removal while they have added a new region to minimize the dust residuals in the Galactic plane (denoted as Region 1 in Fig. \\ref{fig:WILC12}$a$). They have used the Lagrange multiplier method as used in \\citet{eri04:ilc} to determine the ILC coefficients. To quantify the bias due to the residual foreground emissions, they have performed one hundred Monte Carlo simulations using the one-year Galaxy model based on the Maximum Entropy Method (MEM). The three-year ILC map, WILC3YR, has been produced by subtracting this bias prediction.  All ILC methods require several disjoint regions that contain map pixels with similar properties. Otherwise, the efficiency of the ILC method becomes very low, leaving significant residual foregrounds. The Galactic foreground emission appears as the sum of the individual emissions with different physical natures, and its spectral behavior varies over a wide range of frequencies and over the whole sky. The disjoint regions defined by the WMAP team and \\citet{teg03} do not fully reflect the Galactic foreground properties.  In this paper, we derive a new foreground-reduced CMB map by applying the minimum-variance ILC method including information about the foreground spectral properties. The outline of this paper is as follows. In $\\S2$, we describe our simple minimum-variance ILC method, and present a new definition of sky regions for ILC. In $\\S3$, we derive a new foreground-cleaned CMB map from the WMAP three-year data. Measurements of statistics that are related to the quadrupole and the octopole are given in $\\S4$. We discuss our results in $\\S5$.  Throughout this paper, the power spectrum is calculated by \\begin{equation} \\delta T_l^2 = l(l+1) C_l / 2\\pi, \\label{eq:power} \\end{equation} with \\begin{equation} C_l = (2 l + 1)^{-1} \\sum_{l=-m}^{m} |a_{lm}|^2, \\label{eq:cl} \\end{equation} where $a_{lm}$ are the spherical harmonic coefficients in the Galactic coordinate system.  \\begin{figure*} \\mbox{\\epsfig{file=s1.eps,width=180mm,clip=}} \\caption{The three-year version of Galactic foreground emission maps at five WMAP frequencies (K, Ka, Q, V, and W bands) made by co-adding synchrotron, free-free, and dust emission components derived from the MEM. The resolution of each map is $1\\degr$ FWHM. The Mollweide projection in Galactic coordinates is used to display all maps. In each figure, the Galactic center is located at the center and the Galactic plane lies at the horizontal line that passes through the Galactic center. Galactic longitude increases from center to left ($l=0\\degr$--$180\\degr$) and from right to center ($l=180\\degr$--$360\\degr$) while Galactic latitude increases from bottom to top ($b=-90\\degr$--$+90\\degr$).} \\label{fig:MEMG} \\end{figure*} ",
        "conclusions": "\\label{sec_discuss}  In this paper, we have derived a new foreground-reduced CMB map by applying a simple internal linear combination method to the WMAP three-year data. Rather than using the disjoint sky regions as in WILC3YR, TCM, and LILC, we have defined a new group index map composed of four hundred pixel-groups that contain pixels with similar foreground spectral properties, and obtained a CMB map with foreground emission effectively reduced (SILC400).  Two hundred ILC simulations show that the residual foreground emission in the SILC400 is very small in amplitude, which is at $10$--$30~\\mu\\textrm{K}$ levels at high Galactic latitude regions (Fig. \\ref{fig:ResILC}$a$). On high latitude regions where the foreground emission is very weak at high frequency bands (V and W), the spectral index is inaccurately measured due to the MEM reconstruction noise in the Galactic foreground maps. If ILC regions are defined based on the inaccurate spectral index information, the performance of foreground removal with the ILC method becomes low. To reduce such noises, we have smoothed the foreground maps with a Gaussian filter whose width varies depending on the signal-to-noise level of intensities in the MEM Galaxy maps.  Such residual foreground features do not affect significantly the statistics of large-scale modes of CMB anisotropy. For example, when the average residual emission map (Fig. \\ref{fig:ResILC}$a$) is added to or subtracted from the SILC400 map, the estimated quadrupole power and angular separation are $\\delta T_2^2=288~\\mu\\textrm{K}^2$ and $\\theta_{23}=13\\degr$ for the addition map and $\\delta T_2^2=264~\\mu\\textrm{K}^2$ and $\\theta_{23}=11\\degr$ for the subtraction map. Both are within the $68\\%$ confidence limits of the SILC400 results. This implies that the residual foreground emission in our analysis map is not statistically important to the large-scale modes of CMB anisotropy.  We have also shown that the SILC400 recovers the true low $l$-mode powers, angular separation between the quadrupole and the octopole, and the planarity parameter (Figs. \\ref{fig:Plot_L2_SILC400}--\\ref{fig:Plot_powers}, \\ref{fig:Plot_pol23}, and \\ref{fig:planar}). Our SILC400 correctly reconstructs the spherical harmonic coefficients with minimal biases. Our ILC method is insensitive to the number of pixel-groups of common spectral index and on the smoothing scale as demonstrated in Figures \\ref{fig:stat_compare} and \\ref{fig:stat_compare_vs}.  The quadrupole and octopole powers measured from SILC400 are $244_{-203}^{+84}~\\mu\\textrm{K}^2$ and $859_{-109}^{+61}~\\mu\\textrm{K}^2$ (68\\% C.L.), respectively. According to the SILC400 simulations, the minimum variance ILC method tends to underestimate the low $l$-mode powers. Removing the effect of bias due to the residual Galactic emission on the spherical harmonic coefficients of the SILC400, we obtain $\\delta T_2^2 = 276_{-126}^{+94}~\\mu\\textrm{K}^2$ and $\\delta T_3^2 = 952_{-83}^{+64}~\\mu\\textrm{K}^2$. We confirm that the CMB quadrupole power is still lower than the theoretical value of the concordance $\\Lambda$CDM model, with $p = 5.7\\%$. Our estimate is consistent with that of WILC3YR, and is located between those measured from the TCM3YR and LILC.  The quadrupole power of the SILC400 is consistent with the previous measurements from the high latitude part of CMB data. \\citet{efs04} applied the maximum likelihood analysis method to the foreground-reduced CMB maps, and measured the quadrupole powers, $223~\\mu\\textrm{K}^2$ for WILC1YR and $250~\\mu\\textrm{K}^2$ for TCM for the Kp2 sky coverage. \\citet{bie04} applied the power equalization filter to the WMAP data with Kp2 mask, and obtained quadrupole powers $248\\pm45~\\mu\\textrm{K}^2$ for WILC1YR and $279\\pm49~\\mu\\textrm{K}^2$ for TCM.  The angular separation between the quadrupole and the octopole and the their planar properties have also been investigated. We have confirmed that the quadrupole and the octopole are aligned with high planarity. The probabilities of observing such anomalies from the bias-corrected SILC400 map are 4.3\\% for angular separation $\\theta_{23}$ and over 18\\% for planarity parameters (see Tables \\ref{tab:quo_dir} and \\ref{tab:planar}). The observed angular separation is marginally statistically significant.  As observed in Figure \\ref{fig:Plot_pscompare}, the power spectra of the CMB and the Galactic emission maps show zigzag patterns at $l=2$--$8$, with the zigzag directions opposite to each other, which strongly implies that the large scale CMB signal estimated from SILC400 and other analyses may be contaminated by the Galactic foreground emission and is anti-correlated with the Galactic signal (see also Tables \\ref{tab:Galfore} and \\ref{tab:alm_SILC400}). First, we have investigated whether the observed zigzag pattern at $l=2$--$8$ in the CMB power spectrum is statistically significant or not. Among the 200 true input CMB signal maps, we have found 16 cases with the CMB zigzag pattern, 10 cases with the Galaxy zigzag direction, and totally 26 cases (13\\%).  Secondly, we have measured cross-correlation at zero lag $C(0) \\equiv \\sum_{i} \\Delta T_{\\textrm{S}}(\\hat{\\mbox{\\boldmath $n$}}_i) T_{\\textrm{G}} (\\hat{\\mbox{\\boldmath $n$}}_i)/N_{p}$ between the SILC400 (bias-corrected) containing only $l=2$--$8$ modes and the K-band MEM Galactic foreground map smoothed by the $\\textrm{FWHM}=7\\degr$ Gaussian filter. Both maps have been degraded to $\\textrm{Nside}=32$ before cross-correlation measurement. The $C(0)$ values estimated from pixels at $|b| \\le 20\\degr$ are $-0.008$ and $-0.020~\\textrm{mK}^2$ for northern and southern hemispheres, respectively. We see that the CMB signal on the southern hemisphere has strong anti-correlation with the Galactic emission. We have calculated the cross-correlation functions between the 200 true input CMB signal maps and the Galactic emission map. The median values together with 68\\% confidence limits for $C(0)$ are $0.002_{-0.025}^{+0.024}$ and $0.000_{-0.023}^{+0.024}$ for $|b| \\le 20\\degr$ at the northern and southern hemispheres, respectively. The estimated probabilities that the large scale modes of true input CMB signal map have anti-correlation with the Galactic emission map as low as the measured values by chance are 36.5\\% and 19.5\\% for northern and southern hemispheres, respectively. The large scale modes of SILC400 on the southern hemisphere also shows anti-correlation with the Galactic emission even at $|b|>20\\degr$. However, whether such anti-correlation occurs due to the residual Galactic emission or by chance is not clear and further investigation is needed.  In this study, to measure the spectral indices of the foreground, we have used the MEM-derived Galactic foreground maps that contain some level of reconstruction errors and fail to model the Galaxy on the Galactic plane region. For the best reconstruction of CMB anisotropy through the ILC method, it is essential to use the Galactic emission maps that model the Galaxy realistically for precise measurement of the foreground spectral indices."
    },
    "0608/astro-ph0608403_arXiv.txt": {
        "abstract": "We derive a direct general map from the luminosity distance $D_{L}(z)$ to the inhomogeneous matter distribution $M(r)$ in the Lemaitre-Tolman-Bondi (LTB) cosmology and compute several examples. One of our examples explicitly demonstrates that it is possible to tune the LTB cosmological solution to approximately reproduce the luminosity distance curve of a flat FRW universe with a cosmological constant. We also discuss how smooth matter distributions can evolve into naked singularities due to shell crossing when the inhomogeneous ``curvature'' $E(r)$ is a function which changes sign. ",
        "introduction": "Standard Friedmann-Robertson-Walker (FRW) cosmology is characterized by the following features: homogeneity and isotropy when averaged on $\\mathcal{O}(100)$ Mpc length scales, negligible spatial curvature, and a stress tensor which today has a chemical composition of approximately 70\\% dark energy and 30\\% pressureless dust. The standard inflationary embedding of FRW cosmology currently provides a successful picture of the universe.  However, since the invocation of dark energy leads to a new coincidence problem, which is that dark energy dominance roughly coincides with the epoch of nonlinear structure formation, and since we still do not understand the cosmological constant problem, there has been a renewed interest in exploring alternative inhomogeneous cosmological models (which are not perturbatively related to FRW cosmologies) to see whether they might offer a competitively plausible picture of the universe \\cite{Alnes:2005rw,Celerier:1999hp,Vanderveld:2006rb,Chuang:2005yi,Moffat:2005yx,Moffat:2005ii,Nambu:2005zn,Ishibashi:2005sj,Apostolopoulos:2006eg,Kai:2006ws,Biswas:2006ub}.  Supernova luminosity distance measurements as a function of redshift offer compelling evidence for dark energy when interpreted in the context of FRW cosmologies \\cite{Perlmutter:1998np,Riess:1998cb}.  To interpret such data in the context of inhomogeneous cosmologies, it is often not particularly useful to average the underlying inhomogeneous variables to obtain a forced interpretation in terms of FRW cosmology.  This is because generically there is no preferred spatial slicing to compute averages and no meaningful map between the physical observables and the time derivatives of the spatially averaged variables. In some sense, the observables contain more information than that which could be carried by averaged variables, and hence determining the appropriate smearing map on the observables which matches the information that could be carried by averaged variables requires a knowledge of the underlying inhomogeneities in the absence of special symmetries. Therefore, to characterize inhomogeneous cosmologies, it is useful to compute observables such as the luminosity distance function directly in terms of variables describing the underlying inhomogeneous geometry.  The Lemaitre-Tolman-Bondi (LTB) solution corresponds to a spherically symmetric exact solution to the Einstein equations with pressureless ideal fluid. Extensive analyses have been carried out for this model because it allows for investigations of inhomogeneities that cannot be represented as perturbative deviations from FRW cosmologies. The LTB solution is fixed by choosing three smooth functions $\\{ E(r),M(r),R_{0}(r)\\}$, which allow an infinite number of different radial inhomogeneities. Most of the previous attempts to compute a cosmologically plausible luminosity distance $D_{L}(z)$ as a function of redshift $z$ consisted of finding a map from $\\{ E(r),M(r),R_{0}(r)\\}$ to $D_{L}(z)$ or a Taylor expansion of $D_L(z)$ about $z=0$ \\cite{Celerier:1999hp,Alnes:2005rw,Sugiura:1999dp,Mustapha:1997xb}.  In this paper, we derive a map from $\\{ E(r),D_{L}(z),R_{0}(r)\\}$ to $M(r)$. This has the advantage that the observed luminosity distance function $D_{L}(z)$ more directly dictates the underlying cosmological model, as opposed to having to guess the right $\\{ E(r),M(r),R_{0}(r)\\}$ to produce the desired $D_{L}(z)$. In \\cite{Vanderveld:2006rb}, a similar inverse problem (with a different choice of variables) was considered which focused on the situation with $E=0$, while in this paper, we will keep $E$ general. In using this new map, we find the interesting fact that the luminosity distance function is typically (depending on the choice of $E(r)$) an effective probe of the LTB geometry only for $z\\lesssim1$, since the luminosity distance function has a universal behavior in the limit that $\\{ R_{0}=0,E\\rightarrow0\\}$.  More precisely, $D_{L}(z)$ is fixed independently of $M(r)$ for $\\{ R_{0}=0,E=0\\}$, since that limit corresponds to the $\\Omega_M=1$ FRW universe.  A negative consequence of this feature is that the differential equation map from $\\{ E(r),D_{L}(z),R_{0}(r)\\}$ to $M(r)$ fails to be a numerically accurate map beyond $z \\sim 1$.  Nonetheless, we show how the numerical solution can be patched to a semi-analytic solution beyond $z=1$ to obtain a good fit (to within around $5\\%$ for the redshifts of interest) to obtain an LTB cosmology which reproduces the luminosity distance function of an FRW cosmology with $\\Omega_{\\Lambda}=0.7$ and $\\Omega_{M}=0.3$.  Furthermore, we make a subsidiary observation about the class of LTB solutions which can mimic the observed $D_{L}(z)$. We find that if the radial inhomogeneity profile goes from $E(r)>0$ to $E(r)<0$ as $r$ increases while $M'(r)$ is positive in that region, there is generically a danger of forming naked singularities, which can be interpreted as due to shell crossing.  The order of presentation will be as follows.  In the next section, we review the conventional approach to obtaining the luminosity distance as a function of redshift in the LTB cosmologies.  In Sec.~\\ref{sec:Inversion-method}, we construct a set of differential equations (which we will refer to as the inversion method) which can be used to map the luminosity distance function into a particular LTB geometry.  Afterwards, we apply the method to several examples. In Sec.~\\ref{sec:Naked-Singularity}, we discuss how smooth geometries can evolve into naked singularities when $E(r)$ switches signs.  We then summarize and conclude.  For completeness, we present the FRW luminosity distance with $\\Omega_M+\\Omega_\\Lambda=1$ in Appendix A. In Appendix B, we write the LTB solution explicitly in a form that is not commonly found in the literature. ",
        "conclusions": "In this paper, we have derived a set of differential equations in the context of LTB cosmologies (spherically symmetric pressureless dust solutions to the Einstein equations) which can be numerically solved to obtain almost any luminosity-distance-redshift relationship that can be produced by a homogeneous and isotropic FRW model.  We have solved this set of differential equations for several examples to demonstrate the feasibility of our inversion method.  Unlike many other methods in the literature, our method can be used to dial in the geometry that generates the desired luminosity distance {\\em exactly} in a finite redshift interval.  We have also given explicit examples of naked singularity formation in LTB cosmologies in the region where $E(r)$ changes sign.  Although our work serves as a step towards building cosmological models competitive to the standard FRW cosmology, the research program is far from completion. Toy models such as the LTB cosmologies are arguably not yet convincing contenders for compelling alternatives to standard inflationary cosmology, since being at the center of the universe requires giving up the Copernican principle (though recently there has been some effort to alleviate this problem \\cite{Kai:2006ws}) and there is not yet a convincing structure formation history that could explain such radial inhomogeneities. Nonetheless, it is not obvious whether this class of LTB models can be ruled out from current observations \\cite{Alnes:2005rw}.  Furthermore, there is some evidence that our galaxy is in a void, qualitatively similar to the voids presented in two of our numerical examples. However, the existence of the void is still currently being investigated (see for example \\cite{Frith:2005et}), and whether the magnitude of the void is sufficient to cause the redshift effects presented in this paper is unclear.  Regarding tests of LTB cosmologies, any physical probe testing the isotropy of the universe from a point separated sufficiently far away from us could in principle be useful.  We leave this problem for future investigation."
    },
    "0608/hep-th0608202_arXiv.txt": {
        "abstract": "We address a detailed non-perturbative numerical study of the scalar theory on the fuzzy sphere. We use a novel algorithm which strongly reduces the correlation problems in the matrix update process, and allows the investigation of different regimes of the model in a precise and reliable way. We study the modes associated to different momenta and the r\\^ole they play in the ``striped phase'', pointing out a consistent interpretation which is corroborated by our data, and which sheds further light on the results obtained in some previous works. Next, we test a quantitative, non-trivial theoretical prediction for this model, which has been formulated in the literature: The existence of an eigenvalue sector characterised by a precise probability density, and the emergence of the phase transition associated with the opening of a gap around the origin in the eigenvalue distribution. The theoretical predictions are confirmed by our numerical results. Finally, we propose a possible method to detect numerically the non-commutative anomaly predicted in a one-loop perturbative analysis of the model, which is expected to induce a distortion of the dispersion relation on the fuzzy sphere. ",
        "introduction": "The study of quantum field theory in non-commutative spaces has attracted considerable attention over the last years~\\cite{Doplicher:1994tu, Landi:1997sh, Douglas:2001ba, Szabo:2006wx}. This research area has a long history, since the possibility of a quantised structure of spacetime at short distances was first mentioned as early as in the 1930's in some correspondence among Heisenberg, Peierls, Pauli and Oppenheimer~\\cite{Pauli1,Pauli2}, and in the papers published by Snyder~\\cite{Snyder:1946qz}, by Yang~\\cite{Yang:1947ud} and by Moyal~\\cite{Moyal:1949sk} during the 1940's. Although the original motivation to use non-commutativity as a tool to regularise QFT was soon frustrated --- while the renormalisation approach proved to be a  successful method to handle the divergences encountered in the formulation in commutative spacetime ---, non-commutative spaces have attracted renewed interest in more recent years, with the application of this formalism to solid-state physics, to the problem of the quantum Hall effect~\\cite{Karabali:2004xq}, to the study of MHD waves in astrophysics~\\cite{Bourouaine:2006mp}, and with the discovery of the relevance of such spaces to string theory~\\cite{Connes:1997cr, Douglas:1997fm, Alekseev:1999bs, Seiberg:1999vs, Kar:1999tu, Kar:2000pb, Lizzi:2006te} and to a possible quantum theory of gravity~\\cite{Chamseddine:1992yx, Lizzi:2002ib}.  Groenewold-Moyal $\\R^n_\\theta$ spaces are among the most extensively studied non-commutative spaces. The properties of QFT defined in these spaces --- including those related to renormalisability, causality, non-locality, Poincar\\'e invariance \\emph{et c\\oe tera} --- have been addressed in several works~\\cite{Doplicher:1994zv, Doplicher:2001qt, Filk:1996dm, Chen:2001qg, Jain:2003xs, Chaichian:2004za, Szabo:2001kg, Minwalla:1999px, Hinchliffe:2002km, Gubser:2000cd, Castorina:2003zv, Langmann:2002cc, Steinacker:2005wj, Grosse:2004yu, Rivasseau:2005bh, Govindarajan:2006vh, Grosse:2005ig, Lizzi:2006xi, Stern:2006zt, Grosse:2006qv, Gracia-Bondia:2006yj, Grosse:2006tc, Disertori:2006uy, Grosse:2007dm}; a lattice-like regularised formulation has also allowed to investigate numerically various aspects of these models~\\cite{Ambjorn:1999ts, Ambjorn:2000nb, Ambjorn:2000cs, Azuma:2004zq, Bietenholz:2002ch, Bietenholz:2002vj, Ambjorn:2002nj, Bietenholz:2004xs, Bietenholz:2005iz, Bietenholz:2006cz}. In particular, one of the most interesting --- albeit troublesome --- features is the fact that the effective action describing QFT in a Groenewold-Moyal space is divergent when the external momentum along the non-commutative directions vanishes: this effect arises from the integration of the high-energy modes in non-planar loop diagrams, and is henceforth called ``ultra-violet/infra-red (UV/IR) mixing''.  Another class of non-commutative spaces is given by fuzzy spaces: they are built approximating the infinite-dimensional algebra of functions on some particular manifold by means of a finite-dimensional algebra of matrices. Under some conditions, this construction is possible for even-dimensional co-adjoint orbits of Lie groups which are symplectic manifolds --- see~\\cite{Madore:1991bw, Alexanian:2001qj, Balachandran:2001dd, Hammou:2001cc, Medina:2002pc, Grosse:1995pr, GrosseReiter, Balachandran:2002jf, Vaidya:2003ew, Dolan:2003kq, Lizzi:2003ru, Lizzi:2005zx, Balachandran:2005ew, Sheikh-Jabbari:2006bj} and references therein. In particular, co-adjoint orbits of semi-simple Lie groups are adjoint orbits; examples include the $\\C P^n$ complex projective spaces. The most-widely known example of a fuzzy space is the fuzzy two-sphere $S^2_F$~\\cite{Madore:1991bw}, built truncating the algebra of functions on the commutative sphere $S^2$ to a maximum angular momentum $l_{\\mbox{\\tiny{max}}}$. The fuzzy sphere depends on two parameters: the matrix size $N=l_{\\mbox{\\tiny{max}}} +1$ and the radius $R$; the commutative sphere and the non-commutative plane can be obtained in different limits of $N$ and $R$.  A one-loop perturbative calculation shows that, for every finite $N$, QFT on the fuzzy sphere is finite; furthermore, the theory is not affected by the UV/IR mixing problem~\\cite{Chu:2001xi} (although --- due to a non-commutative anomaly --- the latter re-emerges once a double-scaling limit is taken, in which the fuzzy sphere goes over to the non-commutative plane $\\R^2_\\theta$). This feature, as well as the fact that the fuzzy approach explicitly preserves the symmetries of the original manifold for any value of $N$ and allows a well-defined treatment of the topological properties~\\cite{Baez:1998he, Grosse:1995jt, Balachandran:1999hx, Presnajder:1999ky, Carow-Watamura:1998jn, Steinacker:2003sd, Behr:2005wp, Nair:2006qg, Valtancoli:2006dw, Dolan:2006tx, Murray:2006pi, Saemann:2006gf, Aoki:2006wv,  Delgadillo-Blando:2006dp,  Steinacker:2007iq}, has led to suggest the fuzzy space as a potentially interesting candidate for regularisation of quantum field theory.  As a matter of fact, QFT on the fuzzy sphere is mathematically well-defined and finite~\\cite{Grosse:1995ar}, and the formulation is amenable to a non-perturbative approach and to numerical studies using Monte Carlo simulations, with fields represented as finite-dimensional matrices. This approach has been followed in various recent works~\\cite{Martin:2004un, GarciaFlores:2005xc, Medina:2005su, Azuma:2004yg, Anagnostopoulos:2005cy, Azuma:2005pm, O'Connor:2006wv}.  In the present paper we address a detailed Monte Carlo study of the $\\Phi^4$ scalar field theory on the fuzzy sphere; among other issues of interest, this simple model provides a laboratory to test the possibility to use the fuzzy space approach as a potential regularisation scheme for more realistic field theories. As it concerns the practical implementation of Monte Carlo simulations of the model, we shall present a novel algorithm, which reduces the autocorrelation time, combining overrelaxation steps with ergodic configuration updates.  Preliminary results of this study have been presented in~\\cite{Panero:2006cs}.  This manuscript has the following structure: in section~\\ref{theorysect} the theoretical framework underlying the model is recalled, and the basic notations are introduced. In section~\\ref{resultsect} we discuss the implementation of the numerical simulations of the model, and present the results obtained with our algorithm. In section~\\ref{discussionsect} we comment on the implications of these results and on possible research perspectives. A technical discussion of the algorithm is presented in the Appendix~\\ref{overrelsect}. ",
        "conclusions": ""
    },
    "0608/astro-ph0608365_arXiv.txt": {
        "abstract": "We probe the scale dependence of the primordial spectrum in the light of the three-year WMAP (WMAP3) alone and WMAP3 in combination with the other cosmological observations such as  galaxy clustering and Type Ia Supernova (SNIa). We pay particular attention to the combination with the Lyman $\\alpha$ (Ly$\\alpha$) forest. Different from the first-year WMAP (WMAP1), WMAP3's preference on the running of the scalar spectral index on the large scales is now fairly independent of the low CMB multipoles $\\ell$. A combination with the galaxy power spectrum from the Sloan Digital Sky Survey (SDSS) prefers a negative running to larger than 2$\\sigma$, regardless the presence of low $\\ell$ CMB ($2\\le \\ell \\le 23$) or not. On the other hand if we focus on the Power Law $\\Lambda$CDM cosmology with only six parameters (matter density $\\Omega_m h^2$, baryon density $\\Omega_b h^2$, Hubble Constant $H_0$, optical depth  $\\tau$, the spectral index, $n_s$, and the amplitude, $A_s$, of the scalar perturbation spectrum) when we drop the low $\\ell$ CMB contributions WMAP3 is consistent with the Harrison-Zel'dovich-Peebles scale-invariant spectrum ($n_s=1$ and no tensor contributions) at $\\sim 1\\sigma$. When assuming a simple power law primordial spectral index or a constant running, in case one drops the low $\\ell$ contributions ($2\\le \\ell \\le 23$) WMAP3 is consistent with the other observations better, such as the inferred value of $\\sigma_8$. We also find, using a spectral shape with a minimal extension of the running spectral index model, LUQAS$+$ CROFT Ly$\\alpha$  and SDSS Ly$\\alpha$ exhibit somewhat different preference on the spectral shape. ",
        "introduction": "It is commonly believed that the early Universe has experienced a stage of accelerated expansion which is known as inflation \\cite{Sato:1980yn,Guth:1980zm,Linde:2005ht}. During inflation one of the tiny patches of the very early Universe was superluminally stretched to become our observable Universe today. The theory of inflation can naturally explain why the universe is flat, homogeneous and isotropic. Inflation is driven by the potential energy of a scalar field called inflaton and its quantum fluctuations turn out to be the primordial density fluctuations which seed the observed large-scale structures (LSS) and anisotropy of the Cosmic Microwave Background (CMB) Radiation. Inflation typically ends with a period of reheating when the inflaton starts rapid oscillations and it decays into standard particles. Then the universe starts the standard hot Big Bang evolution.   Over the past decade, the theory of inflation has successfully passed several nontrivial tests. In particular, the first year Wilkinson Microwave Anisotropy Probe(WMAP) \\cite{Bennett:2003bz,Kogut:2003et,Peiris:2003ff,Spergel:2003cb,Hinshaw:2003ex,Verde:2003ey} has detected a large-angle anti-correlation in the temperature-polarization cross-power spectrum, which is the signature of adiabatic superhorizon fluctuations at the time of decoupling \\cite{Peiris:2003ff}.  The recently released three year WMAP (WMAP3) \\cite{Spergel:2006hy,Page:2006hz,Hinshaw:2006,Jarosik:2006,WMAP3IE} has marked another milestone on the precision cosmology of CMB Radiation. The simplest six-parameter power-law $\\Lambda$CDM cosmology, namely with matter density $\\Omega_m h^2$, baryon density $\\Omega_b h^2$, Hubble Constant $H_0$, optical depth  $\\tau$, the spectral index and amplitude of the scalar perturbation spectrum, $n_s$ and $A_s$, turn out to be in good agreement with WMAP3 together with the small-scale CMB observations such as BOOMERANG  \\cite{Montroy:2005yx}, the Arcminute Cosmology Bolometer Array Receiver (ACBAR  \\cite{Kuo:2002ua}), the Cosmic Background Imager (CBI  \\cite{Readhead:2004gy}) and the Very Small Array(VSA  \\cite{Dickinson:2004yr}), LSS as measured by Two degree Field (2dF)  \\cite{Cole:2005sx} and the Sloan Digital Sky Survey (SDSS)  \\cite{Tegmark:2003uf} and with the Type Ia Supernova (SNIa) as measured by the Riess ``gold'' sample  \\cite{Riess:2004nr} and the first year Supernova Legacy Survey (SNLS)  \\cite{snls}. This agreement between the above ``canonical'' cosmological model and observations can be used to test a number of possible new physics, such as the equation of state of dark energy, neutrino masses, cosmic CPT violation, etc.   Although the temperature-temperature correlation power (TT) of WMAP3 is now cosmic-variance limited up to $\\ell \\sim 400$  and the third peak is now detected, the features discovered by the first year WMAP \\cite{Peiris:2003ff,Spergel:2003cb} are still present: the low TT quadrupole and localized oscillating features on TT for $\\ell \\sim 20-40$ \\cite{Hinshaw:2006}. Although the signatures of glitches on the first peak discovered by the first year WMAP have now become weaker, they do exceed the limit of cosmic variance \\cite{Hinshaw:2006}.  In fact, even before the release of the first year WMAP Ref.  \\cite{Wang:2002hf} claimed oscillating primordial spectrum could lead to oscillations around the first peak of CMB TT power. These data  have hence revealed many interesting detailed features which cannot be explained in the simplest version of the inflation model, and we are making much effort to construct a realistic particle physics model of inflation.   As for the shape of the primordial power spectrum, it is also noteworthy that a significant deviation has been observed by WMAP3 from the simplest Harrison-Zel'dovich-Peebles scale-invariant spectrum ($n_s=1$ and no tensor contributions), and that this feature is more eminent with the combination of all the currently available CMB, LSS and SNIa (dubbed the case of ``All'' in \\cite{WMAP3IE}). On the other hand, a nontrivial negative running of the scalar spectral index $\\alpha_s$, whose existence was studied even before WMAP epoch \\cite{Feng:2003nt}(for relevant study see also  \\cite{Lewis:2002ah}), was favored by the first-year WMAP papers  \\cite{Bennett:2003bz,Peiris:2003ff,Spergel:2003cb}. But its preference was somehow diminished as corrections to the likelihood functions were made \\cite{Cline:2003ve,Efstathiou:2003tv,Slosar:2004fr}. Dramatically, the new WMAP3 data prefers again a negative running in the ``All'' combination  \\cite{WMAP3IE}. If  confirmed, a nonvanishing  running of $\\alpha_s$ would not only constrain inflationary cosmology significantly \\cite{Feng:2003mk,Kawasaki:2003zv,Huang:2003zp,Yamaguchi:2003fp,Yamaguchi:2004tn,Chen:2004nx,Easther:2006tv,Yokoyama:1998rw,Yokoyama:1998pt}, but also affect the cosmological constraint on the neutrino mass \\cite{Feng:2006zj,toappear}.   In the present paper, using the Markov Chain Monte Carlo (MCMC) method, we aim to probe the scale dependence of the primordial spectrum in the light of WMAP3 alone and that in combination with the other cosmological observations  from galaxy clustering and SNIa. We pay particular attention to the combination with the Lyman $\\alpha$ (Ly$\\alpha$) forest.  We will show that, different from the first-year WMAP (WMAP1), WMAP3's preference on the running of the scalar spectral index on the large scales is now fairly independent of the low CMB multipoles $\\ell$. We  also find that a combination of WMAP3 with the galaxy power spectrum from SDSS prefers a negative running with more than 2$\\sigma$, regardless of the presence of low $\\ell$ CMB ($2\\le \\ell \\le 23$).  On the other hand, if we focus on the Power Law $\\Lambda$CDM cosmology with only six parameters, we find that when one drops the low $\\ell$ CMB contributions WMAP3 would be consistent with the Harrison-Zel'dovich-Peebles scale-invariant spectrum at $\\sim 1\\sigma$. Assuming a simple power-law primordial spectral index or a constant running, in cases one drops the low $\\ell$  contributions ($2\\le \\ell \\le 23$) WMAP3 is consistent with the other observations better, such as the value of $\\sigma_8$. %  On the other hand, using a spectral shape with a minimal extension of the running spectral index model, we find LUQAS$+$ CROFT Ly$\\alpha$ \\cite{Viel:2004bf} and SDSS Ly$\\alpha$ \\cite{McDonald:2004xn} exhibit somewhat different preference on the spectral shape, although it is difficult to draw any definite conclusion because the result depends on how we model the spectral shape.   The rest of the paper is structured as follows. In Section II we describe the method and the data. In Section III we analyze the effects of low CMB multipoles on the determination of cosmological parameters using WMAP3 \\cite{Spergel:2006hy,Page:2006hz,Hinshaw:2006,Jarosik:2006,WMAP3IE}, SNIa  \\cite{Riess:2004nr,snls}, 2dF \\cite{Cole:2005sx} and Sloan Digital Sky Survey 3-D power spectrum (SDSS P(k)) \\cite{Tegmark:2003uf}, LUQAS+CROFT sample of Lyman $\\alpha$ \\cite{Viel:2004bf} and SDSS Lyman $\\alpha$ sample \\cite{McDonald:2004xn}  by global fittings using the MCMC technique. In Section IV scale dependence of fluctuations is analyzed with priors on the Hubble parameter, $\\Omega_bh^2$, and the cosmic age.  Discussions and conclusions are presented in the last section. ",
        "conclusions": "Through our detailed investigations we have found that values of cosmological parameters obtained by MCMC change considerably depending on whether we use low multipole components of CMB or not. Currently, the WMAP team use a very large number of pixels (666+1172 pixels for low multipole temperature and polarization spectra \\cite{Spergel:2006hy,Page:2006hz,Hinshaw:2006,WMAP3IE}), which may take too large a weight for the determination of the likelihood and hence cosmological parameters. This may be somewhat problematic because these low multipoles suffer from relatively large cosmic variance. On the other hand, if we incorporate low $\\ell$ polarization components, the optical depth $\\tau$ can be determined with a smaller ambiguity because the degeneracy between $n_s$ and $\\tau$ is partially removed, which in turn result in a better determination of the spectral index. Hence it is difficult to conclude which approach gives more precise results.  We obtain different values of cosmological parameters with different datasets.  For example, we find a very low $H_0$ in the combination of WMAP3 (without $\\ell \\le 23$) and SDSS is rather interesting, even though we have adopted the HST prior in Table 6.  Nonetheless these discrepancies are typically within 2$\\sigma$ and do not imply the datasets we have used are incompatible with each other, including the differences between the Riess ``gold'' sample and the first year SNLS, SDSS and 2dF, or LUQAS+CROFT sample of Lyman $\\alpha$ as analyzed in Ref. \\cite{Viel:2004bf} and the SDSS Lyman $\\alpha$ sample from Ref. \\cite{McDonald:2004xn}  While our results as a whole support the concordance cosmology, they also indicate we are still far from being able to determine all of the six cosmological parameters with two digits' accuracy in the context of this simple power-law $\\Lambda$CDM model. Hence we should be open-minded to new physics which may be hidden in the error bars of the current observational data and continue blind analysis without any theoretical prejudices.  A good example is the possible running of the spectral index, which is preferred to more than 2$\\sigma$ in a number of combinations of datasets.    {\\bf{Acknowledgments:}} We acknowledge the use of the Legacy Archive for Microwave Background Data Analysis (LAMBDA). Support for LAMBDA is provided by the NASA Office of Space Science. We have performed our numerical analysis on the Shanghai Supercomputer Center (SSC). We used a modified version of CAMB \\cite{Lewis:1999bs,IEcamb} which is based on CMBFAST \\cite{cmbfast,IEcmbfast}. We thank Patrick Mcdonald and Anze Slosar for kind discussions regarding their online codes for SDSS Lyman alpha fittings \\cite{IESlya1,IESlya2}. We are grateful to Kohji Yoshikawa and Ting Zhang for computational assistance. We thank Kevork Abazajian, Long-long Feng, Julien Lesgourgues, Antony Lewis, Hiranya Peiris, David Spergel, Yasushi Suto, Masahiro Takada, Max Tegmark, Matteo Viel, Yongzhong Xu, Pengjie Zhang, Xinmin Zhang and Gongbo-Zhao for helpful discussions. The work of J.Y. is supported partially by the JSPS Grant-in-Aid for Scientific Research No. 16340076 and B.F. is by the JSPS fellowship program. This work is supported in part by National Natural Science Foundation of China under Grant Nos. 90303004, 10533010 and 19925523 and by Ministry of Science and Technology of China under Grant No. NKBRSF G19990754.   {\\it Note Added: } During the finalizing of our manuscript, Refs. \\cite{Eriksen:2006xr,Huffenberger:2006xx,Bridges:2006zm,Finelli:2006fi} appeared on the arXivs which are somewhat relevant to our study."
    },
    "0608/astro-ph0608686_arXiv.txt": {
        "abstract": "Recent measurements are suggesting that we live in a flat Universe and that its present accelerating stage is driven by a dark energy component whose equation of state may evolve in time. Assuming two different parameterizations for the function $\\omega(z)$, we constrain their free parameters from a joint analysis involving measurements from X-Ray luminosity of galaxy clusters and SNe type Ia data. ",
        "introduction": "%  \\hspace{0.5cm}In the framework of general relativity, the present accelerating stage of the Universe (as indicated by SNe type Ia observations) can be explained by assuming the existence of a substantial amount of an exotic dark energy component with negative pressure, also known as quintessence\\cite{Lima04,Ries04}. The existence of this extra component filling the Universe (beyond the cold dark matter) has also been indirectly suggested by independent studies based on fluctuations of the 3K relic radiation, large scale structure, age estimates of globular clusters or old high redshift objects, as well as by the X-ray data from galaxy clusters\\cite{data,Allen02}.  A cosmological constant ($\\Lambda$), the oldest and by far the most natural candidate for dark energy, faces some theoretical difficulties. The most puzzling of them is the so-called cosmological constant problem: the present cosmological upper bound, $\\Lambda_o/8\\pi G \\sim 10^{-47} GeV^{4}$, differs from natural theoretical expectations from quantum field theory, $\\sim 10^{71} GeV^{4}$, by more than 100 orders of magnitude. Actually, such a problem has also inspired many scenarios driven by a $\\Lambda(t)$ or a time varying decaying vacuum with constant equation of state\\cite{LT96}. Among the remaining candidates to dark energy, the most promising ones lead to a time dependent equation of state (EOS), usually associated to a dynamical scalar field component. Such quintessence models may also parametrically be represented by an equation of state, $\\omega(z)$, as proposed by Cooray and Huterer\\cite{CoHu99}, as well as the one discussed by Linder\\cite{Lin03}, and, independently, by Padmanabhan and Choudhury\\cite{PaCh03}. In principle, the time variation of the EOS parameter, $\\omega(z) \\equiv p/\\rho$, may allow a clear distinction between a cosmological constant model and the one driven by a rolling scalar field.  In actual fact, the exploration of the expansion history of the universe using $\\omega(z)$ gave origin to a heated debate with growing interest in the recent literature. The SNe type Ia test is the most promising one related to this subject. However, Maor et al.\\cite{MaBrSt01}, and Weller and Albrecht\\cite{WeAl01}, have also observed  that in order to constrain the evolution of the EOS with SNe observations, it is necessary to use a tight prior on the mean matter density of the Universe. A natural way to circumvent such a problem is to consider  the constraints on the density parameter from measurements of the X-Ray luminosity of galaxy clusters together in a joint analysis involving SNe Ia observations.  In this work we investigate the cosmological implications from X-ray of galaxy clusters and SNe data by considering two different classes of EOS evolving with redshift. In the first scenario (hereafter Model 1), the EOS parameter is defined by\\cite{CoHu99} \\begin{equation} Model\\,\\, 1 : \\,\\,\\, \\omega(z) = \\omega_o + \\omega_1 z, \\end{equation} whereas in the second, the EOS parameter reads\\cite{Lin03,PaCh03} \\begin{equation} Model\\,\\, 2: \\,\\,\\, \\omega(z) = \\omega_o + \\frac{\\omega_1 z}{1 + z} \\end{equation} where $\\omega_o$ and $\\omega_1$ are constants.  It should be noticed that the linear expression of model 1 yields a reasonable approximation for most quintessence models out to redshift of a few, and, of course, it should be exact for models where $\\omega(z)$ is a constant or varies slowly. The unsuitable aspect of the first expression is that it  grows with no limit at high  redshifts $z>1$ (for example, distance to the last scattering surface at $z_{lss}\\simeq 1100$). In order to fix such a problem, some authors\\cite{Lin03,PaCh03} have proposed the second form which has the advantage of giving finite $\\omega(z)$ for all $z$. In both cases, $\\omega_o$ is the current value of the EOS parameter and $\\omega_1$ defines its variation rate for $R$ close to the present epoch ($z=0$).  The paper is outlined as follows. Next section we set up the FRW equations for both parameterizations and the angular diameter distance necessary for our analysis using the X-ray luminosity of galaxy clusters. In section 3 we obtain our main results, namely: the observational constraints to the free parameters trough a joint analysis involving X-ray luminosity of galaxy clusters and SNe type Ia data. Finally, in section 4 we summarize the main conclusions and perspectives. ",
        "conclusions": "Nowadays, the signature of a dark energy or quintessence component is the most impressive observational fact of cosmology. As remarked elsewhere, we are living at a special moment where the emergence of new kind of ``standard cosmology\" seems to be inevitable. In the last few years, a growing attention has been paid for models with a time varying EOS parameter $\\omega(z)$. With basis on this sort of cosmological scenario, we have discussed here two simple possible parameterizations of the EOS obeyed by quintessence models as recently presented in the literature. Our results suggest that it is worthwhile to use the estimates of the gas mass fraction from galaxy clusters in joint analysis with the SNe Ia data since the derived constraints for $\\Omega_M$ (and other quantities) do not require any prior to this parameter. More important, we have also obtained constraints for $w_o$ and $w_1$ which have not been obtained before without a prior in $\\Omega_M$.  The parameterizations seems to be more efficient to explain these data set, once they get a lower $\\chi^2$; however they also have an additional parameter, and that is worthy of a study with more specific statistical criteria (Akaike or Bayesian information criteria, for example). A more detailed analysis of this kind will be investigated in the near future."
    },
    "0608/astro-ph0608153_arXiv.txt": {
        "abstract": "We present the effects of ellipticity of matter distribution in massive halos on the observation of supernovae. A pseudo elliptical Navarro-Frenk-White (NFW) mass model is used to calculate the introduced gain factors and observation rates of type Ia supernovae due to the strong lensing. We investigate how and to what extent the ellipticity in mass distribution of the deflecting halos can affect surveys looking for cosmologically distant supernovae. We use halo masses of $1.0 \\times 10^{12} h^{-1} M_{\\odot}$ and $1.0 \\times 10^{14} h^{-1} M_{\\odot}$ at redshifts $z_{d}=0.2$, $z_{d}=0.5$, and $z_{d}=1.0$, with ellipticities of up to $\\epsilon=0.2$. ",
        "introduction": "Supernovae have emerged as the most promising standard candles. Due to their significant intrinsic brightness and relative ubiquity they can be observed in the local and distant universe. Observational efforts to detect high-redshift supernovae have proved their value as cosmological probes. The systematic study and observation of these faint supernovae (mainly type Ia) has been utilized to constrain the cosmic expansion history \\citep{Goobar95, Perlmutter99, Schmidt98}. Light emitted from any celestial object is subject to lensing by intervening objects while traversing the large distances involved \\citep{KVB95} and the farther the light source, the higher its chance of being significantly lensed. Apart from the fact that gravitational lensing can limit the accuracy of luminosity distance measurements \\citep{Perlmutter03}, it can change the observed rate of supernovae as well.  Studying supernovae and their rates at high redshifts provide us with much needed information for constraining the measurements of the ellusive dark energy, as well as understanding the cosmic star formation rate and metal enrichment at high redshifts. In order to observe and, hence, study the faint high-redshift supernovae, one can raise the chance of observation by looking through clusters of galaxies or even massive galaxies (see \\citet{Sma02} and the references therein). These `gravitational telescopes' amplify the high-redshift supernovae and thereby increase the chance of their detection. However, this boost in observation is offset by the competing effect of depletion (Fig. 1), due to the field being spread by the deflector (amplification bias). For an assumed lens model and a given field of view it is not obvious which effect dominates the observation of supernovae through the halo. The net result depends on the deflector and source parametrs as well as the observational setup \\citep{GunnGoo03}.  Some research has been conducted on the feasibility of observing supernovae through cluster of galaxies \\citep[see, for instance, ][]{SaiRaySch00, GalMaoSha02, GunnGoo03}. These studies have not taken into account how the morphology (mainly the ellipticity) of these clusters as gravitational telescopes could change the expected supernova rate. In this paper, we investigate whether introducing ellipticity into the mass distribution of the deflecting halos can affect the observation of supernovae. For this purpose, we use a pseudo elliptical Navarro-Frenk-White (NFW) halo model with different values of ellipticity. Throughout the paper we assume the so-called concordance cosmology where $\\Omega _{m}= 0.3$, $\\Omega _{\\Lambda }= 0.7$, and $h_{100}= 0.67$, with $h_{100}=H_{0}/100$ km s$^{-1}$Mpc$^{-1}$. In $\\S$ 2 we briefly go over the NFW model and show how an analytical formalism for a pseudo elliptical NFW mass profile can be introduced. Strong lensing by thin deflectors as well as the way ellipticity can afffect the amplification is explained in $\\S$ 3. We present and discuss the results of our calculations in $\\S$ 4. ",
        "conclusions": "Aiming behind massive halos seem to be a good way to enhance the high-redshift supernovae surveys. The cumulative gains of such surveys seem insignificant at low redshifts ($z_{s}<0.2$) but the results are remarkable at higher redshifts. For deep observations where $m_{lim}>25$, the geometry of the intervening halo cannot be ignored. We have shown that introducing ellipticity in the (gravitational potential of) the mass distribution of a deflecting halo (here, for a galactic halo of mass $1.0 \\times 10^{12} M_{\\odot} h^{-1}$ as well as a middle-size cluster of galaxies with a mass of $1.0 \\times 10^{14} M_{\\odot} h^{-1}$) can affect the rate of observed supernovae by a few percent. It was shown that the farther the supernova survey probes, the more significant the effects of introduced ellipticity become.  It should be noted that this work does not involve a broad range of mass profiles for the halos (although we specify that the survey is limited to the virial area of the halo), nor does it address the much needed k-correction. Our calculations are actually an oversimplification due to the fact that a large, massive halo like a galaxy cluster has substructure which consists of the member galaxies, as well as large clouds of gas. A more sophisticated lens model with ellipticity should be employed to calculate the number rate of observed supernovae."
    },
    "0608/astro-ph0608479_arXiv.txt": {
        "abstract": "We have used very deep \\xmm\\ observations of the \\chandra\\ Deep Field-South to examine the spectral properties of the faint active galactic nucleus (AGN) population.  Crucially, redshift measurements are available for $84\\%$ (259/309) of the \\xmm\\ sample.  We have calculated the absorption and intrinsic luminosities of the sample using an extensive Monte Carlo technique incorporating the specifics of the \\xmm\\ observations.  Twenty-three sources are found to have substantial absorption and intrinsic X-ray luminosities greater than $10^{44}$~erg~s$^{-1}$, putting them in the ``type-2'' QSO regime.  We compare the redshift, luminosity and absorption distributions of our sample to the predictions of a range of AGN population models.  In contrast to recent findings from ultra-deep \\chandra\\ surveys, we find that there is little evidence that the absorption distribution is dependent on either redshift or intrinsic X-ray luminosity.  The pattern of absorption in our sample is best reproduced by models in which $\\sim 75\\%$ of the AGN population is heavily absorbed at all luminosities and redshifts. ",
        "introduction": "Recent X-ray spectral studies using \\xmm\\ and \\chandra\\ \\citep{piconcelli03,civano05,mateos05a} have begun to unravel the nature of the faint X-ray sources that make up the bulk of the extragalactic X-ray background (XRB) below 10~keV. Most of these sources are accretion powered AGN \\citep{barger03,page03}.  However, the population is far from homogeneous; the AGN exhibit a range of spectral properties, both in X-rays, and at other wavelengths. Unified models of AGN \\citep[e.g.][]{antonucci93}, have been reasonably successful in explaining the observational differences between the various classes of AGN/Seyfert galaxies. The X-ray spectra of the AGN can, in the majority of cases, be adequately described by a power-law model, in many cases with some degree of absorption.  The AGN with significant X-ray absorption are predominantly identified with narrow emission line galaxies, and are seen as higher redshift, higher luminosity analogues of nearby Seyfert 2 galaxies.  Population synthesis models such as those of \\citet{comastri95}, \\citet{gilli01} and \\citet{ueda03} use a superposition of faint AGN to reproduce the observed XRB. These models incorporate a large fraction of absorbed AGN in order to reproduce the hard spectrum of the XRB. \\citet{gilli01} found that the XRB could be reproduced using a model in which the absorbed and unabsorbed AGN shared a common intrinsic X-ray luminosity function.  A similar AGN population model was found to represent adequately the multi-wavelength properties of AGN in the GOODS dataset \\citep{treister04}.  In our previous work \\citep{dwelly05}, we compared the X-ray colour distribution of sources detected in the \\thirt\\ \\xmm\\ deep field with the X-ray colour distributions predicted by a number of model $N_H$ distributions. We found that the best match to the data was made by a $N_H$ distribution model in which absorbed and unabsorbed AGN share the same intrinsic luminosity function, and where the number of AGN per unit $\\log N_H$ is proportional to \\betaeight.   However, the results from other studies based upon deep X-ray surveys (e.g. \\citealt{piconcelli03,ueda03,cowie03,barger03,barger05}), appear to contradict some of the foundations of these simple synthesis models. These studies all found that on average, the absorbed AGN which are optically identified lie at lower redshifts, and have lower intrinsic luminosities, than their unabsorbed counterparts.  Clearly the relationship between absorption and luminosity in the AGN population requires further examination.  Therefore, in order to constrain the relationship between AGN luminosity, redshift and absorption we have undertaken a study of the X-ray properties of AGN in the \\chandra\\ Deep Field-South (\\cdfs).  In this study we examine the 500~ks \\xmm\\ European Photon Imaging Camera (EPIC) observations of the \\cdfs\\ (hereafter \\xcdfs) which provide a superb dataset for measuring the spectral properties of faint X-ray sources.  The EPIC imaging reaches to fluxes well below the break in the 2--5~keV source counts, covers around 0.19~deg$^2$ and contains enough photons to permit broad band X-ray spectral analysis for even the faintest sources.  The studies of \\citet{streblyanska04} and \\citet{braito05} have also used the \\xcdfs\\ dataset to investigate the X-ray spectra of a number of the brighter AGN in the field.  However, until this work, there has been no investigation of the X-ray properties of the entire source population detected in the \\xmm\\ imaging.  The EPIC data are complimented by \\chandra\\ observations \\citep{giacconi02,lehmer05}, which provide sub-arcsecond positions for the majority of the EPIC sources, allowing us to identify uniquely the X-ray sources with optical counterparts.  What is more, the entire EPIC field of view is covered by the \\combo\\ survey \\citep{wolf04}, providing photometric redshift estimates for nearly all $R<24$ optical counterparts.  We use an extensive Monte Carlo simulation process to recover the $N_H$ and intrinsic luminosity of the AGN detected in our sample. We also use the simulations to compare directly the distribution of sources with the predictions of a number of AGN population models.   This paper is laid out as follows.  In section \\ref{sec:xmm_data} we describe the reduction of the \\xmm\\ data, and the source detection process.  In section \\ref{sec:other_data} we detail how we have correlated this with the other datasets available in the \\cdfs.  In section \\ref{sec:nh_method} we introduce our Monte Carlo method for calculating the absorption and luminosity of the sample, and we demonstrate its fidelity.  In section \\ref{sec:results} we present the distribution of absorption, luminosity and redshift in the \\xcdfs\\ sample and compare it to the predictions of a number of AGN population models. Finally, in section \\ref{sec:discussion} we compare our results to those found by other studies and discuss the implications for AGN population models.  Throughout the paper we use a lambda-dominated flat cosmology with $H_0 = 70$~km~s$^{-1}$Mpc$^{-1}$, $(\\Omega_M,\\Omega_{\\Lambda }) = (0.3,0.7)$.  $S_{E_a-E_b}$ denotes the flux of a source in the observed $E_a$--$E_b$~keV band, corrected for Galactic absorption.  $L_{2-10}$ refers to an object's {\\em intrinsic} X-ray luminosity (that is, before absorption), in the rest-frame 2--10~keV energy band.  $N_H$ is the equivalent hydrogen column density in units of cm$^{-2}$. ",
        "conclusions": "We have analysed the distribution of absorption in a sample of AGN detected in the very deep \\xmm\\ observations of the \\cdfs. Importantly, spectroscopic or photometric redshift determinations are available for 84\\% of the X-ray sample. We determined the absorption and intrinsic luminosity of each AGN using a novel method which takes advantage of the high photon throughput and broad bandpass of \\xmm\\ EPIC.  The AGN in the \\xcdfs\\ sample were compared to the predictions of a number of model AGN populations, using Monte Carlo simulations which allow for the selection function and particulars of the \\xmm\\ observations.  We find no evidence for a decline in the fraction of AGN with significant absorption at high luminosities that has been reported by many authors.  The $N_H$ distribution models which most closely match the pattern of absorption in the \\xcdfs\\ sample are independent of redshift and luminosity.  Our sample contains at least 23 heavily absorbed AGN with QSO-like luminosities.  We postulate that nearly half of the objects without redshift determinations are also absorbed QSOs; the reason they are without \\combo\\ redshift estimates is because they lie in the ``photo-z desert'' ($1.5 < z < 2.5$).  In order to confirm this hypothesis, the redshifts of these optically faint objects must be determined; photo-$z$s incorporating near- and mid-infrared photometry should make this possible."
    },
    "0608/astro-ph0608223_arXiv.txt": {
        "abstract": "A laser diffraction experiment was conducted to study light propagation in air. The experiment is easy to reproduce and it is based on simple optical principles. Two optical sensors (segmented photo-diodes) are used for measuring the position of diffracted light spots with a precision better than  $0.1\\;\\mu m$. The goal is to look for signals of anisotropic light propagation as function of the laser beam alignment to the Earth's motion (solar barycenter motion) obtained by COBE. Two raster search techniques have been used. First, a fixed laser beam in the laboratory frame that scans due to Earth's rotation. Second, an active rotation of the laser beam on a turntable system. The results obtained with both methods show that the course of the light rays are affected by the motion of the Earth, and a predominant quantity of first order with a $\\Delta c/c=-\\beta (1+2a)\\cos \\theta$ signature with $a=-0.4106\\pm 0.0225$ describes well the experimental results. This result differs in a amount of 18\\% from the Special Relativity Theory prediction and that supplies the value of $a=-1/2$ (isotropy). ",
        "introduction": "There are several physical reasons, theoretical and experimental, that could justify a search for anisotropies in light propagation. It is well known that Lorentz and Poincar$\\acute{e}$  were the first ones to build the major part of the relativity theory on the basis of the ether concept as an inertial rest frame, and it is fully compatible with the Einstein's Special Relativity Theory (SRT). There are also some test theories of SRT, where the Lorentz transformations are modified. For example, an ether theory that maintains the absolute simultaneity and is kinematically equivalent to Einstein SRT was constructed \\cite{mansouri77}. These test theories are considered useful to examine potential alternatives to SRT. On the other hand, the reconstruction of the SRT, on the basis of the Lorentz-Poincar$\\acute{e}$ scheme implies in an undetectable ether rest frame (non ether drift) at least in the first order \\cite{lorentz}.  This comportment of the Lorentz-Poincar$\\acute{e}$, as well as, of the Einstein theories arise because they do not rule the whole physics, for instance they do not involve gravitation. It is also well known that the presence of a gravitational field breaks the Lorentz symmetry. For instance, the Sagnac effect \\cite{sagnac13} shows that two opposite light beams travel in different time intervals the same closed path on a rotating disk. This effect is difficult to reconcile with SRT, because on the basis of the Lorentz transformation an isotropic light propagation is required in each point of the rotating disk, and a null effect is predicted. It is necessary to invoke the equivalence principle between an accelerated system and a gravitational field in order to take into account this effect. However, in the generalized Sagnac effect \\cite{wang04} observed in a light waveguide loop consisting of linearly and circularly segments, any segment of the loop contributes to the phase difference between two counter-propagating light beams in the loop, and the situation is again difficult to reconcile with relativistic theories. On the other hand, according to Fox \\cite{fox64},it is possible to preserve the general Lorentz Poincar$\\acute{e}$ symmetry group without assuming the constancy of light speed. There are also the so called extended theories, where the SRT is modified in order to including the Planck scale parameters \\cite{camelia01}(double relativity theories), suggesting several dispersion relations that include theories where an energy dependent speed of light \\cite{magueijo99} is claimed.  There is also evidences suggesting that the propagation of light over cosmological distances has anisotropic characteristics\\cite{nodland97}, with dependence on direction and polarization. This picture is in agreement with the interpretations of the COBE\\cite{smoot91} measurements giving the Earth's ``absolute'' velocity in relation to the uniform cosmic microwave background radiation(CMBR). Of course, there are also interpretations claiming that the COBE measurements give only a velocity for the ``relative'' motion between the Earth and the CMBR \\cite{yaes93}. For instance, it is possible to remove the Earth motion to obtain a ``virtual'' image, where an isotropic distribution of CMBR with small fluctuations ($\\delta T/T \\sim 10^{-5}$) can be seen. There is also some indirect evidences \\cite{hatch02} via analysis of the Global Positioning System (GPS) for the presence of an ether-drift velocity. The assumption of a preferred frame agrees also with a previous analysis made by Brans and Stewart \\cite{brans73} where a description of the topology of the universe has imposed a preferred state of rest so that the principle of special relativity, although locally valid, is not globally applicable.  So far, several tests about violation of the isotropy of the speed of light have been made. In most cases, the tests involve the so called round-trip test of light-speed isotropy like Michelson-Morley experiment and all its variants. Particularly, Miller \\cite{miller33,cahill05} has claimed a non-null results in the M-M experiments. These aspects are presented in Appendix A.  On the other hand, there are also several one-way test of light isotropy experiments. In most cases, they have claimed a null result \\cite{cialdea72, krisher90, turner64, gagnon88}, and some have claimed success \\cite{kolen82,silvertooth86}. Particularly, Silvertooth has claimed an experimental detection of the ether drift velocity using a device capable of detecting the beams arriving in opposite directions\\cite{silvertooth86}. Silvertooth report in 1986 a light anisotropy toward the direction of Leo constellation and compatible with COBE results. The experiment is an unusual double interferometer, an arrangement of light paths and detectors hard to be reproduced. In addition, the presence of a feedback into the laser is quite probable.  In this paper, we report results of a search for anisotropic light propagation as a function of the laser beam alignment relative to the Earth's velocity vector, using a diffraction device. The method is based on simple optical principles. Initial attempts have used digital images of the diffraction spots. However, this method was working in the limit of sensitivity. In other words, the signal's size was close to the measurement resolution. Now, our results are obtained by using the highly sensitive segmented photo-diodes to measure the position of diffracted light spots. In Section 2, the experimental setup and the basic operating principles of the diffractometer are presented. The Earth's velocity vector on the basis of the Doppler shift of the CMBR results are presents in Section 3. In Section 4, the two scanning methods and their results are presented, and finally Section 5 contains our conclusions. ",
        "conclusions": "The discovery of a dipole anisotropy in the CMBR is interpreted as a Doppler shift produced by the Earth' motion (solar barycenter). An experimental survey has been made in order to test if the Earth's velocity is relevant on light propagation in a quantity of first order. The measurements have been obtained by using a laser diffraction experiment mounted on turntable system. Two optic sensors (segmented photo-diodes) were used for measuring the position of diffracted light spots with a precision better than  $0.1\\;\\mu m$. The experiment is easy to reproduce, and it is based on simple optical principles. Two raster search techniques have been used to look for signals of anisotropic light propagation as a function of the laser beam alignment relative to the Earth's motion. The results obtained with both methods show that the course of the rays is affected by the motion of the Earth.  In the scan subjected to Earth's rotation, the sensitivity of the diffraction experiment to the Earth's velocity begins when the altitude on the horizon of the Leo-CC is $h_1 \\sim 20^0$, and it finishes when the altitude is $h_2 \\sim -20^0$. In this passage on the horizon, the Leo-CC has a variation in its azimuthal angle of around $17^0$, and it is equivalent to an azimuthal rotation of the diffractometer by the same angle (around $17^0$). This effect provides a dipole behavior ($\\propto \\cos \\theta_A$) in a narrow angular window ($\\Delta \\theta_A \\sim 17^0$), and it is observed as a peak followed by a depression in the first passage, or a depression followed by a peak in the second passage (see Fig.6 upper).  In the scan subjected to an active rotation, from zero to 180 degree relative to the Earth's velocity direction, a complete set of measurements can be done in $10$ minutes. Consequently it is free from DRIFT-long-term timing variations. The experimental data are susceptible of being interpreted by the test theory of Mausouri and Sexl\\cite{mansouri77} where the one way speed of light is anisotropic by the amount $c(\\theta)=c-v(1+2a)cos \\theta$. The parameter $a$ extract from a fit of the data is $a=-0.4106\\pm 0.0225$, and it differs from SRT prediction with $a=-0.5$ by an amount of 18\\%.  We remark that the CMBR dipole is a frame dependent quantity. We can thus determine the ``absolute rest frame'' of the universe as that in which the CMBR dipole would be zero. In short, our results point out that it is not possible to neglect the preferred frame imposed by  cosmology."
    },
    "0608/astro-ph0608709_arXiv.txt": {
        "abstract": "{} {The aim of this paper is to characterise the star formation activity in the poorly studied embedded cluster Serpens/G3-G6, located $\\sim$~45$\\arcmin$ (3~pc) to the south of the Serpens Cloud Core, and to determine the luminosity and mass functions of its population of Young Stellar Objects (YSOs). } {Multi-wavelength broadband photometry was obtained to sample the near and mid-IR spectral energy distributions to separate YSOs from field stars and classify the YSO evolutionary stage. ISOCAM mapping in the two filters LW2 (5-8.5 $\\mu$m) and LW3 (12-18 $\\mu$m) of a $19\\arcmin \\times 16\\arcmin$ field was combined with JHK$_S$ data from 2MASS, K$_S$ data from Arnica/NOT, and L$\\arcmin$ data from SIRCA/NOT. Continuum emission at 1.3~mm (IRAM) and 3.6~cm (VLA) was mapped to study the cloud structure and the coldest/youngest sources. Deep narrow band imaging at the 2.12~$\\mu$m S(1) line of H$_2$ from NOTCam/NOT was obtained to search for signs of bipolar outflows. } {We have strong evidence for a stellar population of 31 Class\\,II sources, 5 flat-spectrum sources, 5 Class\\,I sources, and two Class\\,0 sources. Our method does not sample the Class\\,III sources. The cloud is composed of two main dense clumps aligned along a ridge over $\\sim$~0.5~pc plus a starless core coinciding with absorption features seen in the ISOCAM maps. We find two S-shaped bipolar collimated flows embedded in the NE clump, and propose the two driving sources to be a Class\\,0 candidate (MMS3) and a double Class\\,I (MMS2). For the Class\\,II population we find a best age of $\\sim$~2~Myr and compatibility with recent Initial Mass Functions (IMFs) by comparing the observed Class\\,II luminosity function (LF), which is complete to 0.08~$L_{\\odot}$, to various model LFs with different star formation scenarios and input IMFs. } {} ",
        "introduction": "\\label{intro}  The Serpens Cauda Clouds are among the darkest regions in a complex of clouds called the Aquila Rift \\citep{dam85}. The young embedded cluster in the Serpens Cloud Core has been well studied at most wavelengths, see \\citet{eir06} for a recent review. About 45' to the south there is also a region with active star formation. \\citet{coh79} found 4 optical T Tauri stars in a 30$\\arcsec$ field and named them Ser/G3-~G6. \\citet{cla90,cla91} mapped the area surrounding Ser/G3-~G6 in the ammonia 1,1 emission line and found two NH$_3$ cores, one on each side of the optically visible stellar group: Ser/G3-~6NE and Ser/G3-~G6SW. \\citet{zie99} surveyed 3.14 square degrees in the optical lines of [S\\,II] and found the Herbig-Haro object HH\\,476 close to the Ser/G3-G6SW core. The energy source of HH\\,476 was identified by \\citet{wu02} to be IRAS18265+0028. At this position \\citet{per94} found a H$_{2}$O maser. Recently, \\citet{har06} found 24 YSO candidates in a 12' $\\times$ 12' area in this region (which they refer to as Cluster B) from a mid-IR survey with Spitzer.  When preparing the {\\it ISOCAM star formation survey},\\footnote{One of the guaranteed time programs with the camera aboard the Infrared Space Observatory (ISO).} this region was defined based on the ammonia maps of \\citet{cla90} and named {\\em Serpens NH3}. The field is relatively opaque at optical wavelengths and has a peak brightness at 60 $\\mu$m \\citep{zha88}.  We will use the distance of 225 $\\pm$ 55 pc to the Serpens Cauda Clouds according to \\citet{str03}, keeping in mind the caveat of the same authors that the cloud complex is possibly 80 pc deep.  Preliminary ISOCAM results from this region were presented in \\citet{kaa99c} and \\citet{kaa99b}. In this paper we present the full results of a 17'$\\times$ 19' ISOCAM survey in LW2 (6.7 $\\mu$m) and LW3 (14.3 $\\mu$m), cross-correlated with the 2MASS point source catalogue. In addition we present deep follow-up $H_2$ line (2.122 $\\mu$m), $K_S$ band (2.15 $\\mu$m), and $L'$ band (3.8 $\\mu$m) imaging from the Nordic Optical Telescope, as well as IRAM 1.3~mm continuum and VLA 3.6~cm mapping of the most active part of the region.  We describe the various observation sets with reductions and results in Sect.~\\ref{obs}, the use of IR excesses to extract YSO candidates in Sect.~\\ref{irex}, the classification of the YSO population in Sect.~\\ref{ysopop}, the luminosity and mass function in Sect.~\\ref{lfmf}, and the cloud structure and the outflow sources in Sect.~\\ref{flow}. ",
        "conclusions": "\\label{con}  An embedded cluster of YSOs was found from an ISOCAM survey around the optical stellar group Ser/G3-G6. Mid-IR photometry in two broadband filters centred on 6.7 and 14.3 $\\mu$m obtained with ISOCAM was combined with data from various telescope/instrument configurations to present a multi-wavelength study. This reveals low-mass YSOs in the evolutionary stages from Class\\,0 to Class\\,III. We propose 2 Class\\,0 candidates, 5 Class\\,I candidates, 5 flat-spectrum sources, 31 Class\\,II objects, and 3 Class\\,IIIs (with some IR-excess). Our selection criterium is IR excess, which means that our sample is severly incomplete for Class\\,III sources in general. This also means that we cannot estimate the disk fraction in the YSO population. The number fraction of Class\\,I to Class\\,II sources in the Ser/G3-G6 IR cluster (5/31) is quite typical for star-forming regions and comparable to the $\\rho$ Ophiuchi and Chamaeleon~I clusters, i.e., quite different from the unusually large fraction of Class\\,Is found in the Serpens Cloud Core \\citep{kaa04}.  The IRAM 1.3~mm continuum mapping detects several bright continuum sources in the two ammonia cores Ser/G3-G6NE and Ser/G3-G6SW, and one outside that is suggested as a candidate prestellar core. MMS1-c and MMS3 are proposed to be Class\\,0 candidates. Both of them are also VLA 3.6~cm radio continuum sources. Deep NOTCam imaging in the 2.122 $\\mu$m line of H$_2$ reveals signs of two complex outflows in the Ser/G3-G6NE core. A deep H$_{\\alpha}$ image shows that the flows are embedded and only the previously known HH 476 is seen optically. We suggest that the driving sources of these two outflows are MMS2 and MMS3.  Comparing with existing pre-main sequence evolutionary models, we find that for any reasonable assumptions on age for the Class\\,II sources, our sample extends well into the BD mass regime. The Class\\,II sample was joined with the Class\\,II sample from the Serpens Cloud Core to improve statistics. Model LFs calculated for a set of ages and IMFs give a best age of 2~Myr for which both the \\citet{sca98} and \\citet{kro01} IMFs are compatible with the observed LF down to the completeness limit at $0.08~L_{\\odot}$. The Class\\,II LF can be described by co-eval star formation $\\sim$~2 Myr ago, continuous star formation over the last 2~Myr, or a burst of star formation of various duration that took place 2~Myr ago. There is a peak in the observed LF that cannot be used to constrain the age because it coincides with the completeness limit. Nevertheless, the general shape of the LF does not permit a much younger or older population. The completeness limit must be shifted to $0.04~L_{\\odot}$ to possibly distinguish between different IMFs and star formation scenarios. This should be feasible with the deeper mid-IR surveys now available from Spitzer."
    },
    "0608/astro-ph0608015_arXiv.txt": {
        "abstract": "{} {Our goals are to compare the CS, N$_2$H$^+$ and dust distributions in a representative sample of high-mass star forming dense cores and to determine the physical and chemical properties of these cores.} {We compare the results of CS(5--4) and 1.2~mm continuum mapping of twelve dense cores from the southern hemisphere presented in this work, in combination with our previous N$_2$H$^+$(1--0) and CS(2--1) data. We use numerical modeling of molecular excitation to estimate physical parameters of the cores.} {Most of the maps have several emission peaks (clumps). Mean sizes of 17 clumps, having counterparts in continuum and CS, are 0.30(0.06)~pc (continuum) and 0.51(0.07)~pc (CS). For the clumps with IRAS sources we derived dust temperatures: 24--35~K, masses: 90--6900~$M_{\\odot}$, molecular hydrogen column densities: $(0.7-12.0)\\times 10^{23}$~cm$^{-2}$ and luminosities: ($0.6-46.0)\\times 10^4~L_{\\odot}$. LVG densities towards CS peaks within the 50$''$ beam (0.56~pc at 2.3~kpc, the average distance of our sample source) vary from source to source in the range: (3--40)$\\times$\\,$10^5$~cm$^{-3}$. Masses calculated from LVG densities are higher than CS virial masses and masses derived from continuum data, implying small-scale clumpiness of the cores. The molecular abundances towards IRAS sources in eight objects are $X$(CS)=$(0.3-2.7)\\times10^{-9}$ and $X$(N$_2$H$^+$)=$(0.3-4.4)\\times10^{-10}$. The CS and continuum maps have been compared with each other and with the N$_2$H$^+$(1--0) maps. For most of the objects, the CS and continuum peaks are close to the IRAS point source positions. The CS(5--4) intensities correlate with continuum fluxes per beam in all cases, but only in five cases with the N$_2$H$^+$(1--0) intensities. The study of spatial variations of molecular integrated intensity ratios to continuum fluxes per beam reveals that $I$(N$_2$H$^+$)/$F_{1.2}$ ratios drop towards the CS peaks for most of the sources, which can be due to a N$_2$H$^+$ abundance decrease. For CS(5--4), the $I$(CS)/$F_{1.2}$ ratios show no clear trends with distance from the CS peaks, while for CS(2--1) such ratios drop towards these peaks. Possible explanations of these results are considered. The analysis of normalized velocity differences between CS and N$_2$H$^+$ lines has not revealed indications of systematic motions towards CS peaks. } {} ",
        "introduction": "The first step in most studies of dense cores of molecular clouds is their identification, either by a visual inspection or by rather sophisticated numerical analysis of molecular line maps. Yet, the maps of ``traditional\" tracers of dense gas (CS, HCN, HCO$^+$, NH$_3$, N$_2$H$^+$ etc.) appear to be different from each other in many cases. Most distinct differences are often observed between CS and N$_2$H$^+$ maps. In low-mass starless cores (e.g. Tafalla et al. 2002) the CS emission vanishes towards the core center while the N$_2$H$^+$ emission is still high in the interior regions. Some more turbulent starless clumps show strong CS emission in comparison with N$_2$H$^+$ (e.g. Williams \\& Myers 1999, Olmi et al. 2005). The optically thick CS lines in the latter case are usually asymmetric and blueshifted, implying that they are probably tracing infall motions. Both these cases can be explained by time dependent chemical models including depletion molecular species onto grains (Bergin \\& Langer 1997, Bergin et al. 1997, Li et al. 2002, Shematovich et al. 2003, Aikawa et al. 2003), suggesting different evolutionary stages of clumps within parent molecular cloud. In the cores with embedded stellar objects, especially those of high mass, the morphologies of molecular maps become more complicated. This is due to the overlapping of the emission coming from the quiescent gas with that from outflows and hot cores, where dust grain mantles evaporate because of their proximity to YSOs. The N$_2$H$^+$ intensities towards these objects may drop, while CS peaks towards YSO positions (e.g. Ungerechts et al. 1997, Bottinelli \\& Williams 2004). Thus, chemical models of starless clumps cannot be used to explain the observed differences between CS and N$_2$H$^+$ maps in high-mass star forming regions.  In the last decade, a large sample of high-mass star forming cores associated with water masers have been mapped in the CS(2--1) line (Zinchenko et al. 1994, 1995, 1998). Many of these cores have subsequently been mapped in different molecular lines, including N$_2$H$^+$(1--0) (Pirogov et al. 2003, hereafter Paper~I). Thus, the maps can be compared and important quantities can be determined for these regions, such as density distribution and chemical composition. Yet, in contrast to the N$_2$H$^+$(1--0) lines, which are optically thin in most cases (Paper~I), the CS(2--1) lines are likely optically thick and possible differences between maps in these lines could be connected both with optical depth effects and chemical differences.  In this paper we present the results of 1.2~mm dust continuum and CS(5--4) observations in several southern sources. Optically thin dust emission is known to trace closely the gas component in high density regions and it is insensitive to molecular abundance variations and radiative transfer effects. This can be used to separate optical depth and chemical effects in molecular line maps.  The goal of the paper is to analyze the dust continuum and CS(5--4) data in comparison with our previous CS(2--1) and N$_2$H$^+$(1--0) results in order to get reliable information on the density and chemical structure of the cores. The paper also contains estimates of physical parameters including sizes, masses as well as density distributions derived from CS(5--4) and CS(2--1) data and molecular abundances for several sources. ",
        "conclusions": "In order to get reliable information on the density and chemical structure of the dense cores associated with high-mass star forming regions twelve objects from southern hemisphere have been mapped in the CS(5--4) line and in dust continuum at 1.2~mm. We have compared CS(5--4) and continuum data with each other as well as with the N$_2$H$^+$(1--0) (Paper~I) and the CS(2--1) (Zinchenko et al. 1995) data. Besides, physical parameters of the cores have been derived from the CS(5--4) and the continuum data.  The results can be summarized as follows:  \\begin{enumerate} \\item Most of the maps have several emission peaks (clumps). Mean sizes of 17 clumps, having counterparts in continuum and CS, are 0.30(0.06)~pc (continuum) and 0.51(0.07)~pc (CS). The CS virial masses lie in the range $\\sim$140\\,--\\,1630~$M_{\\odot}$. For the clumps with IRAS sources we derived dust temperatures: 24\\,--\\,35~K, masses: 90--6900~$M_{\\odot}$, molecular hydrogen column densities: (0.7\\,--\\,12.0)$\\times$\\,$10^{23}$~cm$^{-2}$ and luminosities: (0.6\\,--\\,46.0)$\\times$\\,$10^4~L_{\\odot}$.  \\item Using the CS(5--4) and the CS(2--1) data (Zinchenko et al. 1995) LVG densities in eight sources have been calculated. Densities towards CS peaks within the 50$''$ beam (0.56~pc at 2.3~kpc, the average distance of our sample source) vary from source to source in the range: (3\\,--\\,40)$\\times$\\,$10^5$~cm$^{-3}$. In four sources, the density falls by about an order of magnitude at $\\sim$\\,0.8\\,--\\,2~pc from the CS peaks which could be connected either with the existence of lower density clumps or density gradients within clumps. The masses calculated from LVG densities are higher than virial masses and masses derived from continuum data implying small-scale clumpiness of the cores. The CS abundances have been calculated towards IRAS positions within 50$''$ Gaussian beam for eight sources: $X$(CS)=(0.3\\,--\\,2.7)$\\times$\\,$10^{-9}$. The N$_2$H$^+$ abundances for the same positions lie in the range: $X$(N$_2$H$^+$)=(0.3\\,--\\,4.4)$\\times$\\,$10^{-10}$.  \\item For most of the objects the CS and continuum peaks are close to the IRAS point source positions. The CS(5--4) integrated intensities correlate with continuum fluxes per beam almost linearly in all cases while significant correlations between the CS(5--4) and the N$_2$H$^+$(1--0) integrated intensities are found only in five sources. The highest CS to N$_2$H$^+$ integrated intensity ratios ($\\sim 10$) are found in the vicinity of the most luminous IRAS sources. The lowest ratios ($\\la 0.1$) are observed towards positions close to the centers of the N$_2$H$^+$ clumps without CS counterparts and IRAS sources.  \\item The study of spatial variations of molecular integrated intensity ratios to continuum fluxes reveals that for most of the sources the $I$(N$_2$H$^+$)/$F_{1.2}$ ratios drop towards CS and dust emission peaks near IRAS sources, likely implying underabundance of N$_2$H$^+$ towards these peaks. The $I$(CS)/$I$(N$_2$H$^+$) ratios are clearly enhanced towards the three most luminous sources, whereas the $I$(N$_2$H$^+$)/$F_{1.2}$ ratios drop towards these sources. Possible explanations of these results could involve dissociative recombination of N$_2$H$^+$ or accelerated collapse processes. For CS(5--4) the $I$(CS)/$F_{1.2}$ ratios show no clear trends with distance from the CS peaks while for CS(2--1) such ratios drop towards these peaks implying line saturation effects in several sources.  \\item The analysis of normalized velocity differences between CS and N$_2$H$^+$ lines has not revealed indications of systematic motions in the sources towards CS peaks. A possible indication for infall motions has been found towards one position in G285.26. The CS(5--4) line widths in several sources have a tendency to decrease with distance from CS peaks which could be connected both with optical depth effects and with enhanced dynamical activity towards centers of the cores.  \\end{enumerate}"
    },
    "0608/astro-ph0608564_arXiv.txt": {
        "abstract": "Some tests of fundamental physics - the equation of state at supra-nuclear densities, the metric in strong gravity, the effect of magnetic fields above the quantum critical value - can only be measured using compact astrophysical objects: neutron stars and black holes. The Extreme Physics Explorer is a modest sized ($\\sim$500~kg) mission that would carry a high resolution ($R\\sim$300) X-ray spectrometer and a sensitive X-ray polarimeter, both with high time resolution ($\\sim$5~$\\mu$s) capability, at the focus of a large area ($\\sim$5~sq.m), low resolution (HPD$\\sim$1~arcmin) X-ray mirror. This instrumentation would enable new classes of tests of fundamental physics using neutron stars and black holes as cosmic laboratories. ",
        "introduction": "\\label{sect:intro}  %  All scientific fields go through three phases: {\\em discovery}, {\\em understanding}, and {\\em tool}. In the {\\em discovery} phase there is widespread excitement over the new field and prizes are won as the basic properties of the field are explored; in the {\\em understanding} phase interest narrows to those specialists willing to commit to the intensive research that leads to physical understanding; at some point though the remaining problems become so small that the field becomes boring; not until the experts have achieved that understanding does the field become of wide interest again, but this time as a {\\em tool} to unlock other fields of study.  X-ray astronomy is now over 40 years old and the first area of study within X-ray astronomy - the neutron star and black hole X-ray binaries - have long been in the {\\em understanding} phase, and interest in the field has shrunk to a relatively small expert group. The progress made by this group however, largely thanks to the {\\em Rossi X-ray Timing Explorer} (Swank 1998), has been impressive, so that many of the key astrophysical properties of these exotic systems have been pinned down. This progress has paved the way for the opening of the third phase: the use of compact X-ray binaries as a {\\em tool} to explore the {\\em extreme physics} for which neutron stars and black holes are natural laboratories.  Here I propose the {\\em Extreme Physics Explorer}, a moderate cost, moderate risk, mission to allow precision measurements of the fundamental physics that is only probed with these exotic objects. ",
        "conclusions": "The {\\em Extreme Physics Explorer} is a modest sized ($\\sim$500~kg) mission that would carry an X-ray spectrometer and an X-ray polarimeter, both with high time resolution ($\\sim$5~$\\mu$s), at the focus of a large area ($\\sim$5~sq.m), low resolution (HPD$\\sim$1~arcmin) X-ray mirror.  This instrumentation would enable a new class of tests of fundamental physics using neutron stars and black holes as cosmic laboratories."
    },
    "0608/astro-ph0608087_arXiv.txt": {
        "abstract": "We illustrate the energy transfer during planetary flybys as a function of time using a number of flight mission examples.  The energy transfer process is rather more complicated than a monotonic increase (or decrease) of energy with time.  It exhibits temporary maxima and minima with time which then partially moderate before the asymptotic condition is obtained.  The energy transfer to angular momentum is exhibited by an approximate Jacobi constant for the system. We demonstrate this with flybys that have shown unexplained behaviors: i) the possible onset of the ``Pioneer anomaly\" with the gravity assist of Pioneer 11 by Saturn to hyperbolic orbit (as well as the Pioneer 10 hyperbolic gravity assist by Jupiter) and ii) the Earth flyby anomalies of small increases in energy {\\it in the geocentric system} (Galileo-I, NEAR, and Rosetta, in addition discussing the Cassini and Messenger flybys).  Perhaps some small, as yet unrecognized effect in the energy-transfer process can shed light on these anomalies. ",
        "introduction": "\\label{intro}  The use of planetary flybys for gravity assists of spacecraft became of wide interest during the 1960's, when the Jet Propulsion Laboratory (JPL) first started thinking about what became the ``Grand Tours\" of the 1970's and 1980's (the Voyager missions).  The concept was to use flybys of the major planets to both modify the direction of the spacecraft and also to add to its heliocentric velocity. At the time many found it surprising that energy could be transferred to a spacecraft from the orbital-motion angular-momentum of a planet about the Sun, despite the fact it had been known since the works of Lagrange, Jacobi, and Tisserand on the three-body problem \\cite{moulton, danby}, that the energies of comets could be affected by passing near Jupiter.\\footnote{In that same period observations that the solar system could evolve \\cite{hills} were similarly met with skepticism.}  Even in the simplest, circular restricted 3-body problem \\cite{danby}, it is {\\it not} that the energy of each object is conserved, only that the total energy of the entire system is conserved. Flybys can both give kinetic energy to a spacecraft (to boost its orbital velocity) and also can take kinetic energy from it (to slow it down for an inner body encounter).  Hohmann developed a powerful analysis tool for gravity-assist navigation, the method of patched conics \\cite{hohmann,wiesel}.  At JPL clearer understandings of gravity assists were obtained from the works of \\cite{min} and  \\cite{flandroref}. In this technique, a two-body (Kepler) orbit about the Sun is first kinematically patched onto a two-body hyperbolic orbit about the planet.  Then the spacecraft proceeds along the planet-centered hyperbolic orbit during the flyby, and finally a new two-body orbit about the Sun is patched onto the post-flyby hyperbolic orbit.  The question arises in such an approach as to where to do the patch; obviously somewhere along the hyperbolic asymptotes, but where? This can be addressed by considering how large a region of space is controlled by the flyby planet. When the region is reduced to a planet-centered sphere of radius $r$, and the distance between the Sun and planet is given by $R$, a useful {\\it sphere of influence} is given by the {\\it Roche limit}, $r/R = (3m/M_\\odot)^{1/3}$, where $m/M_\\odot$ represents the ratio of the planet's mass to the Sun's mass. However, the point where the errors in the respective two-body orbits about the Sun and planet are equal is determined not by the exponent 1/3, but by the exponent 2/5 as derived by Lagrange \\cite{wiesel}. It actually makes little difference which exponent is used to define this matching point, as this whole calculation is an approximation to the actual dynamics and, in practice, the gravity assist trajectory used for space navigation is computed by numerical integration of the complete equations of motion.  However, the patched-conic technique is useful for purposes of searching for a fuel-conserving gravity-assist trajectory from a large family of possibilities, and especially when multi flybys of more than one planet are used, for example in the Galileo mission \\cite{russell}. In \\cite{flandroref} and elsewhere \\cite{jdaEGA, VA} the process has been simply and intuitively described and we use it here to gain insights into the nature of the flyby anomaly. (See Figure \\ref{flandrofig}.)    \\begin{figure}[h!] \\begin{center} \\noindent \\psfig{figure=flandroNA4.eps,width=3in}  % \\end{center} \\caption{  \\small{Geometry of a flyby, modified from \\cite{flandroref}.} \\label{flandrofig}} \\end{figure}   The simple vector velocities in the heliocentric system are added to the orbital velocity of the planet (taken to be constant).  The initial and final velocities in the heliocentric system are ${\\bf v}_i$ and ${\\bf v}_f$.    The initial and final velocities in the planetary system are ${\\bf v}'_i$ and ${\\bf v}'_f$.  The velocity of the planet in the solar system is ${\\bf v}_p$.  The change in kinetic energy per unit mass is \\begin{equation} \\Delta {\\cal K} = ({\\bf v}_f \\cdot {\\bf v}_f - {\\bf v}_i\\cdot {\\bf v}_i)/2.  \\label{fly1} \\end{equation} A little algebra \\cite{flandroref,VA} gives one \\begin{equation} \\Delta {\\cal K} = {\\bf v}_p \\cdot ({\\bf v}_f' - {\\bf v}_i').                 \\label{fly2} \\end{equation}  Roughly speaking, in the planetary system which rotates anticlockwise, if a satellite in the ecliptic comes from inside the planetary orbit, travels behind the planet, and then goes around it counter-clockwise, kinetic energy will be added to the spacecraft.  Contrarily if a satellite comes from inside the planetary orbit, travels in front of the planet, and then goes around it clockwise, kinetic energy will be taken away from the spacecraft.  Of course one is not violating conservation of energy.  The energy (and angular-momentum change) is absorbed by the planet that is  being flown by.  However, for such a massive body the relatively infinitesimal change is not noticeable.  Further, there {\\it is} (in high approximation) a conserved quantity for the spacecraft in the barycentric system, Jacobi's integral \\cite{moulton,danby}: \\be J = - C/2 = ({\\cal K} + {\\cal V}) + {\\cal L} = ({\\cal E}) - \\mathbf{\\omega \\hat{z} \\cdot r \\times v}, \\label{jacobi} \\ee where $\\{{\\cal V},{\\cal L},{\\cal E}\\}$ are the potential, rotational-potential, and total energies, respectively, per unit mass, $\\mathbf{\\omega}$ is the angular velocity of the planet (system) whose vector is aligned with $\\mathbf{\\hat{z}}$, the unit vector normal to planet's rotational plane. Eq. (\\ref{jacobi}) is exactly a constant  in the circular restricted 3-body problem, and shows how kinetic energy can be exchanged with angular momentum during a flyby.  In this paper we are going to discuss two unusual results that are associated with planetary flybys.  The first is the possible ``onset\" of the Pioneer anomaly when a major planet gravity assist was used to reach hyperbolic orbit.   The second result is even more surprising, that during Earth flybys at least three craft have exhibited a small velocity increase so that the outbound hyperbolic orbits {\\it in the Earth system} have different energies than the inbound orbits.  In the following sections we will discuss these anomalies from the energy transfer viewpoint.   For the orbits we discuss we obtained information from JPL's Solar System Dynamics web site.\\footnote{{\\tt http://ssd.jpl.nasa.gov/} \\label{JPLdynamics}} Note that, although this information is precise enough to demonstrate the general energy-transfer properties, to glean out the small anomalies we are ultimately interested in understanding, more precise orbit determinations are necessary, particularly with regard to the Pioneer anomaly.   We  conclude with a discussion.  (In Appendix A we describe how the original trajectory data was obtained and in Appendix B we give further details on the accuracy of the calculations used in this paper.) ",
        "conclusions": "\\label{disc}  Having described the dynamical situations of these anomalous flybys, the question arises if there is any insight that might be obtained.   We note again that the best-fit values of the rotational frequencies obtained for the Jacobi constants differ bt 10 to 20 \\% from the physical values.  {\\it A priori} this could in principle be due to missing mass, a misorientation of the invariable plane, an added force, or a light speed anomaly, all of which could be mapped into one another without better knowledge.  In this light, it is useful to recall another result from the study of the Pioneer anomaly, that there are apparent annual and diurnal terms on top of the constant anomaly \\cite{pioprd}.  These are most apparent in clear data sets that came when the Pioneers were at large distances from the Sun.  In Figure \\ref{annual-diur} we show examples of this type of signal.   \\begin{figure}[h!] \\begin{flushleft} \\noindent \\epsfig{figure=ODP_1d_batch.eps,width=2.95in}% \\psfig{figure=opp96.eps,width=2.45in}% \\end{flushleft} \\vskip -10pt \\caption {First we give ODP 1-day batch-sequential acceleration residuals using the Pioneer 10  data set from 1987.0 to 1998.5 \\cite{pioprd}. Maneuver times are indicated by the vertical dashed lines. The annual term is apparent as an overlay to the constant baseline.  Next we give CHASMP program  acceleration residuals from 23 November 1996 to 23 December 1996 \\cite{pioprd}. A clear modeling error is represented by the solid diurnal curve. (An annual term maximum is also seen as a background.) \\label{annual-diur}} \\end{figure}   In discussions of the Pioneer anomaly, these annual/diurnal terms should be kept in mind.  Going on to the Earth flybys, since the mass distribution of the Earth and the dynamics of EGAs are supposedly well known, people have started to ask if there is at least some phenomenological pattern to the EGA anomalies \\cite{awill,cdit}.  Note that the simplest Newtonian version of the problem is the hyperbolic orbit in the geocentric system: \\bea r &=&\\frac{a(\\epsilon^2 -1)}{1+\\epsilon \\cos\\theta} = \\frac{r_F(\\epsilon + 1)}{1+\\epsilon \\cos\\theta}, ~~~~~~~~ v_\\infty = \\left[\\frac{M_\\oplus G}{a}\\right]^{1/2}, \\\\ v_F &=& \\left[v_\\infty^2 + \\frac{2M_\\oplus G}{r_F}\\right]^{1/2}, ~~~~~~~~~~~~~~ b= \\frac{M_\\oplus G}{v_\\infty^2 \\tan(\\Theta/2)}. \\eea In principle three of these parameters, say $v_\\infty$, $\\epsilon$, and $r_F$, determine the trajectory.  Therefore, one might consider looking for phenomenological patterns using such variables.  Another consideration is the largest non-Earth perturbation, which comes from the Moon. (Jupiter has a larger potential contribution but it is the potential gradient (force) that is important.)  In Figure \\ref{moon} we show the potential energies from the Moon during the GEGA1, NEAR, Cassini, Rosetta, and Messenger flybys.  The Moon's positional dependence could be important, as it is part of a separate three-body problem (Sun, Earth, Moon), making the actual spacecraft motion closer to a restricted four-body problem. Also, in Figure \\ref{sun} we show the potential energies from the Sun during the respective Earth flybys.  These are large but do not change much during the flybys.    \\begin{figure}[h!] \\begin{flushleft} \\noindent \\psfig{figure=gEGA-VM.eps,width=1.75in}% \\hskip 5pt \\psfig{figure=nEGA-VM.eps,width=1.75in}% \\hskip 5pt \\psfig{figure=cEGA-VM.eps,width=1.75in}% \\end{flushleft} \\begin{flushleft} \\hskip 1in \\noindent \\psfig{figure=rEGA-VM.eps,width=1.75in}% \\hskip 20pt \\psfig{figure=mEGA-VM.eps,width=1.75in}% \\end{flushleft} \\vskip -10pt \\caption{{\\small In the solar barycenter system,  the potential energies (all per unit mass) from the Moon during the (in order) GEGA1, NEAR, Cassini, Rosetta, and Messenger flybys in units of (km/s)$^2$ vs. time from periapsis in hours.} \\label{moon}} \\end{figure}   \\vskip .25in   \\begin{figure}[h!] \\begin{flushleft} \\noindent \\psfig{figure=gEGA-VS.eps,width=1.75in}% \\hskip 5pt \\psfig{figure=nEGA-VS.eps,width=1.75in}% \\hskip 5pt \\psfig{figure=cEGA-VS.eps,width=1.75in}% \\end{flushleft} \\begin{flushleft} \\hskip 1in \\noindent \\psfig{figure=rEGA-VS.eps,width=1.75in}% \\hskip 20pt \\psfig{figure=mEGA-VS.eps,width=1.75in}% \\end{flushleft} \\vskip -10pt \\caption{{\\small In the solar barycenter system,  the potential energies (all per unit mass)from the Sun during the (in order) GEGA1, NEAR, Cassini, Rosetta, and Messenger flybys in units of (km/s)$^2$ vs. time from periapsis in hours.} \\label{sun}} \\end{figure}     The data points currently available on possible flyby anomalies are limited to the two Pioneer flybys of a giant planet and to the Earth flybys of Table \\ref{haletable}. Unfortunately, the flybys of large satellites orbiting Jupiter and Saturn are not reliable because of uncertainties in their gravitational fields, as evidenced by Ganymede \\cite{anderson2004, palguta}. In addition, the detection of any small flyby anomaly, with respect to the absolute scale of the velocity in the flyby trajectory, requires a spacecraft trajectory relatively free of systematic accelerations. These accelerations can be generated by onboard systems, by non-conservative forces such as atmospheric drag, or by conservative forces from uncertainties in the planet's gravitational field.  Further, the radio Doppler data must be of high quality and referenced to ground-based atomic frequency standards by means of coherent two-way tracking. This limits the number of candidates for further flyby tests, whether of the Earth or other planets. However, given the importance of the gravity-assist technique for the conservation of rocket fuel, more flybys will occur in the future. Now that possible anomalies have been detected, it is just a matter of waiting for opportunities and then doing the necessary tweaking of the flyby orbit, the onboard systems, and the Doppler tracking in order to optimize the possibility for another measurement of an effect.  In summary, we have presented a description and brief discussion of the physics of the energy-transfer process that occurs during planetary flybys. We have also presented a series of intriguing real-world results associated with flybys that belie our current understanding of the underlying physics.  It is hoped that further study, which we encourage, can reconcile this situation."
    },
    "0608/astro-ph0608278_arXiv.txt": {
        "abstract": "We develop a first-order approximation method for the influence of spin on the motion of extended spinning test masses in a gravitational field. This approach is illustrated for approximately circular equatorial motion in the exterior Kerr spacetime. In this case, the analytic results for the first-order approximation are compared to the numerical integration of the exact system and the limitations of the first-order results are pointed out. Furthermore, we employ our analytic results to illustrate the gravitomagnetic clock effect for spinning particles. ",
        "introduction": "\\label{Sec1}  Imagine the motion of an extended spinning body in the exterior vacuum region of an astronomical source. Let $T^{\\mu \\nu}$ be the energy-momentum of the extended test mass; then, the motion is governed by the four equations that are given by the dynamical law $T^{\\mu \\nu}{}_{; \\, \\nu} = 0$. Consider a representative point inside the extended test body (the ``center of mass'') such that $u_s^\\mu = dx^\\mu/d\\sg$ is its four-velocity vector and $\\sg$ is the proper time along its worldline. The equations of motion can then be expressed relative to the chosen worldline. This has been accomplished in an elegant manner by Dixon \\cite{Dixon1}. Using Synge's world function \\cite{Synge}, Dixon has defined the infinite set of multipole moments of $T_{\\mu \\nu}$ in a way that is qualitatively similar to the standard nonrelativistic theory and thereby expressed the equations of motion as \\cite{Dixon1} \\be {DP^\\mu \\over d\\sg} & = & - {1 \\over 2} \\, R^\\mu{}_{\\nu \\alpha \\beta} \\, u_s^\\nu \\, S^{\\alpha \\beta} + {\\cal F}^\\mu, \\label{Dixon-p} \\nl {DS^{\\mu \\nu} \\over d\\sg} & = & P^\\mu \\, u_s^\\nu - P^\\nu \\, u_s^\\mu + {\\cal T}^{\\mu \\nu}. \\label{Dixon-s} \\ee Here, the momentum vector $P^\\mu$ and the spin tensor $S^{\\mu \\nu}$ are the first two moments of $T^{\\mu \\nu}$, while the Dixon force ${\\cal F}^\\mu$ and torque ${\\cal T}^{\\mu \\nu}$ are given in terms of the quadrupole and higher moments of the test body. It is necessary to add a supplementary condition to equations (\\ref{Dixon-p})--(\\ref{Dixon-s}) in order to fix an appropriate center-of-mass trajectory. It turns out that for an extended body the appropriate condition is \\cite{Dixon2} \\be S^{\\mu \\nu} \\, P_\\nu & = & 0. \\label{spin-condition} \\ee The work of Dixon provides a natural generalization of the previous results of Mathisson \\cite{Mathisson} and Papapetrou \\cite{Papapetrou} that were limited to pole-dipole particles.  In the present paper, we are interested in the dynamics of an extended body of mass $m$ and spin $s$ in the field of a rotating mass $M \\gg m$; therefore, we consider equations (\\ref{Dixon-p})--(\\ref{spin-condition}) without the quadrupole and higher moments, i.e., ${\\cal F}^\\mu = 0$ and ${\\cal T}^{\\mu \\nu} = 0$. The resulting reduced Mathisson-Papapetrou-Dixon (``MPD'') equations then lead to natural definitions for the mass $m$ and spin $s$ of the extended test particle that are preserved throughout the motion. We assume that the metric tensor has signature $+2$ and we use units such that $c = G = 1$, unless otherwise specified. It then follows from Eqs.~(\\ref{Dixon-p})--(\\ref{spin-condition}) that in this case the mass of the extended particle is given by $m = \\lt(-P^\\mu \\, P_\\mu\\rt)^{1/2}$, which is a constant of the motion. The spin vector is then defined by \\be S_\\mu & = & -{1 \\over 2 \\, m} \\, \\epsilon_{\\mu \\nu \\rho \\sg} \\, P^\\nu \\, S^{\\rho \\sg}, \\qquad S^{\\mu \\nu} \\ = \\ {1 \\over m} \\, \\epsilon^{\\mu \\nu \\rho \\sg} \\, P_\\rho \\, S_\\sg, \\label{S-vec} \\ee where $\\epsilon_{\\mu \\nu \\rho \\sg}$ is the Levi-Civita tensor given by $\\epsilon_{\\mu \\nu \\rho \\sg} = \\sqrt{-g} \\, \\sg_{\\mu \\nu \\rho \\sg}$ in terms of the alternating symbol with $\\sg_{0123} = 1$. It follows from the reduced MPD equations that $s^2 = S_\\mu \\, S^\\mu = {1 \\over 2} \\, S_{\\mu \\nu} \\, S^{\\mu \\nu}$ is a constant of the motion. Moreover, the requirement that the dipole force be much smaller than the monopole force implies that $s \\ll m \\, r$, where $r$ is the distance between $m$ and $M$; that is, the M{\\o}ller radius of the extended test particle $\\lt(s/m\\rt)$ must be very small compared to its distance from the source.  In a recent paper \\cite{Chicone1}, an approximation scheme has been introduced to deal with this extended pole-dipole system in most astrophysical situations. This scheme is based on the circumstance that $m^{-1} \\, P^\\mu - u_s^\\mu$ is small and of order $\\lt(M/r\\rt)\\lt[s/(mr)\\rt]^2 \\ll 1$. In this approach, it then turns out that to first order in $s/(mr)$ one can assume that $P^\\mu$ and $u_s^\\mu$ are parallel, namely, $P^\\mu \\approx m \\, u_s^\\mu$; therefore, it follows from Eqs.~(\\ref{Dixon-p})--(\\ref{spin-condition}) in this case that \\be {Du_s^\\mu \\over d\\sg} & \\approx & - {1 \\over 2m} \\, R^\\mu{}_{\\nu \\alpha \\beta} \\, u_s^\\nu \\, S^{\\alpha \\beta}, \\qquad {DS^{\\mu \\nu} \\over d\\sg} \\ \\approx \\ 0, \\qquad S_{\\mu \\nu} \\, u_s^\\nu \\ \\approx \\ 0. \\label{Dixon-approx} \\ee The second section of Ref.~\\cite{Chicone1} should be consulted for the details of the derivation of these equations.  In the absence of spin, system (\\ref{Dixon-approx}) simply reduces to the geodesic equation for the motion of the center of mass of the spinless extended particle, as expected~\\cite{Ehlers}. In Sec.~\\ref{Sec2}, we solve system (\\ref{Dixon-approx}) assuming that in the zeroth order of approximation the worldline for a spinless particle is a timelike geodesic. The motion away from the geodesic is then of first order in $s/(mr)$ in our approach. This method is employed to study approximately circular motion near the equatorial plane of a Kerr source of mass $M$ and angular momentum $J = Ma$ in Sec.~\\ref{Sec3}. That is, we assume that in the absence of spin, the test mass follows a stable equatorial circular geodesic orbit in the exterior Kerr spacetime. Our approximate analytic results based on system (\\ref{Dixon-approx}) are compared with the numerical solution of the extended pole-dipole system, Eqs.~(\\ref{Dixon-p})--(\\ref{spin-condition}), in Kerr spacetime in Sec.~\\ref{Sec4}. Section~\\ref{Sec5} is devoted to the gravitomagnetic clock effect. A brief discussion of our results is contained in Sec.~\\ref{Sec6}. The appendices contain further developments of our main results. ",
        "conclusions": "\\label{Sec6}  We have developed in this paper a general first-order approximation scheme for the effect of spin on the motion of an extended spinning test particle in a gravitational field in accordance with the MPD equations that are based on the multipole expansion method. We have neglected the influence of quadrupole and higher moments on the motion. For the astrophysically interesting case of motion in the field of a central source, we must have $s/(mcr) \\ll 1$, where $s/(mc)$ is the M{\\o}ller radius of the extended test particle. Thus the influence of the spin of the particle on its motion is treated to linear order throughout. It turns out that in this approximation the spin is parallel transported along the nongeodesic path of the test particle. To illustrate our general approach, the motion of spinning test particles along nearly circular equatorial orbits in the exterior gravitational field of a Kerr source has been investigated in detail and the results of the first-order approximation scheme have been numerically compared with the corresponding solution of the full MPD equations. As expected, the results are essentially identical for $s/(mcr)$ sufficiently small compared to unity; moreover, we have described the deviations that appear when $s/(mcr)$ is not so small.  The astrophysical implications of our results have been briefly mentioned in Sec.~\\ref{Sec3}. For the motion of the Earth about the Sun, for instance, the spin-induced deviations from Keplerian motion are too small to be observationally significant at present; however, the spin-dependent terms could become important in binary pulsars \\cite{Iorio,Barker}. Furthermore, we have elucidated the gravitomagnetic clock effect for spinning particles in accordance with our first-order approximation method."
    },
    "0608/astro-ph0608422_arXiv.txt": {
        "abstract": "Sub-TeV gamma-ray emission from the North-West rim of the supernova remnant RX J0852.0$-$4622 was detected with the CANGAROO-II telescope and recently confirmed by the H.E.S.S.\\ group. In addition, the H.E.S.S.\\ data revealed a very wide (up to two degrees in diameter), shell-like profile of the gamma-ray emission. We have carried out CANGAROO-III observations in January and February 2005 with three telescopes and show here the results of three-fold coincidence data. We confirm the H.E.S.S.\\ results about the morphology and the energy spectrum, and find the energy spectrum in the NW-rim is consistent with that of the whole remnant. ",
        "introduction": "RX J0852.0$-$4622 (G266.2$-$1.2) is one of the few supernova remnants (SNRs) from which strong non-thermal X-ray emission has been detected \\citep{aschenbach,tsunemi,slane}. A distance of $\\sim$200\\,pc and an age of $\\sim$700~years \\citep{aschenbach99,tsunemi} would make RX J0852.0$-$4622 one of the closest young SNR \\citep[although see][for a dissenting view]{slane}. The non-thermal emission suggests particle acceleration to $\\sim$100\\,TeV, with the possibility of TeV gamma-ray emission from inverse Compton scattering off electrons, and/or $\\pi^0$ decay from proton interactions, with clear implications for the origin of cosmic rays in our galaxy \\citep{volk05}.  Sub-TeV gamma-rays from the northwest rim (NW rim) of RX J0852.0$-$4622 were detected with the 10\\,m CANGAROO-II imaging atmospheric Cherenkov telescope (IACT) \\citep{katagiri}. Gamma-ray emission was recently confirmed by the H.E.S.S.\\ group \\citep{hess0852}, who mapped the spatial distribution of gamma-ray emission over the whole remnant, using data from at least two of four telescopes in their array.  RX J0852.0$-$4622 is the second spatially resolved SNR at TeV gamma-ray energies \\citep{enomoto_nature,hess_nature}. The CANGAROO-II result indicated a soft energy spectrum $\\propto E^{-4.3^{+1.7}_{-4.4}}$, however, in contrast the H.E.S.S.\\ group reported a harder spectrum $E^{-2.1\\pm0.1\\pm0.2}$. The gamma-ray fluxes also appeared discrepant: at 1\\,TeV the CANGAROO-II flux level was 15\\% of the Crab nebula, whereas the H.E.S.S.\\ result was 100\\% of the Crab. The clear differences between these two reports are; \\begin{itemize} \\item CANGAROO-II was only sensitive to within 0.4$^\\circ$ of the X-ray maximum in the NW rim \\citep{aschenbach,tsunemi,slane} whereas H.E.S.S.\\ was sensitive to the whole remnant. \\item the CANGAROO-II observation used a single telescope whereas the H.E.S.S.\\ observations used stereo data from two or more telescopes. \\item the CANGAROO-II flux was derived based on the subtraction of a background derived from the surrounding area which included part of the remnant, whereas the H.E.S.S.\\ background was derived from an off-source region outside the remnant. The CANGAROO-II result, therefore, was dependent on the spatial distribution of gamma-rays inside the SNR. \\end{itemize} The H.E.S.S.\\ result (from stereo observations) showed very extended emission, up to one degree in radius from the center of the remnant. In this paper we present the results of CANGAROO-III stereoscopic observations and determine how much of the apparent discrepancies between these two earlier results arise from these differences. ",
        "conclusions": "We have observed the supernova remnant RX J0852.0$-$4622 with the CANGAROO-III stereoscopic imaging Cherenkov telescope. It is clear from the results for RX J0852.0$-$4622 that CANGAROO-III's stereoscopic performance is significantly improved from that of the single CANGAROO-II telescope. We have confirmed the H.E.S.S.\\ group's result about the morphology and the energy spectrum. The NW rim is the brightest feature. RX J0852.0$-$4622  has been confirmed to be one of the strongest TeV gamma-ray sources in the southern hemisphere and shows a clear shell-type structure at these energies. The energy spectrum around the NW-rim was measured to be consistent with that of the whole remnant, not in good agreement with the previous CANGAROO-II result. The difference can be partially explained by the deterioration of the hardware of the CANGAROO-II telescope."
    },
    "0608/astro-ph0608614_arXiv.txt": {
        "abstract": "\\vspace{-3mm}  We have recently shown that the 3-parameter density-profile model from Prugniel \\& Simien provides a better fit to simulated, galaxy- and cluster-sized, dark matter halos than an NFW-like model with arbitrary inner profile slope $\\gamma$ (Paper~I).  By construction, the parameters of the Prugniel-Simien model equate to those of the S\\'ersic $R^{1/n}$ function fitted to the projected distribution. Using the Prugniel-Simien model, we are therefore able to show that the location of simulated ($10^{12} M_{\\sun}$) galaxy-sized dark matter halos in the $\\langle \\mu \\rangle_{\\rm e}$--$\\log R_{\\rm e}$ diagram coincides with that of brightest cluster galaxies, i.e., the dark matter halos appear consistent with the Kormendy relation defined by luminous elliptical galaxies.  These objects are also seen to define the new, and equally important, relation $\\log(\\rho_{\\rm e}) = 0.5 - 2.5\\log(R_{\\rm e})$, in which $\\rho_{\\rm e}$ is the internal density at $r=R_{\\rm e}$.  Simulated ($10^{14.5} M_{\\sun}$) cluster-sized dark matter halos and the gas component of real galaxy clusters follow the relation $\\log(\\rho_{\\rm e}) = 2.5[1 - \\log(R_{\\rm e})]$. Given the shapes of the various density profiles, we are able to conclude that while dwarf elliptical galaxies and galaxy clusters can have dark matter halos with effective radii of comparable size to the effective radii of their baryonic component, luminous elliptical galaxies can not.  For increasingly large elliptical galaxies, with increasingly large profile shapes $n$, to be dark matter dominated at large radii requires dark matter halos with increasingly large effective radii compared to the effective radii of their stellar component. ",
        "introduction": "Although Jaffe (1983), Hernquist (1990), and Dehnen (1993) introduced their highly useful density-profile models to match the deprojected form of de Vaucouleurs' (1948) $R^{1/4}$ model, these models do not immediately yield the types of structural quantities measured nightly by observers. For example, observers typically refer to a galaxy's (projected) half-light radius and surface density at this radius.  On the other hand, modelers frequently use the Navarro, Frenk, \\& White (1995, hereafter NFW) model --- a modified version of Hernquist's model --- and report scale radii\\footnote{Due to the divergence of the mass of the NFW model, a half-light radii cannot be defined.}  and scale densities ($r_{-2}$ and $\\rho_{-2}$) where the slope of the (internal) density profile equals $-2$.  Or, typically, they might report on the `concentration', related to the ratio of $\\rho_{-2}$ with the average background density of the universe, or the ratio of $r_{-2}$ with the halo's virial radius.  Recently, alternatives to both the NFW model, and it's generalization with arbitrary inner profile slope $\\gamma$, have been shown to provide better fits to the density profiles of simulated dark matter halos (e.g., Merritt et al.\\ 2006, hereafter Paper~I, and references therein). Considering fits to galaxy- and cluster-sized dark matter halos built from hierarchical $\\Lambda$CDM simulations, and fits to dark matter halos constructed from cold spherical collapses, two density models stand out (Paper~I).  The first is Einasto's (1965) model; see Tenjes, Haud, \\& Einasto (1994) for a more recent application.  This model has the same functional form as S\\'ersic's (1963, 1968) model but is applied to internal (3D) density profiles rather than projected surface density profiles.  Although we will not be using Einasto's model in this paper, having studied it in Paper~I and Graham et al.\\ (2006, hereafter Paper~II), we note that it has recently been applied to halos by Navarro et al.\\ (2004), Diemand, Moore, \\& Stadel (2004b), and Merritt et al.\\ (2005).  The second density model, by Prugniel \\& Simien (1997), which {\\it is} used here, is of particular interest because it is defined using two of the three parameters contained in S\\'ersic's $R^{1/n}$ (light-profile) model.  Indeed, the Prugniel-Simien model was developed to match the deprojected form of S\\'ersic's model. Specifically, the (projected) effective radius $R_{\\rm e}$ and the (projected) profile shape $n$ are common to both.  The third parameter in Prugniel \\& Simien's model, $\\rho_{\\rm e}$, an internal density term at $r=R_{\\rm e}$, can readily be used to obtain the associated surface density, $\\mu_{\\rm e}$, at the projected radius $R=R_{\\rm e}$.  This model therefore allows one to directly compare the structural properties of simulated dark matter halos with the structure of real galaxies and clusters, and to do so using a better fitting function than the generalized NFW model. In addition, Terzi\\'c \\& Graham (2005) have already shown that the Prugniel-Simien model describes the deprojected light-profiles of real elliptical galaxies better than the Jaffe, Hernquist and Dehnen models.  For both real spheroidal stellar systems and simulated dark matter halos alike, we have been able to plot their location on diagrams of mass versus: profile shape, effective radius, effective surface density, and effective internal (3D) density.  In the latter three figures, the galaxy-sized dark matter halos are observed to be consistent with the high-mass extension of ordinary elliptical galaxies.  Moreover, the cluster-sized dark matter halos are seen to coincide with the location of real galaxy clusters in all four diagrams, and possibly define a trend in the $M-n$ and $M-\\mu_{\\rm e}$ diagrams.  This is suggestive that the same (mass-dependent) physical processes are in operation in the real and simulated systems.  While it remains unclear as to why the simulated galaxy-sized halos do not follow the mass versus profile shape ($M-n$) relation defined by real elliptical galaxies (although see Nipoti \\& Ciotti 2006), we have been able to use this observation to constrain the galaxy-to-halo size ratio. In order for (dwarf and giant) elliptical galaxies to be dark matter dominated at large radii, we show that this ratio must increase with galaxy mass.  Due to the use of the Prugniel-Simien density-profile model, we have also been able to show, for the first time, the location of simulated dark matter halos in the Kormendy (1977) diagram.  Within this plane of effective surface density versus effective radius, luminous elliptical galaxies are observed to follow the relation $\\mu_{\\rm e} \\sim 3\\log R_{\\rm e}$ (e.g., Hamabe \\& Kormendy 1987; Graham 1996; La Barbera et al.\\ 2004).  Intriguingly, the galaxy-sized dark matter halos also appear consistent with this relation. Furthermore, we present a new diagram, and relation, using the internal (3D) effective density. Not only do the galaxy-sized dark matter halos follow the same linear trend defined by luminous elliptical galaxies, but the real galaxy clusters and cluster-sized dark matter halos follow their own linear relation which is offset in density by two orders of magnitude.  Such scaling laws refelct important regulatory mechanisms that dictate the formation and growth of galaxies and their dark matter halos, and as such these relations provide both inisght into these processes and useful constraints for galaxy modelling.  In the following Section we present the mathematical form of the Prugniel-Simien model.  In Section~3 we introduce the simulated and real data sets which are subsequently compared with each other in Section~\\ref{results}.  Section~\\ref{mass} shows trends with mass while Section~\\ref{SecComp} presents the $\\langle\\mu\\rangle_{\\rm e}$--$\\log R_{\\rm e}$ and the $\\log\\rho_{\\rm e}$--$\\log R_{\\rm e}$ diagrams. A brief summary is given in Section~\\ref{SecSum}. ",
        "conclusions": "\\label{SecSum}  Simulated galaxy-sized dark matter halos appear largely consistent with the location of brightest cluster galaxies in the $\\langle\\mu\\rangle_{\\rm e}$--$\\log R_{\\rm e}$ plane. Indeed, the halos appear congruent with the Kormendy relation.  Interestingly, the galaxy-sized halos also appear to follow the same relation as luminous elliptical galaxies in the $\\log \\rho_{\\rm e}$--$\\log R_{\\rm e}$ plane, defining a new relation $\\log (\\rho_{\\rm e}[M_{\\sun} \\hskip4pt {\\rm pc}^{-3}]) = 0.5 - 2.5\\log (R_{\\rm e}[{\\rm kpc}])$. Using this information, coupled with knowledge of the stellar and density profile shapes, we are able to make statements about the relative effective radii of stellar-to-dark matter distributions in galaxies and clusters. Specifically, while large elliptical galaxies require a small ratio of their stellar-to-dark matter effective radii, dwarf galaxies and galaxy clusters could have a size ratio of unity or larger and still be dark matter dominated. The galaxy clusters and simulated cluster-sized dark matter halos appear to define a new relation given by $\\log (\\rho_{\\rm e}[M_{\\sun} \\hskip4pt {\\rm pc}^{-3}]) = 2.5 - 2.5\\log (R_{\\rm e}[{\\rm kpc}])$."
    },
    "0608/hep-ph0608321_arXiv.txt": {
        "abstract": "The Sun is a source of high energy neutrinos ($E \\gtrsim 10$~GeV) produced by cosmic ray interactions in the solar atmosphere. We study the impact of three-flavor oscillations (in vacuum and in matter) on solar atmosphere neutrinos, and calculate their observable fluxes at  Earth, as well as their event rates in a kilometer-scale detector in water or ice. We find that peculiar three-flavor oscillation effects in matter, which can occur in the energy range probed by solar atmosphere neutrinos, are significantly suppressed by averaging over the production region and over the neutrino and antineutrino components. In particular, we find that the relation between the neutrino fluxes at the Sun and at the Earth can be approximately expressed in terms of phase-averaged ``vacuum'' oscillations, dominated by a single mixing parameter (the angle $\\theta_{23}$). ",
        "introduction": "The Sun is a well-known source of low-energy electron neutrinos ($E\\lesssim 20$ MeV), continually produced in  its core by nuclear reactions. The study of this flux has been one of the most exciting enterprises in astroparticle physics, and has provided us with crucial evidence for $\\nu$ flavor oscillations (see, e.g., \\cite{Ba89,Fogli:2006fu} and references therein).  High-energy neutrinos and antineutrinos (with $E\\gtrsim 10$~GeV) are also continually produced by cosmic ray interactions in the solar atmosphere~\\cite{SSG91,Moskalenko:1993ke,Ingelman:1996mj, Hettlage:1999zr,Learned:2000sw}, which lead to the production of secondary mesons and their subsequent decay into $\\nu_e$, $\\overline\\nu_e$, $\\nu_\\mu$, and $\\overline\\nu_\\mu$, collectively referred to as ``solar atmosphere neutrinos'' (SA$\\nu$) hereafter. From pion production dominance,  at the source the SA$\\nu$ flavor ratio is expected to be $\\phi_e:\\phi_{\\mu}:\\phi_{\\tau}\\simeq 1:2:0$, with approximately equal fluxes of $\\nu$ and $\\overline\\nu$. Concerning the $\\san$ energy spectrum, the Monte Carlo simulation in \\cite{Ingelman:1996mj} finds a power-law decrease ($\\sim E^{-3}$) for $E\\gtrsim 100$~GeV. At lower energies, heliomagnetic effects are expected to make the spectrum less steep (as $\\sim E^{-2}$) \\cite{SSG91,Hettlage:1999zr}.  At the Earth, within the solid angle subtended by the Sun, the $\\san$ flux significantly exceeds the more familiar ``atmospheric neutrino'' flux generated by cosmic rays interacting with the Earth atmosphere. In fact, because of the different scale heights, the density at the first interaction point is lower in the solar atmosphere than in the terrestrial one, and therefore a larger fraction of mesons can decay into neutrinos instead of being absorbed. The expected $\\san$ signal can be quantified in terms of $O(10)$ events per year in a km$^3$ of water (or ice) above a threshold $E\\gtrsim 100$~GeV~\\cite{Ingelman:1996mj}, and might thus be detected by large-volume $\\nu$ telescopes like IceCube~\\cite{Halzen:2006mq}.   The characterization of the $\\san$ signal is an important goal, both in itself and  because it might represent a ``background'' in searches for signatures of possible weakly interacting massive particle (WIMP) annihilation in the Sun~\\cite{Srednicki:1986vj, Ellis:1992df,Bergstrom:1998xh,deGouvea:2000un,Crotty:2002mv,Cirelli:2005gh,Halzen:2005ar}. (Predicted WIMP event rates are rather model-dependent, ranging from $10^{-12}$ to $10^3$ events/year/km$^3$ \\protect\\cite{Halzen:2005ar}.)  Neutrino oscillations can significantly affect the $\\san$ flux, and must be taken into account in the calculation of the corresponding lepton event rates in detectors at the Earth. Earlier studies have stressed, e.g., the relevance of $\\nu_\\mu\\to\\nu_\\tau$ vacuum oscillations \\cite{Hettlage:1999zr}, which might produce a few $\\tau$ events  per year in a km$^3$ detector (for $E>100$~GeV), comparable to the $\\tau$ event rate in OPERA, the CERN-to-Gran Sasso long-baseline experiment~\\cite{Sirignano}. Although very challenging from an experimental viewpoint, the identification of $\\san$-induced $\\tau$ events in future $\\nu$ telescopes might thus supplement laboratory tests of $\\nu_\\mu\\to\\nu_\\tau$ appearance with their astrophysical counterpart. Given the potential importance of the $\\san$ signal detection, we think  useful to revisit the impact of flavor transitions (in vacuum and in matter) on solar atmosphere neutrinos, in the light of the most recent advances in $\\nu$ oscillation physics. In particular, taking into account updated values for the mass-mixing $\\nu$ parameters, we investigate the three-flavor matter effects associated to $\\san$ propagation in the Sun.   The plan of our work is as follows. In Sec.~\\ref{nufluxes} we describe the input $\\san$ fluxes at the Sun, basically taken from Refs.~\\cite{Ingelman:1996mj,Hettlage:1999zr}. In Sec.~\\ref{nuprob} we study $\\san$ oscillations in vacuum and matter within the standard $3\\nu$ mixing framework, and then calculate their flavor transition probabilities at  Earth by solving numerically the flavor evolution equations. In Sec.~\\ref{detection} we evaluate the impact of flavor oscillations on the observable events rates in a km$^3$ neutrino telescope. We summarize our results in Section~\\ref{conclusions}.  In a nutshell, we find that interesting, genuine three-flavor transitions in matter can take place during the $\\san$ propagation, but also that their impact is significantly suppressed by a ``conspiracy'' of various effects, including spatial and energy smearing, and $\\nu/\\overline\\nu$ indistinguishability. In practice, the final oscillation effects can be approximated by (phase-averaged) vacuum flavor transitions, dominated by the mixing angle $\\theta_{23}$. ",
        "conclusions": "During the last forty years, the Sun has been the most important source of astrophysical neutrinos, which have provided us with crucial evidence of neutrino oscillations. Neutrino astrophysics is now entering a new era,  where next generation neutrino telescopes will be capable to detect high-energy (extra)galactic astrophysical neutrino fluxes from new sources. These $\\nu$'s would be useful both to probe the properties of the sources~\\cite{Barenboim} and to test for new physics beyond the standard  three-neutrino mixing scenario~\\cite{decay}.  The same telescopes might be able to ``see'' the Sun in a new way, namely, as a source of high energy neutrinos, produced by cosmic ray interactions in the solar atmosphere ($\\san$). Even if only a handful of events per year may be detectable in a km$^3$ telescope, their theoretical fluxes are relatively under control, and thus a measurement might provide a high-energy, astrophysical probe of neutrino flavor mixing. Moreover, the $\\san$ flux represents an unavoidable background in searching possible neutrino fluxes induced by WIMP annhihilation in the Sun core. For these reasons, it appears worthwhile to characterize the $\\san$ flux detectable at the Earth, after neutrino oscillations have occurred in the solar matter and in vacuum.  In our work we have studied solar atmosphere neutrino oscillations in detail, in the light of the most recent determinations of the neutrino mass mixing parameters, and paying particular attention to matter effects (generally neglected a priori in this context). We have found that, in the energy range $E\\in[10,\\,1000]$~GeV relevant for detection, $\\san$ oscillations in matter have in principle a rich phenomenology, and can generate effects which depend on all $3\\nu$ mass-mixing parameters. Unfortunately, these effects are smeared out to a large extent by summing over the neutrino and antineutrino channel, and by integrating over the solar disk angle (as well as over energy). In practice, oscillations are dominated by the ``vacuum'' effect of the $\\theta_{23}$ mixing angle. Therefore, even if matter effects can and do occur,  the original flavor ratio $\\phi_{e}:\\phi_{\\mu}:\\phi_{\\tau}\\simeq 1:2:0$ at the source approximately yields the well-known ratio $\\phi^D_{e}:\\phi^D_{\\mu}:\\phi^D_{\\tau}\\simeq  1:1:1$ at the detector (for maximal 2-3 mixing), just as in the case of phase-averaged vacuum oscillations.  Further refinements may include improved estimates of the input $\\san$ fluxes and of their uncertainties. In recent years, significant progress has been made in determining the cosmic ray composition and in refining the interaction and cascade development models~\\cite{Stanev:2004ys}. More detailed solar atmosphere models \\cite{Gu97} and future gamma ray data might also help to constrain the $\\san$ production mechanism. In particular, observations in gamma rays may sharpen our knowledge of helio-magnetic effects \\cite{SSG91,Moskalenko:2006ta}. Earlier estimates of the unoscillated $\\san$ fluxes \\cite{SSG91,Moskalenko:1993ke,Ingelman:1996mj,Hettlage:1999zr} could thus be usefully revisited. This task will certainly become mandatory if a $\\san$ signal will be found in neutrino telescopes or, even better, if it will become a background for signatures of WIMP annihilation in the Sun.  \\medskip"
    },
    "0608/hep-th0608015_arXiv.txt": {
        "abstract": "We consider a modified gravity theory, $f(R)=R-a/R^n+bR^m$, in the metric formulation, which has been suggested to produce late time acceleration in the Universe, whilst satisfying local fifth-force constraints.  We investigate the parameter range for this theory, considering the regimes of early and late-time acceleration, Big Bang Nucleosynthesis and fifth-force constraints. We conclude that it is difficult to find a unique range of parameters for consistency of this theory. ",
        "introduction": "Various cosmological observations suggest that the universe is pervaded by a new form of energy, dubbed dark energy, giving rise to accelerated expansion at the present epoch \\cite{Riess:1998cb,Perlmutter:1998np,Tonry:2003zg}. These observations are a challenge for fundamental physics, since a well--motivated candidate for dark energy has to be found. The usual candidates include scalar fields or extra dimensions.  Recently, it has been suggested that instead of a new matter form, the cosmic accelerated expansion could be attributed to a modification of General Relativity itself, see e.g. \\cite{Deffayet:2001pu,Freese:2002sq,Dvali:2003rk,Nojiri:2003ni,Carroll:2003wy,Capozziello:2003tk, Nojiri:2003ft,Carroll:2004de,Abdalla:2004sw}. For reviews see e.g. \\cite{Easson:2004fq,Nojiri:2006ri,Trodden:2006qk}. The simplest models of modified gravity are theories in which the Ricci scalar $R$ in the Einstein--Hilbert Lagrangian is replaced by some (arbitrary) function of $R$. Examples of these models include inverse powers $f(R) = R^{-n}$ or $f(R) = R - a R^{-n}$. It has been shown that these models are equivalent to scalar tensor theories with vanishing Brans--Dicke parameter, thereby ruling out these models \\cite{Chiba:2003ir}. It has recently been argued that by adding terms like $bR^m$ ($m$ positive), the field can be made massive and hence short-ranged, allowing the theory to be made consistent with local constraints~\\cite{Nojiri:2003ft}.  However, past work considered only the vacuum of the theory, without taking into account the coupling to matter.  In this paper, we will consider theories with positive and negative curvature terms, taking the matter coupling into account. This is a non-trivial task and we have to make simplifying assumptions in order to proceed. The paper is organised as follows. In the next Section we formulate the theory both in the Jordan frame and in the Einstein frame. We will write down useful equations needed for later Sections. In Section \\ref{seclocal}, local constraints for the choice $f(R)=R+bR^m$ are considered, analytically and numerically for $m=2$ and numerically for more general cases. In Section \\ref{secDE} we discuss consequences for the model $f(R)=R-aR^{-n} + bR^m$. We will summarise our findings in Section \\ref{seccomp}. ",
        "conclusions": "\\label{secconc} In this paper, we have investigated the parameter range of the modified theory of gravity, $f(R)=R-a/R^n+bR^m$.  For the specific case in which $a=0$, we modelled fifth force experiments using the local Earth system and calculated the strength and range of such a force. We find that, assuming a Yukawa form for the potential, the force has $\\alpha>\\frac16$ and $\\lambda=\\frac1\\mu$ where $\\mu$ is approximated by Eqn (\\ref{eqnmass}).  This restricts the allowed parameter space for $b$ such that the interaction range is smaller than experimentally allowed ranges.  For the general model ($a\\neq0$), we assume that $a/R^n$ leads to late time acceleration, setting $a$ for this to be the case.  We considered the allowed range of $b$ required by the following regimes: early and late-time acceleration, Big Bang Nucleosynthesis, existence of minima and fifth-force constraints.  Our results show that it is very difficult to find a consistent parameter range to satisfy two or more regimes.  The tightest constraint seems to come from BBN.  If the curvature values are known today and during the BBN epoch, the parameter $b$ would be constrained to a high degree.  Comparison of this result with fifth force constraints can ultimately decide the viability of this theory. Of course, more complicated cases for $f(R)$ might be considered, which might evade the afore--mentioned problems, see e.g. the discussion in \\cite{Nojiri:2006gh}.  In future, one should also verify our results by solving the full field equations in the presence of an extended object, like the Earth, and study the stability of the solutions, similar to \\cite{Faraoni:2005ie,Cognola:2005sg}."
    },
    "0608/astro-ph0608458_arXiv.txt": {
        "abstract": "We show that the combination of a weak magnetic propeller and accretion disc resonances can effectively halt accretion in short period cataclysmic variables for large fractions of their lifetimes. This may help to explain the discrepancy between the observed and predicted orbital period distributions of cataclysmic variables at short periods. Orbital resonances cause the disc to become eccentric, allowing material to fall back onto the donor star or out of the system. A weak magnetic field on a rapidly spinning primary star propels disc material outwards, allowing it to access these resonances. Numerical and analytic calculations show that this state can be long lived ($\\sim10^{11}$\\,yr). This is because the magnetic propeller is required only to maintain access to the resonances, and not to push matter out of the Roche lobe, so that the spin down time-scale is much longer than for a classical propeller model. ",
        "introduction": "\\label{sec:int}  It is known that accretion disc resonances in cataclysmic variables can remove angular momentum from the outer accretion disc \\citep{whi88}. This occurs when an element of the disc orbits with a frequency which is resonant with that of the binary orbit, typically at a ratio of three to one. This 3:1 resonance drives the accretion disc to become eccentric and causes the superhump modulations which are the defining feature of the superoutbursts of cataclysmic variables. Disc resonances provide a second source of angular momentum transport which is responsible for increasing the mass transfer rate, driving ordinary outbursts to become superoutbursts \\citep*{osa89,tru01}.  In known systems, superhumps normally fade soon after outburst. This is because the tidal torques clear material from the outer disc until there is no material remaining at the resonant radius. Resonant effects cannot switch on again until a new outburst drives the outer disc beyond the resonant radius once more.  In this paper we consider the possibility that access to orbital resonances can be maintained almost indefinitely by the action of a weak magnetic propeller. It is known that rapidly spinning and highly magnetic primary stars can act as magnetic propellers in some cataclysmic variables. Indeed, in the case of AE~Aqr it is believed that the entire mass transfer stream is ejected by a particularly strong magnetic propeller \\citep*{wyn97} and the evolution of AE~Aqr \\citep*{sch04} and WZ~Sge \\citep*{las99} may well have been influenced by magnetic effects. Here we consider the possibility that a {\\it much weaker} magnetic propeller might drive the disc outwards as far as the resonant radius, and so maintain a reservoir of material at this radius. The resulting continuous deposition of angular momentum by the resonance would drive the growth of eccentric modes. Disc material would be therefore be likely to extend beyond the primary Roche lobe. Such material would then either escape the binary or fall back to the donor star, preventing accretion onto the primary star in either case.  We are motivated to study this mechanism because it may modify the evolution of cataclysmic variables, especially if this state can be long lived. This may well be the case, since a much weaker torque is required than in the AE~Aqr case. The mechanism may slow or even halt accretion onto the white dwarf by driving mass out of the primary Roche lobe. At smaller mass ratios the resonant radii move further into the Roche lobe of the primary star, and are more likely to reach the disc, so that any effect is likely to influence short period cataclysmic variables disproportionately. This is interesting because existing models are unsuccessful in reproducing the observed evolutionary behaviour of short period cataclysmic variables.  Binary evolution models predict a minimum orbital period for cataclysmic variables which is reached when the donor star is no longer able to shrink in response to mass loss. Models consistently place this minimum at around 65\\,min \\citep*[e.g.]{kol92,how97,kol99}, whereas the observed period distribution shows a sharp cut off at 78\\,min \\citep*[e.g.][]{rit98,tap02}. The same models also predict that the evolution of the orbital period should slow at short periods, leading to an accumulation of systems in a sharp {\\it period spike} at the minimum period. This feature is absent from the observed distribution, as shown in Fig.~\\ref{fig:klaus}.  A number of explanations have been proposed for these discrepancies, but none are entirely satisfactory. An obvious solution is to hide short period systems from view, but standard selection effects do not act sufficiently strongly to explain the differences \\citep*{pat98,kin02}. Another possibility is that an additional angular momentum loss mechanism could act to increase the minimum period \\citep{pat98} or a range of angular momentum loss rates could blunt the period spike \\citep{kin02}. However, \\citet{bar02} show that even the loss of all the angular momentum carried by the accretion stream cannot increase the minimum period sufficiently. They also argue that blunting the period spike by `smearing'  the angular momentum loss rates would require a great deal of fine tuning to be successful.  In this paper we consider how the effects of propeller-induced access to orbital resonances might prevent some short-period CVs from exhibiting outburst behaviour, and cause them to become fainter. In Section~\\ref{sec:bin} we discuss the model in more detail. In Sections~\\ref{sec:num} and \\ref{sec:res} we explore the parameter space for which the idea may be of interest using smoothed particle hydrodynamics calculations. Finally, in Section~\\ref{sec:dis}, we discuss how this mechanism might influence the evolution of short period cataclysmic variables.  \\begin{figure} \\begin{center} \\resizebox{80.0mm}{80.0mm}{ \\mbox{ \\includegraphics{klaus2.ps}}} \\end{center} \\caption{CV orbital period histogram with the cumulative frequency overplotted as a solid line. This plot is based on fig. 1 from \\citet{kin02}, but uses more recent data from then online version of \\citet{rit98}. The plot contains data for all 215 CVs with periods known to be between $60$ and $120 \\ {\\rm{ mins}}$, from a total of 569 with known periods in the catalogue.} \\label{fig:klaus} \\end{figure}  \\begin{figure} \\begin{center} \\resizebox{80.0mm}{80.0mm}{ \\mbox{ \\includegraphics{radii.ps}}} \\end{center} \\caption{Plot of resonant, Roche lobe, circularisation, and viscous disc radii about the primary star in a close binary as a function of mass ratio $q$.} \\label{fig:radii} \\end{figure}  \\begin{figure*} \\begin{center} \\resizebox{88.1mm}{88.1mm}{ \\mbox{ \\includegraphics{fig1a.ps}}} \\resizebox{88.1mm}{88.1mm}{ \\mbox{ \\includegraphics{fig1b.ps}}} \\end{center} \\caption{This figure illustrates the positions of the 2:1 and 3:1 resonances in binaries with mass ratios of $q=0.05$ in panel a, and $q=0.5$ in panel b. The primary Roche lobe is marked with a solid line while the secondary star, which is assumed to be Roche lobe filling, is coloured grey. The two dashed circles represent the resonant radii, the inner circle representing the 3:1 resonance and the outer line being the 2:1 resonance. The length unit is the binary separation $a$ and the diagram is centred on the centre of mass. The primary star is represented by a cross. In case a, both the resonant radii are within the Roche lobe and accessible to the disc, whereas neither is available in case b.} \\label{fig:res} \\end{figure*} ",
        "conclusions": ""
    },
    "0608/astro-ph0608491_arXiv.txt": {
        "abstract": "We compare the observed peaked iron abundance profiles for a small sample of groups and clusters with the predictions of a simple model involving the metal ejection from the brightest galaxy and the subsequent diffusion of metals by stochastic gas motions. Extending the analysis of Rebusco et al. (2005) we found that for 5 out of 8 objects in the sample an effective diffusion coefficient of the order of $10^{29}$ cm$^{2}$~s$^{-1}$ is needed.  For AWM4, Centaurus and AWM7 the results are different suggesting substantial intermittence in the process of metal spreading across the cluster. There is no obvious dependence of the diffusion coefficient on the mass of the system.  We also estimated the characteristic velocities and the spatial scales of the gas motions needed to balance the cooling losses by the dissipation of the same gas motions. A comparison of the derived spatial scales and the sizes of observed radio bubbles inflated in the ICM by a central active galactic nucleus (AGN) suggests that the AGN/ICM interaction makes an important (if not a dominant) contribution to the gas motions in the cluster cores. ",
        "introduction": "X-ray spectroscopy is widely used to determine the metallicity of the hot gas in galaxy clusters and groups. Using the high energy resolution of ASCA, \\textit{BeppoSAX}, \\textit{Chandra} and \\textit{XMM-Newton} the radial profiles of Fe, S, Si, Ca and other elements (e.g Buote et al. 2003b, Sanderson et al. 2003, Sun et al. 2003, Matsushita, Finoguenov \\& B\\\"ohringer 2002, Gastaldello\\& Molendi 2002, O'Sullivan et al. 2005, Fukazawa 1994, Tamura et al. 2001) have been derived for a sample of objects. Outside the cluster core regions the metallicity of the intracluster medium (ICM) is on average one-third of the solar one and it does not seem to evolve up to a redshift of about $1$ (e.g Mushotzky \\& Loewenstein 1997, Tozzi et al. 2003). This suggests an early enrichment by Type II supernovae (SNII; e.g. Finoguenov et al. 2002). The metallicities in the core regions demonstrate much more diversity. For the so-called 'non-cooling flow' clusters the metallicity and the surface brightness do not vary strongly across the core region. For another group of so-called 'cooling flow' clusters, the metallicity and the surface brightness are both strongly peaked at the center. The clusters from the latter group always have a very bright galaxy (BG) dwelling in their centers, which makes these galaxies a prime candidate for producing the peaked abundance profiles.  These peaked profiles were likely formed well after the cluster/group was assembled (e.g. B\\\"ohringer et al. 2004; De Grandi et al. 2004).  The relative abundances of different elements indicate that Type Ia supernovae (SNIa) explosions and stellar winds within the BG have played a major role in the formation of the central abundance excesses (e.g. Renzini et al. 1993, Finoguenov et al. 2002).  If the metals in the ICM are due to the galaxy stars then one would expect the abundance profiles to follow the BG light profile. However the observed metal profiles are much less steep than the light profiles, suggesting that there must be a mixing of the injected metals, a process which may help to diffuse them to larger radii (e.g.Churazov et al. 2003, Rebusco et al. 2005, Chandran 2005). Stochastic gas motions could provide such a mechanism of spreading metals through the ICM, provided that the characteristic velocities and the spatial scales are of the right order. Furthermore, the dissipation of the kinetic energy of the same gas motions could be an important source of energy for the rapidly cooling gas in the ``cooling flow'' clusters.  In Rebusco et al. 2005 we estimated the parameters (velocity and length scale) of stochastic gas motions in the core of Perseus cluster (A426) required to spread the metals ejected from the central galaxy.  We found that a diffusion coefficient of the order of $2\\times10^{29}$ cm$^2$~s$^{-1}$ is needed to explain the observed abundance profile and that for turbulent velocities of about $300$ km~s$^{-1}$ and eddies sizes of about $20$ kpc the dissipation of turbulent motions compensates for the gas cooling losses. We now extend the same model to a small sample of clusters and groups. The sources we analyzed have mean temperatures which vary by a factor of $\\sim 6$.           The structure of the paper is the following.  In section \\ref{sec_sample} we describe the sample of clusters and groups; in section \\ref{sec_model} the structure and the ingredients of the stochastic diffusion model are described. The results are discussed in section \\ref{sec_results}. The last section summarizes our findings.\\\\ We adopt a Hubble constant of $H_0=70$ km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_M=0.3$ and $\\Omega_\\Lambda=0.7$. ",
        "conclusions": "We show that the iron abundance profiles observed in a small sample of nearby clusters and groups with cool core are consistent with the assumption that metals are spreading through the ICM with effective diffusion coefficients of the order of $10^{29}$cm$^2$~s$^{-1}$. This value does not show any obvious trend with the mass or the ICM temperature of the system.  For characteristic velocities of the order of $300$ km/s and length scales of the order of $10$ kpc, the dissipation of the gas motions would approximately balance the gas cooling. When observations are available, our estimates of the turbulent length scales are consistent with the observed buoyant bubbles size, pointing at the central supermassive black hole as the likely origin of the core gas motions.  Two objects in our sample require very different diffusion coefficients: for AWM7 very small level of gas mixing is needed, while for AWM4 the mixing has to be very intense. It is possible that these two objects hint to the intermittence of the gas mixing and that they represent two different episodes of the cluster evolution: AWM7 is only weakly disturbed now (but it will likely be disturbed soon by the onset of AGN activity), while AWM4 has been already mixed recently. O'Sullivan et al. 2005, using the analogy with M87 (where an abundance gradient is present in spite of clear signs of a radio jets interaction with the ICM) suggested that it is unlikely that AGN activity is responsible for the absence of the abundance peak in AWM4. One cannot however completely exclude the possibility that in the past there was a period of truly violent AGN activity in AWM4, which not only removed the central abundance peak, but also destroyed the whole cooling flow.    \\begin{table*} \\centering \\begin{minipage}{177mm} \\begin{center} \\begin{tabular}{c |c c c |c c c|c} \\hline \\hline Name & $a_b/a_\\odot$ & $M_{Fe}/M_\\odot \\times 10^8$ & $r_m/kpc$ & $k$ & $SR/SNU$ & $t_{age}/Gyr$ & $D/[cm^2/s]\\times 10^{28}$\\\\ \\hline NGC 5044   &  $0.2$ & $1.5$ & $42$ & $1.1$ & $0.26$ & $8$ &$12.2^*$  \\\\ & \\# $0.3$ & $0.5$ & $36$ & $1.1$ & $0.15$ & $5$ & $8.7$ \\\\ & $0.4$ & $0.8$ & $31$ & $2$    & $0.06$ & $8$ & $9$ \\\\ NGC 1550   & \\# $0.3$ & $0.8$ & $67$ & $1.1$ & $0.15$ & $7.6$ & $12.6^*$ \\\\ M87            & \\# $0.2$ & $0.9$ & $46$ & $1.1$ & $0.15$ & $6.2$ & $8.3$ \\\\ & $0.3$ & $0.8$ & $29$ & $1.1$ & $0.24$ & $5$  & $8.8$ \\\\ AWM4         & \\# $0.2$ & $7.2$ & $83$ & $ 2$   & $0.33$ & $10$ & $45$\\\\ &  $0.3$ & $3.8$ & $64$ & $1.1$ & $0.34$ & $10$ & $37.5$ \\\\ & $0.4$ & $1.9$ & $55$ & $2$    & $0.30$ & $6.4$ & $25^*$\\\\ Centaurus   & \\# $0.3$ & $6.0$ & $47$ & $2$    & $0.33$ & $10$ & $2^*$ \\\\ & $0.4$ & $3.6$ & $37$ & $1.1$ & $0.39$ & $10$ & $1^*$\\\\ AWM7         & \\# $0.3$ & $7.6$ & $50$ & $1.7$ & $0.15$ & $10$ & $0.6^*$ \\\\ & $0.4$ & $3.4$ & $32$ & $1.1$ & $0.15$ & $8.7$ & $0.5$ \\\\ A1795         & $0.2$ & $16$ & $140$ & $2$   & $0.30$ & $8.4$ & $36$\\\\ & \\# $0.3$ & $2.9$ & $62$ & $1.1$ & $0.15$ & $5.4$ & $12.5$ \\\\ \\hline \\end{tabular} \\end{center} \\caption{Enrichment models:column (1) Name of the object (2) Total iron excess; the hash indicates the model that was used to produce the plots in Fig. 1 and 2 (3) Half mass radius (4) $k$ (5) Present-day SNIa rate (6) Source lifetime (For the definitions see the text)-(7) Effective diffusion coefficient: the models marked with the asterisk have an uncertainty due to the choice of the criterion to fix $D$ of a factor $\\%/\\times 2$, while for the other models the uncertainty is smaller.} \\end{minipage} \\end{table*}"
    },
    "0608/astro-ph0608034_arXiv.txt": {
        "abstract": "We consider the hypothesis of a limiting minimal curvature in gravity as a way to construct a class of theories exhibiting late-time cosmic acceleration. Guided by the minimal curvature conjecture (MCC) we are naturally lead to a set of scalar tensor theories in which the scalar is non-minimally coupled both to gravity and to the matter Lagrangian. The model is compared to the Lambda Cold Dark Matter concordance model and to the observational data using the ``gold\" SNeIa sample of Riess et.~al.~(2004). An excellent fit to the data is achieved. We present  a toy model designed to demonstrate that such a new, possibly fundamental, principle may be responsible for the recent period of cosmological acceleration. Observational constraints remain to be imposed on these models. ",
        "introduction": "One of the most profound discoveries of observational physics is that the universe is accelerating in its expansion~\\cite{Riess:1998cb,Perlmutter:1998np,Netterfield:2001yq,Halverson:2001yy,Tonry:2003zg,Bennett:2003bz}. There have been many attempts to explain this late-time acceleration, for example, a pure cosmological constant, dark energy associated with some new scalar field and modified gravitational theories, although all current models require some level of fine-tuning and none are considered to be a complete explanation~\\footnote{For a recent reviews on the subject see~\\cite{Nojiri:2006ri,Copeland:2006wr}.}. Whatever is responsible for the current acceleration may arise from some completely new physical principle. This is the possibility we consider in this paper. Our goal is to construct a toy model that represents a late-time accelerating Universe using a new, possibly fundamental, principle. As our guiding principle, we hypothesize the existence of a {\\it minimal curvature \\rm} scale in gravity.  In a Friedmann, Robertson-Walker (FRW) space-time, without cosmological constant $\\Lambda$ and with only standard matter sources such as dust and radiation, the universe will always decelerate as it expands. One way to avoid this is to add matter to the system that violates the strong energy condition (SEC). In a cosmological context this violation constitutes the addition of matter sources satisfying the equation of state $p \\le -1/3 \\rho$. A second possibility is to explicitly remove flat space-time as a solution to the theory. In this case the vacuum of the theory, which is approached at late times as the energy density in matter fields becomes more and more dilute, is not Minkowski space-time, but instead an accelerating Universe~\\cite{Carroll:2003wy,Carroll:2004de,Easson:2004fq}. To remove flat spacetime as a solution we hypothesize the existence of a minimal curvature in our underlying fundamental theory. The simplest example of this is, of course, to introduce a bare cosmological constant into General Relativity. However, in principle there may exist many other ways to achieve this result. Indeed, it appears that many accelerating cosmological models derived from modified gravity theories contain such a minimal curvature~\\cite{Nojiri:2003wx,Schuller:2004nn,Schuller:2004rn}.  The idea of a minimal curvature scale in gravity mirrors that of a maximal curvature scale. In the literature many authors have considered this possibility and used it to remove the curvature singularities of General Relativity by bounding curvature invariants from above at the level of the classical action~\\cite{markov}-\\cite{Easson:2003ia}. In the case of singularity removal, it is necessary to bound {\\it all} curvature invariants in order to cover all possible physical situations in which such a singularity may occur.  By contrast, in the case of a minimal curvature approached at late times in a homogeneous, isotropic universe, symmetry implies that it is only necessary to bound the Ricci scalar $R$ from below. Hence, unlike in the case of a maximal curvature hypothesis, we shall see that one may implement a minimal curvature by using a modified Brans-Dicke theory where the Brans-Dicke field couples non-minimally to the matter Lagrangian.  Within this context we demonstrate that the existence of the minimal curvature (MC) produces a Universe that evolves from a matter dominated period to an accelerating phase mimicking the $\\Lambda$-Cold-Dark-Matter ($\\Lambda$CDM) model.  We emphasize that the model presented here is only a \\it toy construction of the late Universe\\rm. The model is not intended to provide a consistent cosmology from the time of Big-Bang Nucleosynthesis (BBN) until today. It is unlikely that the precise model presented here is compatible with solar system experiments and the tight constraints on the time variation of Newton's constant. However, the model \\it does \\rm provide an example of how the postulated existence of a minimal curvature scale in gravity can provide a new mechanism to generate cosmological acceleration of the late Universe.  Furthermore, the model may capture features of a possibly more fundamental theory that admits a minimal curvature scale~\\footnote{An interesting example of a minimal curvature scale occurs in a certain classical limit of quantum gravity~\\cite{Schuller:2004nn,Schuller:2004rn}.}.  In Section~\\ref{sec:MC}, we describe the minimal curvature construction, first by using a toy example and then by using a class of modified Brans-Dicke theories. We solve the equations of motion for this example and demonstrate how the Universe evolves from a matter dominated phase to an accelerating period as the curvature approaches its minimal value. In Section~\\ref{sec:comp}, we compare the MC model with $\\Lambda$CDM and to the supernovae (SNeIa) ``gold\"  sample of~\\cite{Riess:2004nr}. Finally, we comment on the possibility of constructing more realistic models that satisfy the limiting curvature hypothesis and offer our conclusions and speculations in Section~\\ref{sec:conclusions}. In Appendix A, we provide a detailed analysis of the vacuum MC theory. In Appendix B, we construct an Einstein frame description of the vacuum theory and compare it to the MC vacuum. ",
        "conclusions": "\\label{sec:conclusions} We have shown that a period of late-time cosmic acceleration can follow directly from a simple minimal curvature conjecture (MCC). The model fits the SNeIa data exceptionally well. While the specific formulation considered here is only a toy construction, unlikely to be compatible with constraints from solar system and table top test of the equivalence principle, it may capture phenomenologically interesting features of a more fundamental theory that admits a limiting minimal curvature. Furthermore, the construction successfully demonstrates the possibility that a new fundamental physical principle  may ultimately be responsible for the recent period of cosmological acceleration. It is possible that experimentally viable models based on the minimal curvature conjecture exist. The search for such models within the context of scalar-Gauss-Bonnet gravity is currently underway. Despite the tight theoretical and experimental constraints on scalar-Gauss-Bonnet cosmologies~\\cite{Nojiri:2005jg,Cognola:2006eg,DeFelice:2006pg,Calcagni:2006ye,Koivisto:2006xf}, we remain optimistic that an experimentally and theoretically viable model based on the minimal curvature construction can be discovered."
    },
    "0608/astro-ph0608202_arXiv.txt": {
        "abstract": "{} { We investigate the cosmological growth of dark halos and follow the consequences of coeval growth for the accretion history of associated supermassive black holes.} {The Press-Schechter approximation is used to obtain an analytic expression for the mean rate of growth of dark matter halos. Dark halo accretion rates are compared with numerical work and the consequences for understanding AGN evolution are described. } {The mean accretion rate onto dark matter halos is shown to have a simple analytic form that agrees with previous numerical work and that may easily be calculated for a wide range of halo mass, redshift and cosmological parameters.  The result offers a significant improvement over published fitting formulae deduced from merger trees. We then consider the growth of associated supermassive black holes, and make a basic test of the simple hypothesis of ``Pure Coeval Evolution'' (PCE) in which, {\\em on average}, black hole growth tracks dark halo growth.  We demonstrate that both the absolute value of the integrated AGN bolometric luminosity density and its cosmological evolution derived from hard X-ray surveys are well-reproduced by PCE. Excellent agreement is found at $z \\ga 0.5$, although the observed luminosity density drops by a factor 2 compared with PCE by $z=0$: black hole growth appears to decouple from halo growth at low redshifts, and this may be related to the phenomenon of ``cosmic downsizing''. Overall, AGN evolution appears either to be caused by or to be closely linked to the slow-down in the growth of cosmic structure. We also discuss the mean Eddington ratio averaged over all galaxies, which is predicted to show strong evolution to higher values with redshift. } {} ",
        "introduction": "One of the great mysteries in the study of active galaxies and QSOs is the physical origin of the strong cosmological evolution in their space density at a given luminosity.  For many years it has appeared that, at least in broad terms, the evolution was best described as ``pure luminosity evolution'' in which QSOs appeared on average to have dimmed with time rather than changing in space density \\citep{marshall,boyle88,croom04}.  This is in accord with recent evidence that almost all galaxies with a massive spheroid component contain a supermassive black hole at the present epoch \\citep{magorrian,ferrarese00,gebhardt00,tremaine,onken}, the great majority of which must be largely inactive today.  If these black holes were luminous in a phase of accretion and growth at higher redshift then their luminosities must indeed have declined with time.  Yet this picture is not easily in accord with modern ideas of galaxy formation, since if the matter content of the universe is dominated by cold dark matter we expect galaxies to grow hierarchically, and we expect the mass function of black holes at the centres of galaxies to increase with time - what then is the mechanism that allows the mass function to increase but the luminosity to decrease with cosmic epoch?  One obvious possibility is that black holes may increase in mass but with a decreasing mean accretion rate.  This picture has been incorporated into models such as those of \\citet{haehnelt93} and \\citet{kauffmann} in which simple parameterisations of such cosmological evolution in accretion rate were allowed, and it was shown that with a suitable choice of parameters the evolution in the QSO optical luminosity function could be matched.  The physical origin of this variation in accretion rate has not yet been determined, however, and explanations range from a systematic depletion in available gaseous material with cosmic epoch to long-term variations in the accretion process itself:  these latter explanations would of course only work if black holes were formed at high redshift and the luminosity evolution that has been observed were a reflection of the variation in accretion history of individual black holes.  More recently however it has become clear that optically luminous QSOs are a short-lived phenomenon, at least compared with the Hubble time. Measurement of the clustering of QSOs in the 2dF QSO Redshift Survey (2QZ: \\citealt{croom05}) has shown that their clustering amplitude does not increase with time as would be expected if QSOs were long-lived, given that cosmic structure is expected still to be growing at $z \\sim 2$ in a Universe with a cosmological matter density parameter $\\Omega_M \\sim 0.3$.  This implies that we must look for a universal variation in accretion rate rather than anything intrinsic to an individual accretion ``event''.  The simple picture of QSO luminosity evolution itself has now been shown to be more complicated when active galaxies are selected at X-ray wavelengths. First, it now seems likely that the hard X-ray background is produced by X-rays emitted from active galaxies, but in that case a large fraction of those active galaxies must be highly absorbed, with equivalent X-ray absorption column densities in neutral hydrogen $> 10^{23}$\\,cm$^{-2}$ \\citep{risaliti,comastri,ueda}.  These active galaxies would be likely to be optically obscured also. A second significant discovery from X-ray surveys however has been that low luminosity X-ray selected active galaxies, that would be classified as Seyfert galaxies were they optically selected, display cosmological evolution in their space density that appears to have a maximum at lower redshift than their high luminosity QSO counterparts \\citep{steffen,cowie,ueda,zheng,barger}.  It may be to some extent that the identification of these weak active galaxies may be less complete at higher redshifts, but early indications are that there is a significant shift of the space density maximum to lower redshifts with decreasing luminosity.  Hence it appears that there is much still to understand both observationally and theoretically about the evolution of active galaxies.  Nonetheless, the basic inference that there must have been evolution in accretion rate seems inescapable.  In this paper we test the extent to which the growth of black holes may be linked to the growth of galaxies, and argue that cosmological evolution in accretion rate can be understood primarily as arising from the cosmological evolution in the rate of accretion of matter onto galaxies. In section 3 we present a new calculation of the accretion rate onto dark halos using the extended Press-Schechter approximation. We find that the analytic expression agrees well with both our own and previous numerical estimation of the accretion rate calculated from merger trees.  The analytic approach here circumvents the need for such numerical estimation, and we also find that it provides mathematically better behaviour than fitting functions to the numerical results that have previously been proposed in the literature.  In section 4 we argue that coeval growth of black holes with their associated galaxies and dark halos leads to a simple expression for the total mass accretion rate onto all black holes. We calculate the expected integrated luminosity density arising from accretion onto black holes and show that it agrees remarkably well with observation. When expressed as the Eddington ratio (the ratio of the actual mass accretion rate to the rate required to attain the Eddington limiting luminosity) it shows significant cosmological evolution almost independent of halo mass and without depending strongly on the choice of cosmological parameters. ",
        "conclusions": "The results from this work may be summarised as follows. \\begin{enumerate} \\setlength\\itemsep{0em}  \\item Press-Schechter theory may be used to calculate analytically the mean mass accretion rate of dark matter halos.  The result of the calculation is in good agreement with previous numerical work but offers a significant improvement over previously-published fits to numerical simulations. The mass accretion rate depends almost linearly on mass but is very insensitive to choice of cosmological parameters within the range normally considered. We present fitting formulae to allow easy calculation of the accretion rate.  \\item It seems likely that the deduced strong decline in halo mass accretion will have an effect not only on the build-up of dark halos but also on the baryonic structures they contain.  There may therefore be a close link between the cosmological evolution in star formation rate, in AGN accretion and dark halo accretion.  \\item We have investigated arguably the simplest hypothesis one could make for black hole growth, ``Pure Coeval Evolution'' (PCE), which postulates that, {\\em on average}, supermassive black holes have growth that tracks the growth of their parent dark halos.  In this case it is straightforward to calculate the expected mass accretion rate onto black holes and, estimating the local black hole mass density, we predict a value for the integrated AGN luminosity density in remarkable agreement at $z \\ga 0.5$ with the value deduced from hard X-ray surveys.  \\item The evolution in integrated AGN luminosity density at $z > 0.5$ is well-matched by this model if the integrated black hole mass density tracks the mass density in massive $M_{\\rm H}>10^{11.5}\\msun$ halos, consistent with the value obtained by inferring the clustering bias of QSOs in the 2dF QSO redshift survey \\citep{croom05}  \\item At $z<0.5$ the observed luminosity density falls off faster than predicted, although even by $z=0$ the overprediction is only a factor two.  It seems that in the recent universe the black hole growth has started to decouple from dark halo growth, although a decrease in mean radiative efficiency associated with an increasing prevalence of ADAFs may also occur.  \\item Expressed in terms of the Eddington ratio, at sufficiently high redshifts ($z \\ga 4$), halos would have black holes accreting at a significant fraction of the Eddington rate - we expect all massive galaxies to contain an active AGN.  At higher redshifts still, the nominal Eddington ratio could greatly exceed unity.  The response of the black hole would be to limit its growth to the Salpeter rate, but such systems might produce significant outflows \\citep{king,kingpounds} which in turn would produce a significant feedback effect on their hosts.  It is also likely that in this phase of black hole growth, growth by mergers between halos is faster than growth by accretion.  Merger-dominated growth is one way of avoiding the problem of high-mass black holes at high redshift that have had insufficient time to grow at the Salpeter rate from solar- or intermediate-mass black hole progenitors \\citep{willott}. \\end{enumerate}  \\noindent Further work will model the AGN luminosity function and the contribution to the X-ray background (Babi\\'{c} et al. in preparation)."
    },
    "0608/astro-ph0608216_arXiv.txt": {
        "abstract": "A census of classical T Tauri stars and Herbig Ae/Be stars has been performed around the Orion-Eridanus Superbubble which is ionized and created by the Ori~OB1 association.  This sample is used to study the spatial distribution of newborn stars, hence the recent star formation sequence, in the region that includes two giant molecular clouds (Orion~A and B) and additional smaller clouds (NGC~2149, GN~05.51.4, VdB~64, the Crossbones, the Northern Filament, LDN\\,1551, LDN\\,1558, and LDN\\,1563).  Most of the molecular clouds are located on the border of the Superbubble, and associated with H$\\alpha$ filaments and star formation activity, except the Northern Filament which is probably located outside the Superbubble.  This suggests that while star formation progresses from the oldest Ori\\,OB1a subgroup to 1b, 1c and 1d, the Superbubble compresses and initiates starbirth in clouds such as NGC\\,2149, GN\\,05.51.4, VdB\\,64, and the Crossbones, which are located more than one hundred pc away from the center of the Superbubble, and even in clouds some two hundred pc away, i.e., in LDN\\,1551, LDN\\,1558, and LDN\\,1563.  A superbubble appears to have potentially a long-range influence in triggering next-generation star formation in an OB association. ",
        "introduction": "%  Ori~OB1 is one of the nearest OB associations \\citep{dez99}, harboring some 10 O-type stars \\citep{hum78}.  \\citet{bla64} divides the Ori~OB1 association into four subgroups, Ori~OB1a, 1b, 1c, and 1d, with 1a being the oldest and 1d (the Orion Nebula) being the youngest \\citep{bro94,war78,bla91}. From the $Hipparcos$ catalogue, \\citet{dez99} determine distances of 336$\\pm$16, 473$\\pm$33, and 506$\\pm$37 pc to Ori~OB1a, 1b, and 1c, respectively.  In Ori~OB1, most of the current star formation activity is located in, or close to, the two giant molecular clouds \\citep{wal00,meg05,her07a,bri07a}, Orion~A and B, which have been surveyed in $^{12}$CO by \\citet{mad86} and \\citet{wil05}.  Together with many smaller molecular clouds, the entire region is called the Orion-Monoceros complex \\citep{mad86,wil05}.  As the Ori~OB1 association evolved, the ionization photons and stellar winds from massive stars and supernova ejectra produced the Orion-Eridanus Superbubble \\citep{bro95}.  The Orion-Eridanus Superbubble has been studied in many wavelengths \\citep[reviews by][]{hei99}, e.g., in \\ion{H}{1} \\citep{hei76,bro95}, $IRAS$~100~$\\micron$ \\citep{bur93}, H$\\alpha$ \\citep{siv74,rey79}, X rays \\citep{bur93,sno95,guo95}, and $\\gamma$ rays \\citep{par98}.  Even though the Orion-Eridanus Superbubble is one of the most explored superbubbles, the overall star formation related to the Superbubble is not well investigated, simply because of its large sky coverage. It is indeed a great challenge to identify low-mass pre-main-sequence (PMS) stars in such a large region, so a wide-field survey is needed.  \\citet{wal00} and \\citet{bri07a} summarize different methods to search for low-mass (0.1--1~M$_{\\sun}$) PMS stars in OB associations, e.g., by objective-prisms, X rays, proper motions, and variability surveys.  Each method has its own advantage and drawback and here we summarize briefly each of the methods that have been applied to the Orion region.  An objective-prism imaging finds emission-line stars as young star candidates.  A large number of candidate PMS stars have been recognized by this method, e.g., in Orion \\citep{wir89,wir91,wir93,kog89} by the Kiso H$\\alpha$ survey.  The total number of H$\\alpha$ emission stars is 1157 over an area of 150 square degrees \\citep{wir93} to a limiting magnitude of V=15.  However, an H$\\alpha$ survey detects PMS stars only with strong H$\\alpha$ emission, i.e., classical T Tauri stars (CTTSs) and Herbig Ae/Be (HAeBe) stars, but is not sensitive to weak-line T Tauri stars (WTTSs).  An object-prism sample may also be contaminated considerably by late-type dwarfs that also show H$\\alpha$ in emission, namely the dMe stars.  Indeed, \\citet{bri01} conclude that only about 40\\% of the H$\\alpha$ emission stars in Orion OB1a found in the Kiso H$\\alpha$ survey are really PMS stars.  Because low-mass PMS stars (CTTSs and WTTSs) are X-ray sources \\citep{fei99}, the $ROSAT$ All Sky Survey (RASS), for example, has been used for large-scale searches for low-mass PMS stars.  $ROSAT$ is sensitive to soft X rays, so would find preferentially WTTSs.  The main contaminations in an RASS sample would be chromospherically active K and M dwarfs.  \\citet{ste95} study the density distribution of X-ray-selected PMS candidate stars in Ori OB1 (700 square degrees) and find significant surface density enhancements around Ori OB1a, OB1b, OB1c, and $\\lambda$ Ori.  \\citet{alc96} accomplish spectroscopic and photometric observations of the RASS X-ray sources in the Orion cloud complex and find 112 new WTTSs in about 450 square degrees, and follow up by broad- and narrow-band photometric and high-resolution spectroscopic observations \\citep{alc98,alc00}. \\citet{wal00} compare the spatial distributions of the H$\\alpha$ emission stars in the Kiso H$\\alpha$ survey and RASS X-ray sources in Ori\\,OB1.  The H$\\alpha$ emission line sources (CTTSs) are spatially closer to the Orion molecular clouds than the X-ray sources (WTTSs), a consequence of CTTSs being less evolved than WTTSs so are still in the proximity of their birthplace.  Stars in an OB association share common space motions, so proper-motion measurements can be used to identify probable members in an OB association with no limitation on stellar masses.  Because it is difficult to measure accurately proper motions of faint stars, this method is restricted to nearby OB associations and bright (massive and intermediate-mass) stars \\citep{dez99}.  One vulnerability of this technique is that the proper motions near the solar apex or antapex are small, so it is unsuitable to study regions near these two directions, as is the case for Ori~OB1 which is close to the solar antapex.  PMS stars are variable \\citep{her94}, so a multi-epoch photometric monitoring can be used to identify possible PMS stars in an OB association.  This method ensures that most PMS stars be found, but may also pick up other kinds of variables which are difficult to weed out on the basis of variability alone.  However, to conduct a multi-epoch survey of a large region is time-consuming. Using this method, \\citet{bri05} have identified 56 low-mass young stars in the Ori\\,OB1a subgroup, and 141 in Ori\\,OB1b.  Recently, \\citet{bri07b} apply a similar method and find nearly 200 low-mass PMS stars in the 25 Ori group.  \\citet{car01} use the southern 2MASS telescope and perform a 0.84$\\times$6 square degrees multi-epoch photometric observations near Trapezium.  They find 1235 near-infrared variable stars, most of which ($\\sim$93\\%) are probably associated with the Orion A molecular cloud.  They also show that the spatial distribution of stars with near-infrared excess traces the $^{13}$CO emission well.  None of the methods mentioned above are sensitive to those young stars embedded in clouds because of excessive extinction.  Recently, {\\it Spitzer Space Telescope} is used to investigate disk evolution in the $\\sigma$ Orionis cluster, 25 Orionis aggregate in Ori OB1a, and Ori OB1b \\citep{her07a,her07b}. \\citet{meg05} conduct a survey in the Orion region with {\\it Spitzer} to study star formation associated with the A and B molecular clouds.  In this work, we use 2MASS near-infrared data \\citep{cut03} to select CTTSs and HAeBe stars, so as to diagnose recent star formation related to the Superbubble.  Although near-infrared data miss sources deeply inside a cloud, most CTTSs and HAeBe stars can still be detected in the region we have investigated.  We describe in \\S 2 selection criteria of the young star sample, and follow-up spectroscopic observations.  In \\S 3, we outline the properties of some specific star-forming regions.  We discuss in \\S 4 the structure of the Superbubble, and how it influences the star formation history out to a distance of some 200~pc.  The conclusions are summarized in \\S 5. ",
        "conclusions": "%  We have conducted a census of classical T Tauri and Herbig Ae/Be stars selected by 2MASS colors in the vicinity of the Orion-Eridanus Superbubble produced by Ori~OB1.  Our study results in young star samples in two giant molecular clouds (Orion~A and B) and additional smaller molecular clouds, including NGC~2149, GN~05.51.4, VdB~64, the Crossbones, the Northern Filament, LDN~1551, LDN~1558, and LDN~1563.  Combined with H$\\alpha$ and CO images, we investigate the star formation history in this region.  The main results of our work are summarized as follows:  \\begin{enumerate}  \\item Clouds associated with H$\\alpha$ filaments are preferentially located on the border of the Orion-Eridanus Superbubble, created by UV photons from luminous stars in Ori~OB1.  Clouds outlined by H$\\alpha$ filaments exhibit recent star formation.  \\item The Northern Filament neither accompanies H$\\alpha$ filament, nor is associated with PMS stars. It is likely to be situated outside, hence out of reach of the compression by, the Superbubble to form stars.  \\item Star formation triggered by the expanding Orion-Eridanus Superbubble started from Ori~OB1a and propagated to other subgroups.  The Superbubble not only triggered starbirth in Orion~A and B, but also led to current star formation in NGC~2149, GN~05.51.4, VdB~64, and the Crossbones which are located some one hundred pc away from Ori~OB1a.  \\item The Superbubble likely also prompted star formation in clouds some two hundred pc away from Ori~OB1a, as witnessed in LDN\\,1551, LDN\\,1563, and LDN\\,1558.  \\item The triggered star formation scenario is supported by sequential star formation, ages of PMS stars in the clouds, and the expansion timescale of the Superbubble.  \\end{enumerate}"
    },
    "0608/astro-ph0608399_arXiv.txt": {
        "abstract": "name{\\color{red}\\bf Abstract} \\begin{abstract} {\\footnotesize The observation of several low energy events during SN1987A burst has an important weight in the attempts to learn about the properties of the supernova neutrinos, but is somewhat puzzling if compared with the expectations. In this work, we study the low energy events observed by Kamiokande-II, and consider in particular the possibility that a few of them are due to background. We focus our attention not only on the event 6, that falls below the energy threshold of 7.5 MeV, but also on the other four events  that at 1 $\\sigma$ fall below the energy threshold, namely the events 3,4,10,12. The {\\em a priori} expectations on the number of background events above threshold would suggest 1 or 2 background events only. The volume distribution of the Kamiokande-II events is not a uniform distribution at 3 $\\sigma$, that suggests the presence of background events close or at the border of the volume used for the analysis, including the events 3,4,10. Next we checked the expected energy distribution assuming that the signal is due to $\\bar{\\nu}_e p\\to e^+ n$ and that the average antineutrino energy is 14 MeV. The agreement with the observations is not perfect, but it is acceptable at the 11 \\% confidence level if we include the peak of low energy background events; otherwise, we face a 2.9~$\\sigma$ problem. The expected energy distribution implies that the evidence for supernova neutrinos is at 10 $\\sigma$ and that 1-3 background events are plausible. This conclusion does not change strongly when we model the time distribution of the signal, taking into account the presence of an initial luminous phase of neutrino emission. This suggests however that some of the early events (in particular, event 4) are due to supernova neutrinos and not to background. In summary, our comparison between the expectations and the data lead us to formulate the hypothesis that some of the observed low energy events are due to background and that some among them belong to a peculiar phase of emission, that could be further characterized by low energy neutrinos. Such an interpretation is mostly attractive, since it diminishes to a minimum the postulated number of background events and thus improves the agreement between the {\\em a priori} (model independent) and the {\\em a posteriori} (model dependent) expectations on the number of background events. We argue on these grounds that there is no significant disagreement between the average energy of the supernova neutrinos seen in Kamiokande-II and the conventional expectations. Alternative possibilities of interpretations are mentioned and briefly discussed. } ",
        "introduction": "Motivation and context \\label{abh}} On February 23, 1987, several experiments~\\cita{hirataPRL,imb,baksan,mb} contributed to begin the era of extragalactic neutrino astronomy. These observations had an enormous impact on astrophysics and on particle physics.\\footnote{In order to have an idea of how large was the impact it is sufficient to consider that, on March 2007, there are 172,000 (resp., 3,050) entries on `SN1987A' in {\\sc Google} (resp., in {\\sc Google scholar}). The {\\sc SPIRES} database lists 578 papers typing the same keyword; our first two references are those with more than 500 citations, whereas those that rank more than 100 citations include theoretical works on extra dimensions, CPT violation, axions, magnetic moments and exotic particles; several studies of the role of neutrino oscillations follow. Finally, the 20th anniversary of this observation was celebrated with conferences held in Moscow (Feb.~20-22, 2007), in Hawaii (Feb.~23-25, 2007) and Venice (Mar.~6-9, 2007)~\\cita{webs}.} A straightforward interpretation of these data is yet {\\em difficult}. This is due not only to the fact that we lack a firmly established theory of supernova explosion, but also to certain anomalous features of the data that have been understood, emphasized and analyzed in the course of the time. In particular, this is the case of the average energy deduced by the 12 events observed in Kamiokande-II (KII), which is half of that observed by IMB, and also lower then what expected in theoretical models for supernova neutrino emission, recently reviewed in~\\cita{keil}. The main goal of the present paper is to propose an interpretation of this apparent discrepancy.    The low energy feature of KII dataset is clearly reflected in the outcome of two recent analyses of these data, that explore certain possibilities that deviate strongly from theoretical expectations. The first one is by Mirizzi and Raffelt~\\cita{mira}, who describe the distribution in $\\bar{\\nu}_e$ energies as $E_{\\bar{\\nu}_e}^\\alpha \\exp[-(\\alpha+1) E_{\\bar{\\nu}_e}/\\langle E \\rangle]$ and find that the best fit of KII data is provided by $\\alpha\\sim 0$ ({\\em i.e.}, a monotonically decreasing distribution). The second one is the analysis of Lunardini~\\cita{luna}, who adopts a two-component distribution as suggested by three flavor oscillation scenarios and finds that a component with $\\langle E_{\\bar{\\nu}_e} \\rangle\\sim 5$~MeV permits to fit KII data better.\\footnote{In fact, the question of which is the energy distribution of an `average' core collapse supernova has important implications for future experiments: {\\em e.g.}, the analysis of SN1987A neutrinos of \\cita{luna} was used to argue that future search of relic supernova neutrinos could fail because the emitted neutrinos have much lower energy than expected.} This type of approaches is useful to emphasize features of the data like the excess of low energy events in KII dataset but one should recall that neither $\\langle E_{\\bar{\\nu}_e} \\rangle\\sim 5$ MeV nor $\\alpha\\sim 0$ are compatible with the current expectations for supernova neutrino emission, so that in a conservative analysis, these values should not be allowed.  Stated otherwise, there is a trend in modern studies of SN1987A to accept the opinion that KII data are incompatible with the present theoretical expectations and, consequently, to start from this position to investigate more or less radical departures from the conventional paradigm for neutrino emission. One could even think that the failures to obtain a theory of supernova explosion could motivate such an attitude. But the fact that we do not have yet a definitive theory of the explosion does not mean that all theoretical possibilities are {\\em a priori} equivalent. Furthermore, one should be careful in distinguishing between the problems of getting a theory of supernova explosion and the problem of knowing the distribution of the emitted neutrinos. In the only scenario that has been explored in some details till now, the so called ``delayed scenario'' \\cita{del1,del2}, a large amount of the neutrinos--up to 90~\\%--is emitted in a phase that {\\em follows} the explosion: thus, the two problems are to a certain extent independent (we will better analyze the connection in the following). Comparing~\\cita{keil} and~\\cita{bahcall} we note that the expected range of one crucial parameter, the average $\\bar{\\nu}_e$ energies or temperatures,  did not change much since 1989. What changed is the expectation on the temperature of the other antineutrinos and thus the impact of oscillations, that after~\\cita{keil} can be argued  to modify only slightly the observable $\\bar{\\nu}_e$ signal. To summarize, the main reasons why we are not convinced that we should abandon the conventional paradigm for neutrino emission are that: 1)~the expectations for neutrino emission seem to be stable; 2)~conversely, there is no convincing theoretical argument till now supporting an interpretation of the excess of low energy events in terms of low energy supernova neutrinos.      Also, on general statistical ground one should be aware of the risks of using the data--12 events in the case of KII--to infer the characteristics of the model, rather than asking the significance level at which the null-hypothesis (=the theoretical expectation) is ruled out. In principle it is always possible to obtain a perfect fit ($\\chi^2=0$) to the data declaring that the observed distribution coincides with the expected one, but only if one is ready to renounce to the previous knowledge. In short, we believe that a \\underline{conservative} discussion should address other questions: How severe is the deviation from the conventional theoretical expectations? How reliable is the indication of a large amount of low energy supernova neutrinos, inferred from KII observations? Is it possible to conceive other interpretations (more standard than the one proposed in \\cita{luna} and \\cita{mira}) of the observed low energy events?    This is why we would like to explore the possibility that some of the low energy events in KII are not due to $\\bar{\\nu}_e p\\to e^+ n$ interactions of supernova neutrinos as usually assumed. The possibility that a few events are due to elastic scattering has been recently reconsidered~\\cita{aldo}, finding that, although this cannot be excluded, it helps only marginally to explain the excess of low energy events in KII data set. This forced the authors of~\\cita{aldo} to admit that a few of these events are of a different origin, and possibly are due to background. This is a conservative position, since it is evident that the KII observations could be polluted by some background events: see figures~4, 5, 10, 11, 12 of the second publication of KII collaboration on SN1987A, {\\em i.e.}~\\cita{hirata}. Furthermore, in \\v{C}erenkov type detector it is not possible to know whether a single event is due to signal or to background. In our understanding, the main point of KII collaboration was just that it is not possible that all 12 events are due to background, which means that an observation of supernova neutrinos has been made (as we show in Sect.~\\ref{ed}, the {\\em a priori} knowledge of the signal reinforces significantly this evidence). It should be noted that the KII collaboration did not quantify the probability that the 12 events are due to background on individual basis; the only step in this direction was to exclude occasionally one specific event, occurred below the threshold for solar neutrinos events (more discussion on this issue in Sect.~\\ref{et}). In this note, we analyze quantitatively the possibility that some background events occurred in the KII dataset, extending and complementing the analysis of~\\cita{lambLoredo}. In particular, we focus our attention on the spatial and the energy distribution of the 12 candidate events.   We are aware of the problems of small numbers statistics, that there are many more papers on SN1987A neutrinos than events, and that there is the risk of running into `forensics' (quoting a witticism of John Beacom). Nevertheless, we should also recall that supernovae are rare events on human timescale, and that all we have at the moment is one observation of supernova neutrinos and some theoretical ideas to compare with. In other words, we feel that we have the duty to analyze all possible hints for anomalies and to extract as much information as we can from SN1987A neutrinos.     The plan of this paper is the following: in Sect.~\\ref{sep} we discuss the tools we have to separate signal from background events; in Sect.~\\ref{discu} we formulate definite hypotheses on the background and on the signal events, and show that the average energy of SN1987A neutrinos agrees much better with the expectations when we account for the presence of a few background events; in Sect.~\\ref{ddu} we comment on the approach we used and the results we obtained, and finally draw our conclusions. ",
        "conclusions": ""
    },
    "0608/astro-ph0608350_arXiv.txt": {
        "abstract": "Disk galaxies in a cluster of galaxies are moving in hot gas filling the cluster. Generally, they are moving at transonic or supersonic velocities. If ram-pressure stripping is insufficient to destroy the gas disk, the galaxies should be affected by the wind of the surrounding hot gas similar to an airfoil. In this paper, I consider whether the aerodynamic interaction can be strong enough to force a disk galaxy to deviate from the orbit that it would have been in. I find that while the lift force is not effective, the drag force could affect face-on disk galaxies in poor clusters on long timescales. ",
        "introduction": "The hot gas in clusters is often called the intracluster medium (ICM). The typical temperature and density of the ICM are $T_{\\rm ICM}\\sim 2$--10~keV and $\\rho_{\\rm ICM}\\sim 10^{-27}\\rm\\: g\\: cm^{-3}$, respectively \\citep{sar86}. While most galaxies are moving at transonic velocities ($V\\sim 1000\\rm\\: km\\: s^{-1}$) \\citep{sar86}, the galaxies that fall into a cluster from the outside have a Mach number of ${\\cal M}\\sim 2$--3 \\citep{ghi98}. So far, the interaction between disk galaxies and the surrounding ICM has been considered mostly in terms of `ram-pressure stripping' \\citep{gun72,fuj99,fuj99b,got03}. The disk of a disk galaxy is composed of stars and gas filling the disk. The latter is called interstellar medium (ISM), and is cold and dense ($T_{\\rm ISM}\\lesssim 10^4$~K and $\\rho_{\\rm ISM}\\gtrsim 10^{-24}\\rm\\: g\\: cm^{-3}$) compared with the ICM. The ISM prevents the penetration of the galaxy by the ICM. If the ram-pressure force of the ICM ($\\sim \\rho_{\\rm ICM}V^2$ per unit area) becomes larger than the gravitational restoring force of the galaxy, the ISM is stripped away from the galaxy \\citep{gun72}. Once the ISM is stripped, the ICM flows through the stellar disk of the galaxy without much resistance. Many numerical simulations have been performed to study this ram-pressure stripping \\citep{aba99,mor00,qui00,sch01,vol01,mar03,roe05a,roe05b}.  Another possible interaction between the ICM and ISM may be aerodynamic effects on a galaxy by the flow around the galaxy. In this paper, I discuss a few cases where the aerodynamic forces could operate on the galaxy motion, assuming that a disk galaxy is a thin solid disk for simplicity. Since the aerodynamic effects are available only when the disk ISM is not completely removed, I discuss only the cases where numerical simulations showed that the ram-pressure does not remove the ISM completely. Ram-pressure stripping is ineffective when the attack angle of a galaxy is small, and/or when the galaxy velocity is small. I consider the former in section~\\ref{sec:lift} and the latter in section~\\ref{sec:drag}. ",
        "conclusions": "I have shown that while the lift force on disk galaxies can be ignored, the drag force could not be ignored for face-on disk galaxies in poor clusters.  Equations~(\\ref{eq:lift}) and~(\\ref{eq:drag}) show that the strengths of the lift and drag forces are sensitive to the radius of the ISM disk, $r$. In the above estimates, gas supply from stars in the galaxy was not considered. If the gas ejected from the stars is added to the ISM, $r$ should become larger \\citep{gae87}. It is likely that active disk galaxies, in which large amount of gas is being released from the stars, have much larger $r$ and are more affected by the aerodynamic forces. It is also to be noted that if the velocity or pressure distribution of the ICM around a disk galaxy is measured, the aerodynamic forces can be constrained.  The observations of stripped ISM will be useful to obtain the velocity field around a disk galaxy \\citep{yos04,che06}. X-ray observations will tell us the pressure distribution around a galaxy \\citep{vik01}.  The models presented here are very simple.  In an actual galaxy, the situation would be rather delicate.  One example is how the lift or drag would operate on the collisionless stars and dark matter particles, as required to alter the orbit of the whole galaxy. The connecting force is gravity. If the gas disk is lifted or dragged off the stellar disk, and the center of the dark matter halo, then both will pull back on the gas disk. This would communicate the lift or drag force to all components. Too much separation, and the restoring gravity force will be weakened and the gas will be stripped, too little and nothing will happen. In order to study the process quantitatively, numerical simulations would be required.  \\vspace{5mm}  I am grateful to C. Struck for useful comments.  Y.~F.\\ was supported in part by a Grant-in-Aid from the Ministry of Education, Culture, Sports, Science, and Technology of Japan (17740182)."
    },
    "0608/astro-ph0608166_arXiv.txt": {
        "abstract": "We present near-infrared (nIR) and optical observations of the afterglow of GRB~030115.  Discovered in an infrared search at Kitt Peak 5 hours after the burst trigger, this afterglow is amongst the faintest observed in the R-band at an early epoch, and exhibits very red colors, with $R-K\\approx 6$.  The magnitude of the optical afterglow of GRB~030115 is fainter than many upper limits for other bursts, suggesting that without early nIR observations it would have been classified as a ``dark'' burst. Both the color and optical magnitude of the afterglow are likely due to dust extinction and indicate that at least some optical afterglows are very faint due to dust along the line of sight. Multicolor {\\it Hubble Space Telescope} observations were also taken of the host galaxy and the surrounding field. Photometric redshifts imply that the host, and a substantial number of faint galaxies in the field are at $z \\sim 2.5$.  The overdensity of galaxies is sufficiently great that GRB~030115 may have occurred in a rich high-redshift cluster. The host galaxy shows extremely red colors (R-K=5) and is the first GRB host to be classified as an Extreme Red Object (ERO). Some of the galaxies surrounding the host also show very red colors, while the majority of the cluster are much bluer, indicating ongoing unobscured star formation. As it is thought that much of high redshift star formation occurs in highly obscured environments it may be that GRB 030115 represents a transition object, between the relatively unobscured afterglows seen to date and a population which are very heavily extinguished, even in the nIR. ",
        "introduction": "Most ground-based GRB afterglow searches are being pursued first in optical wavelengths. These have relied on either comparing a newly acquired image to an archival image (normally the Digitized Sky Survey) and locating any new sources, or on performing multi-epoch observations and identifying any variable sources within the GRB error box. In many cases both approaches have been used. Such searches have resulted in the discovery of approximately 60 GRB afterglows with about 50\\% and 35\\% being discovered, respectively, by each of the methods, and $\\sim15\\%$ being discovered by a combination thereof. One GRB afterglow was found by a color-color technique, (Rhoads 2001; Gorosabel \\etal\\ 2002), exploiting the ability of multi-band observations to distinguish blackbody stellar spectra from the power-law spectra displayed by GRB afterglows.  The advent of small robotic telescopes that can slew automatically in response to GRB alert notices (e.g., ROTSE) and larger telescopes that can easily be remotely operated (e.g., the Palomar 48~inch) has increased the discovery rate of young and bright afterglows. For example, the afterglows of GRB~021004 and GRB~021211 were discovered only minutes after the GRB trigger (Fox 2002; Fox \\etal\\ 2003a,b). These rapid-response campaigns have also substantially decreased the fraction of bursts that have no identified optical counterparts: the so called ``dark bursts.'' Of the GRBs well localised so far by the Wide X-ray Monitor (WXM) on {\\it HETE-2,} the optical counterpart discovery percentage is $\\sim 60\\%$, while for bursts localized by the Soft X-Ray Camera (SXC) the discovery rate is $\\sim 90\\%$, (largely due to the smaller error boxes associated with SXC positions - see for example, Lamb \\etal\\ 2003). There are evidently possible selection effects: relatively bright and soft bursts are more easily localised by the SXC, and these smaller error boxes are easier to search to deep limits. Indeed, although {\\it Swift} now allows routine $\\sim 3$ arcminute error boxes to be routinely circulated its increased sensitivity means that it locates fainter bursts, and the recovery rate of optical afterglows is lower than for the SXC. Nevertheless the dark burst fraction does appear to have decreased significantly from the $\\geq 70\\%$ that was typical during the first few years of GRB observations, when slow and less accurate positions made it more difficult to locate afterglows, even for relatively bright bursts (see e.g. Fynbo \\etal\\ 2001; Lazzati \\etal 2002).   The spectral energy distribution (SED) of a GRB afterglow is best described by a set of connected power-laws with breaks due to synchrotron self-absorption, the peak energy of the electrons, and the time-dependent cooling of the fireball (Sari, Piran \\& Narayan 1998). In theory, knowledge of these break frequencies allows us to reconstruct details of the afterglow (as has been done for a number of well sampled GRB afterglows to-date, e.g., by Panaitescu \\& Kumar 2001). Theoretical models are also useful for studying GRBs with no optical afterglows, where observations in some bands (such as X-rays and radio) can be used to place limits on the expected optical flux.  In a prototypical dark burst, such as GRB~970828 (Groot \\etal\\ 1998; Djorgovski \\etal\\ 2001), extrapolation of the X-ray flux, assuming a synchrotron spectrum, led to the conclusion that the burst should have been seen in the optical.  Conversely, Fynbo \\etal\\ (2001) have shown that the classification of many bursts as dark may simply represent the failure to search either early or deep enough. This interpretation is supported by the detection of the faint afterglow of GRB~020124 (Berger \\etal\\ 2003) and the fast fading afterglow of GRB~021211 (Fox \\etal\\ 2003a; Crew \\etal\\ 2003), which would not have been detected but for early, deep observations. Although it is also true that the refinement of afterglow locations via X-ray observations does allow for very deep pointed observations, and thus enables faint afterglows to be located. This is now commonly the case for {\\it Swift} bursts.  For the relatively small number of bursts that are still apparently ``dark'' despite deep and early optical limits, a number of plausible explanations have been widely discussed including i) bursts which occur at high$-z$ where the Lyman break has moved through the optical bands ii) GRBs which originate from behind significant obscuring columns, significantly attenuating the optical light, iii) bursts which are intrinsically faint (i.e. are spectrally similar to bright) afterglows but are a factor 10-100 fainter and iv) fade very rapidly so that, at the time of the first optical observations they have faded below the detection limit. Observations of afterglows to date, including recent observations of {\\it Swift} bursts indicate that all of the above explanations may apply with some bursts originating from beyond z=5 (e.g. Jakobsson et al. 2005; Haislip et al. 2005; Kawai et al. 2005), a fraction being apparently obscured (e.g. De Pasquale et al. 2005; Watson et al. 2005). While the first two scenarios above have been used to motivate afterglow searches in the near-infrared, locating afterglows from bursts akin to the third and fourth options can be equally well done by early, deep optical observations.  As the connection between long-duration GRBs and supernovae is now secure (Hjorth \\etal\\ 2003; Stanek \\etal\\ 2003; Zeh, Klose \\& Hartmann 2004), we expect GRBs to point to locations of star formation throughout the Universe. The determination of which of the explanations above is the dominant origin of dark GRBs may thus shed light into the processes of star formation and its evolution. For example were many dark GRBs genuinely at high redshifts, they would point to strong star formation even at very early ages in the Universe, while if most are actually in obscured environments at lower redshifts they would support models whereby much of the star formation in the Universe is obscured. Moreover, the type of their host galaxies, their environments and relative frequencies may also indicate where much of the star formation occurs.  Here we present the nIR discovery of the afterglow of GRB~030115. This burst was heavily reddened with respect to a model power law SED and, in the optical, was the faintest GRB afterglow ever observed at similar epochs post-burst.  We show that its host galaxy lies at moderate redshift and has very red colors (R$-$K $\\sim 5$). It is therefore likely that GRB~030115 was reddened due to high local extinction. Furthermore there is a significant overdensity of galaxies at the redshift of the GRB host possibly associated with the latter; this putative galaxy cluster is a mix of young, blue, star forming systems and red systems which are either dusty starburst systems or contain an older stellar population. This cluster, if confirmed, is the first associated with a GRB, and one of very few known at $z>2$.  Throughout this paper we assume a $\\Lambda CDM$ cosmology with $\\Omega_{\\Lambda} = 0.73$, $\\Omega_M = 0.27$ and $H_0 = 72$~km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": "We have reported the discovery of the afterglow of GRB~030115. Originating from a (photometrically determined) redshift of $\\sim 2.5$, the afterglow was the faintest seen in the R-band after 24 hours. The decay rate was faster than the mean of most GRBs with an average decay rate over the first 24 hours of $\\alpha \\sim 1.5$.  The faintness of the afterglow in the optical is therefore likely a function of both this rapid decline and its extremely red spectrum.  Comparison of both the afterglow and host galaxy with those observed for other GRBs implies that the latter is the most reddened of any GRB host galaxy for which an afterglow has been seen. This is in contrast to some previous cases where red afterglows have been observed to lie in blue galaxies (e.g. GRB~000418, Klose \\etal\\ 2001; Gorosabel \\etal\\ 2003). However none of these afterglows is as extreme as that of GRB~030115, which has a flux density below the upper limits set for the large majority of dark GRBs at similar epochs. It may therefore be that some of the dark GRBs so far observed evolved in similar ways to GRB~030115 and were simply missed due to the lack of early IR observations. IR-robotic observations (e.g., REM, Liverpool Telescope) may be very succesful at locating such afterglows in the future and will therefore provide valuable statistics on the number of GRBs occuring in obscured galaxies.  In turn these GRB hosts  offer a useful tool in characterising the nature of embedded star formation. Indeed moderate resolution spectroscopy of these sources at early times may be one of the best means of providing a detailed study of the environments of such galaxies, which are too optically faint for normal absorption line studies.  GRB~030115 also lies in a potentially significant overdensity of galaxies at $z \\sim 2.5$. This association consists of a high ERO density, consisting of either dust enshrouded star formation or old, passive systems as well as bluer systems with active unobscured star formation. This association represents one of the richest associations found at $z>2$.  However, further observations are necessary to confirm the existence of this cluster. Thus far, we have found low-redshift GRBs in field galaxies, however, we expect that as we study GRBs at higher redshifts, we should begin to find them, like GRB~030115, in clusters or in far richer regions of the Universe directly associated with star formation."
    },
    "0608/astro-ph0608600_arXiv.txt": {
        "abstract": "The coalescence of a binary black hole system is one of the main sources of gravitational waves that present and future detectors will study.   Apart from the energy and angular momentum that these waves carry, for unequal-mass binaries there is also a net flux of linear momentum that implies a recoil velocity of the resulting final black hole in the opposite direction. Due to the relevance of this phenomenon in astrophysics, in particular for galaxy merger scenarios, there have been several attempts to estimate the magnitude of this velocity.  Since the main contribution to the recoil comes from the last orbit and plunge, an approximation valid at the last stage of coalescence is well motivated for this type of calculation.  In this paper, we present a computation of the recoil velocity based on the close-limit approximation scheme, which gives excellent results for head-on and grazing collisions of black holes when compared to full numerical relativistic calculations. We obtain a maximum recoil velocity of $\\sim 64\\,$ km/s for a symmetric mass ratio $\\eta = M^{}_1M^{}_2/(M^{}_1+M^{}_2)^2 \\sim 0.19$ and an initial proper separation of $4\\,M$, where $M$ is the total ADM mass of the system. This separation is the maximum at which the close-limit approximation is expected to provide accurate results.  Therefore, it cannot account for the contributions due to inspiral and initial merger.  If we supplement this estimate with PN calculations up to the innermost stable circular orbit, we obtain a lower bound for the recoil velocity, with a maximum around $84$ km/s. This is a lower bound because it neglects the initial merger phase.  We can however obtain a rough estimate by using PN methods or the close-limit approximation. Since both methods are known to overestimate the amount of radiation, we obtain in this way an upper bound for the recoil with maxima in the range of $220-265$ km/s. We also provide non-linear fits to these estimated upper and lower bounds. These estimates are subject to uncertainties related to issues such as the choice of initial data and higher effects in perturbation theory. Nonetheless, our estimates are consistent with previous results in the literature and suggest a narrower range of possible recoil velocities. ",
        "introduction": "The inspiral and merger of binary black holes systems is one of the most interesting sources of gravitational waves that both earth-based interferometric antennas (LIGO~\\cite{ligo}, VIRGO~\\cite{virgo}, GEO600~\\cite{geo} and TAMA~\\cite{tama}) and space-based ones (LISA~\\cite{lisa}) will detect. These waves carry both energy and momentum away from the system, leading to the adiabatic shrinking of the orbit, due to the former, and a recoil of the merged object by conservation of momentum, due to the latter. The magnitude of this recoil is of astrophysical importance because it determines whether the merged hole will be ejected from its host galaxy.   Possible observational evidence for such a recoil may be the observations of faint galaxies~\\cite{Binggeli:1998bj,Binggeli:2000bb} where the lack of a dense nucleus has been associated with the central black hole being ejected after merger~\\cite{Merritt:2004xa}. There is also evidence of an ejection of a supermassive black hole in ongoing galaxy mergers, either because of recoil or because of slingshot due to the presence of 3 or more supermassive black holes in the merger~\\cite{Haehnelt:2006hd}. The gravitational recoil has also been shown to have important consequences in hierarchical merging scenarios and the observable structure of galaxy nuclei. Recoil velocities of a few hundred km/s could be large when compared to escape velocities of dwarf galaxies, globular clusters and dark matter halos~\\cite{Haiman:2004ve,Merritt:2004xa}. For super-massive black holes at the centers of galaxies, a kick of this magnitude could potentially transfer energy to the stars in the nucleus~\\cite{Merritt:2004xa}. There are thus very important astrophysical aspects that can be refined or clarified with a better understanding of the black hole kicking process.   Mass distributions without symmetries that undergo gravitational collapse of any sort will exhibit momentum ejection and recoil of the center of mass of the remnant due to the strong emission of gravitational radiation~\\cite{Bonnor:1961br,Papapetrou:1962ap,Peres:1962ap,Bekenstein:1973jd,Redmount:1989rr}. Of particular interest is the case of unequal-mass black-hole binary systems. An intuitive picture of how the system gets a {\\em kick} after the merger is the following~\\cite{Wiseman:1992dv,Favata:2004wz}:  From the {\\em center of mass} point of view, the lighter black hole will move faster than the heavier one, and hence it will beam forward gravitational radiation stronger. Then, there will be a net flux of linear momentum carried by the gravitational radiation in the direction of the lighter black hole, and this will cause a recoil of the {\\em center of mass} in the opposite direction. The first analytic studies of this subject were carried out by Fitchett and Detweiler~\\cite{Fitchett:1983fc,Fitchett:1984fd}; Oohara and Nakamura~\\cite{Oohara:1983on}; Nakamura and Haugan~\\cite{Nakamura:1983nh}; and Wiseman~\\cite{Wiseman:1992dv}. Due to the strong non-linearity of the merger phase, analytic studies have difficulties in obtaining an accurate estimate of the recoil velocity. The first quasi-Newtonian analytic calculations were presented in~\\cite{Fitchett:1983fc,Fitchett:1984fd}, while a post-Newtonian (PN) analysis have been carried out in~\\cite{Wiseman:1992dv,Blanchet:2005rj,Damour:2006tr}.  Estimates using black-hole perturbation theory have been given in~\\cite{Favata:2004wz,Hughes:2004ck}, and a estimate that combines full numerical relativity and perturbation theory, the {\\em Lazarus} approach, is given in~\\cite{Campanelli:2004zw}.   Full numerical relativistic simulations are a natural approach to this problem since they can in principle handle the non-linearities of the gravitational field during the merger.  The challenge is the resolution that the computational resources impose.   Some calculations have already been carried out in different scenarios to estimate recoil velocities.  The first one was done by Anninos and Brandt~\\cite{Anninos:1998wt} for the case of the head-on collision of two unequal-mass black holes.  Their numerical calculations were effectively 2-dimensional since they made use of the axisymmetry of the configuration. Using the same type of numerical calculations they also estimated the gravitational radiation recoil from highly distorted black holes~\\cite{Brandt:1999zh}. More recently, and due to the significant advances in 3-dimensional numerical relativity in the binary black hole problem~\\cite{Pretorius:2005gq,Campanelli:2005dd,Baker:2005vv}, estimates of the radiation recoil velocity have also appeared~\\cite{Herrmann:2006ks,Baker:2006vn}.   Each of the approaches described above has its own limitations. Analytic approaches are able to provide accurate estimates in their region of validity. However, the largest contribution to the recoil velocity occurs during merger, precisely where the approximation methods break down. Numerical simulations, in principle, have the opportunity of producing estimates with a minimal number of assumptions.  However, as we have mentioned, these calculations have also limitations and use initial data that is only an approximation to the actual astrophysical configurations. Therefore, the error bars on the computed distribution of recoil velocities relative to the distribution present in nature are believed to be large and, even worse, are difficult to estimate.  It is then not surprising to find disagreements on the estimated recoil velocity as calculated with different methods.   In this paper, we present estimates of the recoil velocity using an approach not used before, the close-limit approximation (CLA), which combines both analytical and numerical techniques.  The CLA was introduced by Price and Pullin~\\cite{Price:1994pm}, who showed that this approximation method provides accurate results compared to those obtained from numerical relativity~\\cite{Anninos:1993zj} for head-on collisions of two black holes (see also~\\cite{Anninos:1995vf}). The CLA scheme is based on the assumption that in the last stage of coalescence, when the black holes are sufficiently {\\em close} to each other, the system behaves, up to a certain degree of approximation, as a single distorted hole. Then, the CLA scheme consists of establishing an appropriate correspondence between the binary black hole system and a single perturbed hole.  Once this correspondence is made, one can extract initial data that can be evolved by the perturbative relativistic equations. From the outcome of the evolution, one can estimate the fluxes of energy, angular momentum, and linear momentum carried away to spatial infinity by the gravitational radiation emitted. The CLA scheme has been developed and applied by a number of authors~\\cite{Gleiser:1996yc,Gleiser:1998rw,Krivan:1998er,Nicasio:1998aj,Gleiser:2001in,Khanna:2001ch,Sarbach:2001tj,Khanna:2002qp}. In particular, Andrade and Price~\\cite{Andrade:1996pc} used the CLA to estimate the recoil velocity of a head-on collision of unequal-mass black holes starting from rest.   Since the CLA applies to the last stage of the merger of two black holes, it is very appealing to use it to estimate the recoil velocity of the merger of an unequal-mass black hole binary system. With this scheme, we obtain a maximum recoil velocity of $\\sim \\{17,33,64\\}\\,$ km/s for a symmetric mass ratio $\\eta = M^{}_1M^{}_2/(M^{}_1+M^{}_2)^2 \\sim 0.19$ and initial proper separations of $\\{3,3.5,4\\}\\,M$, with $M$ the total ADM mass. Beyond a proper separation of $4\\,M$ the CLA is not expected to provide accurate results~\\cite{Andrade:1996pc}. Therefore, this method cannot account for the contributions during the inspiral and initial merger phase. Supplementing this estimate with PN calculations up to the innermost stable circular orbit (ISCO), we obtain a lower bound for the recoil velocity, with a maximum of $\\sim 84\\,$ km/s. This lower bound neglects the initial merger phase, for which we can obtain an approximate estimate by using either PN methods or the CLA. Since both methods are known to overestimate the amount of radiation during the early merger phase, we obtain, thus, an upper limit for the recoil with maxima in the range of $220-265\\,$ km/s.  We also perform non-linear fits to these bounds and obtain \\begin{equation} v_{\\FIT}  = a \\eta^2 \\sqrt{ 1 - 4 \\eta} \\left(1 + b \\eta + c \\eta^2 \\right), \\label{v-fit} \\end{equation} where $a=7782\\,$ km/s, $b=-2.507$ and $c=2.727$ for the lower bound and $a=14802\\,$ km/s, $b=-1.1339$ and $c=1.4766$ for the upper bound.   The plan of this paper is as follows: Sec.~\\ref{calculation} describes the main procedure of our calculation; Sec.~\\ref{bbhid} constructs initial data for a quasicircular binary black hole system in the $3+1$-formalism; Sec.~\\ref{close-limit} maps these initial data to a single perturbed black hole spacetime, such that it is suitable for a CLA evolution; Sec.~\\ref{resultscla} describes the numerical implementation and presents results from the evolution within the CLA scheme; Sec.~\\ref{totalrecoil} estimates the lower and upper bounds, as well as constructing the non-linear fits to these bounds; we finish in Sec.~\\ref{discussion} with a summary and discussion of the main results and points to future research to obtain improved estimates.   The conventions that we use throughout this work are the following: For the 4-dimensional spacetime, we use Greek letters for the indices and a semicolon for the covariant derivative.  The Schwarzschild metric admits a 2+2 decomposition consisting of the warped product of a Lorentzian 2-dimensional manifold with the 2-sphere (see~\\cite{Gerlach:1979rw,Gerlach:1980tx}). On the 2-dimensional Lorentzian manifold indices are denoted by capital Latin letters, the covariant derivative associated with the 2-dimensional metric is represented by a vertical bar, and the Levi-Civita antisymmetric tensor by $\\epsilon^{}_{AB}\\,$.  On the 2-sphere indices are denoted by the lower-case Latin letters $a,b,\\ldots,h$, the covariant derivate by a colon, and the Levi-Civita antisymmetric tensor by $\\epsilon^{}_{ab}\\,.$ When using the $3+1$ decomposition of spacetime quantities, spatial indices are denoted by the lower-case Latin letters $i,j,k,\\ldots\\,$. Uncontrolled remainders are denoted with ${\\cal{O}}(A)$ or ${\\cal{O}}(A,B)$, which stands for terms of order $A$ and terms of order $A$ or $B$ respectively. Although usually, when dealing with order symbols, $A$ and $B$ must be dimensionless, here they will not be, but can be made to be dimensionless through division by the appropriate factor. We also use physical units in which $G = c = 1$. ",
        "conclusions": "\\label{discussion}  We have calculated the recoil velocity after the merger of an unequal mass binary black hole system using the CLA scheme. This approximation assumes that the black holes are close enough that the system can be approximated by a single perturbed black hole spacetime. In contrast to other approaches, except for full numerical relativity, this approximation allows us to make valid statement about physical process when the system is close to plunge. Therefore, it is of great interest to use this method for the study of gravitational recoil. However, the CLA has the disadvantage that it cannot be used during the beginning of the merger or the inspiral phases.   Initial data for the CLA can be constructed analytically by mapping data suitable for a binary black hole inspiral to that of a single perturbed hole. With such initial data, the Cunningham-Price-Moncrief and the Zerilli-Moncrief master functions can be numerically evolved from some initial proper separation through ringdown. These gauge-invariant master functions contain all the information necessary to evaluate the gravitational metric waveforms and, thus, the energy, angular momentum and linear momentum fluxes carried away from the system.   The results obtained can be summarized as follows. First, the maximum recoil velocity obtained in the CLA scheme is of $v \\sim 64 \\,$ km/s for the maximum initial separation allowed by this method ($D=4 M$). This maximum occurs at a symmetric mass ratio of $\\eta \\sim 0.19$. By supplementing this estimate with PN ones valid in the inspiral regime, we obtain lower and upper bounds with maxima of $v \\sim 84 \\,$ km/s and $v \\sim 220 \\,$ km/s respectively. We have further provided non-linear analytic fit functions that conveniently parameterize these bounds, together with the results from the CLA, and other results in the literature. These results also suggest that there is a region around $D \\sim \\{4,6\\} \\,M$ that greatly contributes to the recoil, but can only be poorly modelled by current approximation schemes.   Ultimately, the estimates presented here suffer of the same predicaments as other calculations. Due to its analytical nature, the CLA relies on certain assumptions that do not hold over the entire history of the binary. Such assumptions introduce an error in the estimated recoil that is difficult to quantify. In particular, the assumptions made here are the following: close separations ; slow-motion; simple initial data. The close-limit assumption is essential to allow a mapping of a binary inspiral to a single perturbed spacetime. The slow-motion approximation supplements the close-limit assumption and can, in principle, be improved on in future extensions of this work. The choice of initial data is assumed to represent the gravitational content of some initial slice, although we know that this fails even at large separation because it does not agree with the deviations from conformal flatness predicted by PN theory. Moreover, it does not contain any radiation, which is not what it should be expected for initial data corresponding to a snapshot of the orbital evolution. Due to these assumptions, the estimate of the recoil velocity will be contaminated by some error. However, experience in CLA calculations indicates that the error made only overestimates the physical quantities calculated, relative to full numerical simulations~\\cite{Nicasio:1998aj,Khanna:1999mh,Khanna:2000dg}.   Future work will concentrate on extending this approach to second order in $P$ and to other, more realistic, initial data sets~\\cite{Yunes:2005nn,Yunes:2006iw}. Ultimately, it would be interesting to compare the CLA approach directly to full numerical relativistic simulations in an attempt to determine the region of validity of the CLA more accurately. Another possibility is to use a multi-parameter perturbation scheme (see~\\cite{Bruni:2002sm,Sopuerta:2003rg,Passamonti:2004je,Passamonti:2005cz}) where perturbations in the linear momentum and separation parameters can be cleanly separated at the different perturbative orders."
    },
    "0608/astro-ph0608093_arXiv.txt": {
        "abstract": "We describe a method for accurately determining M dwarf metallicities with spectral synthesis based on abundance analyses of visual binary stars. We obtained high resolution, high signal-to-noise spectra of each component of five visual binary pairs at McDonald Observatory. The spectral types of the components range from F7 to K3 V for the primaries and M0.5 to M3.5 V for the secondaries. We have determined the metallicities of the primaries differentially with respect to the Sun by fitting synthetic spectra to Fe \\textsc{i} line profiles in the observed spectra. In the course of our analysis of the M dwarf secondaries we have made significant improvements to the PHOENIX cool-star model atmospheres and the spectrum analysis code MOOG. Our analysis yields a RMS deviation of 0.11 dex in metallicity values between the binary pairs. We estimate the uncertainties in the derived stellar parameters for the M dwarfs to be 48 K, 0.10 dex, 0.12 dex, 0.15 km s$^{-1}$, and 0.20 km s$^{-1}$ for {T}$_{eff}$, log \\textsl{g}, [M/H], $\\xi$, and $\\eta$  respectively. Accurate stellar evolutionary models are needed to progress further in the analysis of cool-star spectra; the new model atmospheres warrant recalculation of the evolutionary models. ",
        "introduction": "The lowest mass stars, M dwarfs, are the most abundant stellar objects in our galaxy. They make up over 70\\% of stars in number and contribute over 40\\% of the total stellar mass content \\citep{henry98}. Despite their numbers, these stars remain one of the least understood stellar types. This is due to a lack of empirical data needed to test the predictions of theoretical models. Metallicity, one such parameter for which precise data do not yet exist for this stellar type, is the motivation for this study.  Metallicity, [M/H]\\footnote{We adopt the standard spectroscopic notation: for elements X and Y, log $\\epsilon$(X) $\\equiv$ log$_{10}(N_{X}/N_{H})$ + 12.0, [X/Y] $\\equiv$ log$_{10}(N_{X}/N_{Y})_{\\star}$ - log$_{10}(N_{X}/N_{Y})_{\\sun}$, and $N_{X}$ is the number density of element X. We assume [M/H] = [Fe/H] throughout this paper.}, is of particular importance in theoretical models of low-mass stars because of its role, along with mass, in governing early evolution and main sequence properties \\citep{kroupa97, baraffe98, siess00}. Additionally, stellar metallicity has been shown to correlate with the occurrence of extrasolar giant planets \\citep[most recently][]{fv05, santos05}. As planets are being discovered and confirmed around M dwarfs \\citep{delfosse98,marcy98,marcy01,benedict02,rivera05,butler04,bonfils05}, the question of these stars' metallicities has become increasingly relevant.  The historic lack of data for M dwarfs has been due primarily to their intrinsic faintness, a consequence of their low mass. However, despite the proliferation of large aperture telescopes and increasingly sensitive instrumentation, precise determinations of the fundamental parameters for M dwarfs are still rare. The factor now limiting the determination of accurate metallicities is the lack of a technique reliable enough to interpret the complex spectra of these stars.  Modern analysis techniques applied to high-resolution, R, and high signal-to-noise, S/N, spectra of solar-type stars consistently yield chemical abundances with internal precisions of 10\\% \\citep[e.g.][]{carlos04,vf05}. The application of these methods to M dwarfs is complicated by the effects of significant molecule formation and the resulting opacity in the photospheres of these stars. Molecular band spectra are much more complex than atomic spectra and typically dominate the spectral regions in which they are located. The regions traditionally utilized to derive metallicities for solar-type stars, the visible and red, are affected by TiO absorption bands in M dwarfs. These TiO lines blend with all other lines and create a ``pseudo continuum,'' making equivalent width measurements of atomic lines in this region unreliable for all but the earliest M dwarfs \\citep{woolf06}. Therefore, comprehensive synthetic spectrum analyses must be employed.  \\citet[hereafter V98]{v98} pioneered the use of spectral synthesis to determine M dwarf parameters to high precision. They fit synthetic spectra of TiO and atomic lines based on the NextGen version of the PHOENIX model atmospheres \\citep{haus99} to a high quality observed spectrum of the M3.5 V star Gl 725B to determine its parameters. Their precisions for the parameters effective temperatures, \\textsl{T}$_{eff}$, surface gravity, log \\textsl{g}, metallicity, [M/H], and macroturbulent velocity, $\\eta$, were 71 K, 0.14 dex, 0.07 dex, and 0.7 km s$^{-1}$ respectively.  In this paper we present the results of a test of the V98 technique and PHOENIX cool-star model atmospheres, and our work to improve both based on the abundance analyses of visual binary stars. In the next section we describe our high resolution, high S/N observations of five visual binary pairs containing primaries ranging from spectral type F7 V -\u2013 K3 V and secondaries M0.5 V -\u2013 M3.5 V. In \\S 3 we discuss our differential abundance analysis of the primaries with respect to the Sun. In \\S 4 we describe our application of the V98 technique to derive the metallicities of the M dwarf secondaries. Based on the assumption that stars in a bound system have the same metallicities, we show that the original V98 technique yields metallicities for M dwarfs that are systematically 0.5 dex too low. We then outline our updates and modifications to this technique in \\S 5. We show that a re-analysis of the M dwarfs in our binary sample, using our revised method and new model atmospheres, gives a RMS deviation of 0.11 dex from the primaries' metallicities. In \\S 6 we present an error analysis of the new technique, which yields uncertainties of 48 K for \\textsl{T}$_{eff}$, 0.10 dex for log \\textsl{g}, 0.12 dex for [M/H], 0.15 km s$^{-1}$ for the microturbulent velocity, $\\xi$, and 0.20 km s$^{-1}$ for $\\eta$. Finally, in \\S 7, we discuss the potential ramifications of our validation of the PHOENIX cool-star model atmospheres and our future applications of this technique to large samples of M dwarfs. ",
        "conclusions": "We have carried out a test of the V98 technique to determine M dwarf metallicities using an analysis of binary star pairs. Our result is that the V98 technique yielded values that were systematically too low. Motivated by this, we have made modifications to the original technique and a re-analysis of the M dwarfs secondaries validates our new approach. Our modifications include expanding the stellar analysis code MOOG, adopting relationships for the abundances of the alpha elements and carbon to iron, using TiO line data from \\citet{plez98}, inclusion of new atomic lines and data, introducing a variable microturbulence, and the use of an empirical surface gravity relationship. We have also made improvements in the PHOENIX model atmospheres that could have implications beyond those noted in our focused study.  Our modified technique, in conjunction with new model atmospheres, has yielded a technique for determining M dwarf metallicities consistent to 0.11 dex with the techniques applied to solar-similar stars. We have assumed that the mean offset in derived metallicities between the primary and M dwarf parent samples is actually zero and that the observed offset of -0.08 $\\pm$ 0.04 dex is due to random errors. The observed offset is actually 2$\\sigma$ from zero and might be indicative of an unknown systematic in our technique or the model atmospheres. We estimated the uncertainties in our derived stellar parameters for the M dwarfs to be 48 K, 0.10 dex, 0.12 dex, 0.15 km s$^{-1}$, and 0.20 km s$^{-1}$ for {T}$_{eff}$, log \\textsl{g}, [M/H], $\\xi$, and $\\eta$ respectively based on the assumption of no systematic errors. In future studies of single M dwarfs, we will adopt these values as the standard uncertainties in our analysis method when used with the particular model atmospheres described here.  Although we find good results for our study of the five binary pairs presented here, we are aware of several issues that could affect the results of our analyses of the primaries and the M dwarfs. In our analysis of the solar-similar primaries, there are potential systematics introduced by using the (\\bv) -- \\textsl{T}$_{eff}$ relationship of \\citet{ramirez05}. Applying their relationship to the Sun using the most probable (\\bv)$_{\\sun}$ of 0.64 \\citep{holmberg06} yields a \\textsl{T}$_{eff}$ of 5706 K, which is 71 K lower than the true \\textsl{T}$_{eff\\sun}$. We have estimated the errors in our derived metallicities for the primaries due to uncertainties in our adopted \\textsl{T}$_{eff}$ values on the same order as this difference (\\S 3.3), but our approach ignores the potential systematics in the \\citet{ramirez05} relationship for which it could be evidence. Nonetheless, we consider the \\textsl{T}$_{eff}$ values derived from the \\citet{ramirez05} relationship to be preferable to spectroscopically determined values for solar type stars which have their own potential problems \\citep{ramirez04}.  Additionally, the model atmospheres we used in our analysis of the the primaries used the ``old'' solar abundance values given by \\citet{anders89} and were not consistent with the models we used to analyze the M dwarfs. However, the abundances of the species that differed the most between the model versions used, C and O, affect the models of the primaries much less than the M dwarfs models. Model atmospheres of M dwarfs are heavily influenced by molecular opacities which depend directly and indirectly on the abundances of C and O. In comparison, the 1D model atmosphere structures of solar-similar stars are only minimally affected by opacities related to the C and O abundances at the level at which they have been altered (0.17 and 0.21 dex respectively).  Our technique for deriving the stellar parameters for M dwarfs relies on a number of assumptions that are not needed in analyses of solar-similar stars. These assumption could potentially introduce systematic errors that are not obvious given our test sample size. Improvements in our technique, including determining the abundances of specific elements instead of just the global metallicity, would be possible if there were more robust ways to estimate the stellar {T}$_{eff}$ and log \\textsl{g} values as there are for solar-type stars.  Our derived {T}$_{eff}$ values for the M dwarfs follow the expected trend and decrease smoothly with spectral type. However, when comparing these values with the {T}$_{eff}$ values as functions of spectral type in Table 4.1 of \\citet{reidhawley04}, there is a systematic trend of increasing deviation with spectral type. For the earliest M dwarf in our sample (HIP 32423B, M0.5), there is essentially no difference in our derived {T}$_{eff}$ and that given by interpolating in Table 4.1. For our sample, the deviation increases linearly with spectral type and is roughly 300 K for the latest M dwarf in our sample (HIP 12114B, M3.5). This comparison assumes that there is minimal dispersion in the spectral type -- {T}$_{eff}$ relationship, which is only true at a rough level (N. Reid, private communication) as such a relationship ignores the effects of metallicity. Also, only two M dwarfs, GJ 699 \\citep{dawson04} and GJ 411 \\citep{leggett96, seg03} have known strictly empirical {T}$_{eff}$ values based on measured bolometric fluxes and radii. However, the suggestion of a systematic error in our derived temperature scale should not be ignored lightly. As we have shown in \\S 6, changes in our adopted stellar {T}$_{eff}$ values of $\\sim$ 50 K correspond to differences of 0.12 dex in [M/H]. We plan to apply our analysis to the M dwarfs with known strictly empirical {T}$_{eff}$ values in order to test the accuracy of our spectroscopically derived {T}$_{eff}$ values and attempt to resolve this issue.  Also, we use an empirical relationship to constrain the surface gravity of our M dwarfs. Adopting values derived from fully validated evolutionary models might be more reliable because it would be possible to include the effects of varying stellar ages. We suggest that the biggest opportunity to advance detailed spectroscopic analyses of M dwarfs would by improving low-mass stellar evolutionary models. All stellar structure models depend on realistic model atmospheres as a boundary condition. Therefore, the new PHOENIX model atmospheres that we have validated in our analysis could be used to produce more reliable evolutionary models.  Although there is much left to do in the area of abundance analyses for M dwarfs, the work presented here has been the next step in the detailed study of these neglected objects. We consider our new technique to be robust enough for application to some current outstanding problems. Our research in the near-future will be focused on determining metallicities of M dwarfs with candidate extrasolar planets and those with accurate dynamical masses using the technique developed here."
    },
    "0608/astro-ph0608570_arXiv.txt": {
        "abstract": "Using a theoretical model describing pulse shapes, we have clarified the relations between the observed pulses and their corresponding timescales, such as angular spreading time, dynamic time as well as cooling time. We find that the angular spreading timescale caused by curvature effect of fireball surface only contributes to the falling part of the observed pulses, while the dynamic one in the co-moving frame of the shell merely contributes to the rising portion of them provided the radiative time is negligible. In addition, the observed pulses resulted from the pure radiative cooling time of relativistic electrons exhibit the property of fast rise and slow decay (a quasi-FRED profile) together with smooth peaks. Besides, we interpret the phenomena of wider pules tending to be more asymmetric to be a consequence of the difference in emission regions. Meanwhile, we find the intrinsic emission time is decided by the ratios of lorentz factors and radii of the shells between short and long bursts. Based on the analysis of asymmetry, our results suggest that the long GRB pulses may occur in the regions with larger radius, while the short bursts could locate at the smaller distance from central engine. ",
        "introduction": "Owing to possible overlapping between neighboring pulses, time profiles of gamma-ray bursts (GRBs) usually exhibit very diverse configurations$^{18}$. Pulses as the fundamental elements of bursts provide us a clue which will help us catch their quiddities. According to the standard fireball shock model$^{19, 33, 15}$, the pulses are generally proposed to either originate from the internal shocks produced by collisions between different shells in the fireball$^{24, 37}$ or come from the action of the external shock on the circum-burst medium$^{14, 26, 4, 6}$. However, problems are still in existence not only for their radiation mechanisms but also at the aspect of the differences between short and long pulses in their profiles.  It had been pointed out that GRB light curves especially the pulses in them could reproduce the temporal activity of its inner engine$^{10, 2, 17}$. Dermer \\& Mitman$^{5}$ had found that investigations on pulses were preferred for learning whether GRB sources require engines to be long-lasting or short-lived. The long-lasting central engine should produce more relativistic outflows during this phase which naturally lead to long-lasting emissions in the co-moving frame of the shells. The long time together with the delay time caused by the curvature effect of photosphere can inevitably results in longer duration in the observer frame$^{39,}$\\footnote{hereafter paper I}. Qin and Lu$^{23}$ thought that for most of observed pulses their corresponding co-moving pulses might contain a long decay time relative to the time-scale of the curvature effect. Moreover, studies with BATSE data for GRB pulses also indicate most sources do not display the curvature effects$^{11, 1}$, which suggests that the physics of such pulse formation is dominated or determined by other effects, perhaps the hydrodynamic or radiative cooling process or both of them. It happens that the relations between these timescales and temporal structure of GRBs had been researched by some authors$^{27, 29}$.  For the long and smooth single-peak GRBs, the curvature effects$^{9}$ or external shock$^{4}$ could be the leading action. The result that many short bursts are highly variable demonstrates short GRBs can't be produced by external shocks$^{16}$. So one of the aims in this work is to conjecture the approximate regions where bursts could occur. With this purpose, the structures of observed pulses connecting with different timescales are taken into account in very detail. These results suggest that short and long bursts might take place in different spatial regions from the sources. ",
        "conclusions": "By studying the influences of different timescales on the shape of pulses, we conclude that the profiles of observed pulses arising from angular spreading times would be long-tailed and called standard decay form without rising part, and those only resulting from dynamic timescales would follow a rising form. Which one is more dominant than another is determined by the value of $\\Delta\\tau_{\\theta}$, i.e. $\\Delta\\tau_{\\theta}>1$ for dynamic dominant case, and vice versa. The pure radiative cooling times would lead to smooth and FRED-like temporal profiles without too long tails. Additionally, we find the intrinsic emission time, $t^{\\mathrm{'}}$, is constrained by the ratios of lorentz factors and radii of the shells between short and long bursts.  Spada et al.$^{32}$ found by simulating internal shocks that the angular spreading and the dynamic times are comparable when a shell broadens linearly in some special regions. When it happens, the thickness $\\Delta^{'}$ of the shell could be estimated by $\\Delta^{'}\\sim R_c v'_{sh}/(\\Gamma c)$, which offer us a clue to understand that $T_{dyn}$ will go beyond $T_{ang}$ as $\\Delta^{'}\\gg R_cv'_{sh}/(\\Gamma c)$. In this case, this conclusion just meet those previous viewpoints that the effect of not the geometry but the hydrodynamics governs the temporal and spectral characteristics of GRB pulses$^{3}$. On the contrary, the component of duration from $T_{ang}$ would always be larger than from $T_{dyn}$ no matter how large the radii of fireball are.  On the other hand, for larger radii the radiative cooling time will be the dominant contribution to the pulse duration$^{32}$, thus the shape of observed pulses should reflect the properties dominated by radiation during the whole duration. Moreover, the lorentz factors, thickness of the merged shell, the frequency of emitted photons and the energy equipartition factor $\\varepsilon_{B}$ of the magnetic field will reduce to smaller values with the increasing of radii. On this occasion, there is a more strong contribution from $T_{syn}$ to observed pulses in contrast with either $T_{ang}$ or $T_{syn}$. Even if the outflows ejected from central engine are extremely calibrated in the early phase, the jet will spread sideways quickly$^{20}$ so that its geometry effects on observations are made to be more considerable. Therefore, the observed pulses are caused to exhibit more asymmetric and the trend of FRED are then strengthened.  We have known the leading model of central engine for long GRBs, perhaps including short ones, could be the collapsar model$^{36, 38}$. Besides, other merger models of two compact objects have also been proposed$^{44, 45, 46, 47}$. However, the analysis in this work is independent of the detailed progenitor models. The lorentz factors had been estimated as an order of 100-1000 for long bursts$^{51, 52}$, while its lower limit is about 500 for short ones (Norris and Bonnell 2006). Meanwhile, we have reached a conclusion in our recent works$^{48, 49}$ that the lorentz factors is proportional to $\\Gamma^{-\\omega}$ with $\\omega>2$ for short bursts and $\\omega<2$ for long bursts. In terms of the properties of tiny spectral lags in short bursts as well as their symmetric pulse characteristic, I thus infer that, although short and long GRB pulses could be interpreted with the same emission mechanism$^{12}$ in terms of the possible synchrotron radiation, the long and asymmetric FRED pulses could be produced by external shock in larger radii$^{14}$, while the short and symmetric pulses might be formed by internal shock in smaller emission distance from sources. We expect it to be verified in the future by the SWIFT satellite in view of its capability of precise and prompt localization."
    },
    "0608/astro-ph0608656_arXiv.txt": {
        "abstract": "We present an XMM-Newton observation of the radio galaxy 4C 23.56 at $z=2.48$ which reveals extended X-ray emission coincident with the radio lobes spanning $\\sim 0.5$ Mpc.  These are the largest X-ray-bright lobes known at $z>2$.  Under the assumption that these features result from inverse-Compton scattering of cosmic microwave background photons by relativistic electrons in the radio source lobes, the measured luminosity of $L_{\\mathrm{0.5-8 keV}}=7.5\\times10^{44}$ erg s$^{-1}$ implies a minimum energy stored within the lobes of $\\sim10^{59}$ erg.  We discuss the potential of the large-scale radio/X-ray lobes to influence evolution of the host galaxy and proto-cluster environment at high redshift. ",
        "introduction": "The galaxies hosting powerful radio-loud active galactic nuclei (AGN) at high redshift are believed to evolve into the most massive galaxies in the present-day Universe, which reside locally in rich cluster environments (e.g., Best et al., 1998). High-redshift radio galaxies (HzRGs) are therefore popular targets in studies of the co-evolution of massive galaxies, their central AGN, and their environment on larger scales.  In recent years, observations with Chandra and XMM-Newton have detected extended X-ray emission in an increasing sample of HzRGs (Fabian et al. 2003, 2003b; Carilli et al. 2002; Yuan et al. 2003; Scharf et al. 2003; Belsole et al. 2004; Overzier et al. 2005; Erlund et al. 2006; Blundell et al. 2006).  Generally associated with the radio structure, the extended X-ray emission is believed to arise in most cases partly or wholly from inverse-Compton (IC) scattering of cosmic microwave background (CMB) photons by relativistic electrons in the jets and/or lobes of the radio galaxy.  As noted by Schwartz (2002), the steep $(1+z)^4$ dependence of the density of CMB photons on redshift suggests this process will be increasingly significant toward higher redshifts.  Based on extrapolation from the radio source number counts, Celotti \\& Fabian (2004) estimate that at redshifts above $z\\gae1$ the majority of luminous ($L_X>10^{44}$ erg s$^{-1}$) extended X-ray sources may result from IC scattering of the CMB in HzRGs.  This new class of X-ray source may provide crucial insight into the effect of AGN activity on the evolution of massive galaxies and their larger scale environment.  The identification of bubble-like features in low-and intermediate-redshift clusters indicates AGN jets can inject a significant amount of energy into their environs (e.g., McNamara et al. 2000; Fabian et al. 2006).  Similar processes have been suspected to exert significant influence on the evolution of structure at high redshift, e.g., through the regulation of gas accretion and star formation in massive galaxies (e.g., Silk \\& Rees, 1998) and as a source of non-gravitational heating within forming cluster cores (e.g., Bower et al., 2001). However, the mechanisms and magnitude of AGN feedback at high redshift are unknown.  \\begin{figure*} \\begin{minipage}{\\textwidth} \\caption{{\\it a:} 0.5--8 keV EPIC image of 4C 23.56.  \\label{fig:image_raw} {\\it b:} Smoothed 2--8 keV EPIC image of 4C 23.56, with the region used for spectral analysis of the X-ray point source overlaid.  The darker line seen through the NW corner is a chip gap in the pn camera. \\label{fig:image_soft}{\\it c:} Smoothed 0.5--2 keV EPIC image of 4C 23.56, with regions used for spectral analysis of the extended emission overlaid.  {\\it d:} Difference image obtained by subtracting a smoothed 0.5--2 keV band image from a smoothed 2--8 keV band image.  Soft flux appears light (yellow) and hard flux appears dark (blue).  Labels indicate the Knopp \\& Chambers (1997) 6 cm radio positions of the 4C 23.56 core (C) and NE (B) and SW (A) lobes.  \\label{fig:image_sub}  Images are $2\\arcmin$ on a side and orientated so that North is up and East is to the left.} \\begin{center} \\includegraphics[width=0.23\\textwidth, angle=0]{epic_raw.epsf} \\hspace{1mm} \\includegraphics[width=0.23\\textwidth, angle=0]{pn_circle_colour.epsf} \\hspace{1mm} \\includegraphics[width=0.23\\textwidth, angle=0]{pn_panda_colour.epsf} \\hspace{1mm} \\includegraphics[width=0.23\\textwidth, angle=0]{ep_hmins_colour3.epsf} \\end{center} \\end{minipage} \\end{figure*}  The detection of IC-scattered X-ray emission in HzRGs has begun to clarify the evolution and energetics of high-redshift radio lobes.  As energy losses from the more energetic electron population responsible for the radio emission are more rapid than those of the less energetic population involved in IC scattering, the X-ray lobe emission can outlast the radio emission and provide a measure of past radio activity (e.g., Fabian et al. 2003; Blundell et al. 2006).  X-ray measurements have allowed estimates of the minimum Lorentz factors, $\\gamma_{min}$, in radio hotspots (Blundell et al., 2006) and the total energy contained within X-ray luminous radio lobes (Erlund et al. 2006) at $z\\sim2$.  Scharf et al. (2003), noting the relative positions of X-ray and Ly$\\alpha$ flux in 4C 41.17 at $z=3.8$, suggest that the IC-scattered X-ray photons could themselves give rise to a secondary feedback mechanism by photoionising the surrounding medium.  Both the direct mechanical feedback from the jets and this second radiative mechanism could influence the structure and energetics of forming structures on galactic and cluster scales and, interestingly, would preferentially affect the most massive structures at high redshift.  In this Letter we report the detection of extended X-ray emission associated with the radio galaxy 4C 23.56  ($z=2.483$; R{\\\"o}ttgering et al., 1997) in XMM-Newton observations obtained as part of a larger study of the X-ray properties of HzRGs. We begin with a summary of the multiwavelength properties of 4C 23.56 in \\S2, present the X-ray observations and analysis in \\S3 and briefly discuss our findings in \\S4.  Throughout this letter we assume a $\\Lambda$CDM cosmology with $\\Omega_m$=0.3, $\\Omega_{\\Lambda}=0.7$, and $H_0=70$ km s$^{-1}$ kpc$^{-2}$.  Unless specified otherwise, X-ray fluxes are given in the observed-frame 0.5--8 keV band, X-ray luminosities are corrected for Galactic absorption and given in the rest-frame 0.5--8 keV band, and errors are quoted at the 90\\% confidence level. ",
        "conclusions": "\\subsection{The nature of the extended X-ray emission}  The analyses above indicate soft, extended, X-ray emission coincident with the radio lobes of 4C 23.56 over a scale of $\\sim0.5$ Mpc with a luminosity of  $\\sim7\\times10^{44}$ erg s$^{-1}$.  While there are several possible emission mechanisms for extended X-ray flux in HzRGs (see, e.g., Overzier et al., 2005 and references therein), few can account for the morphology and magnitude of the flux seen in 4C 23.56.  Extrapolating the synchrotron spectrum from the radio waveband, we find the observed X-ray luminosity is two orders of magnitude higher than would be expected from synchrotron emission.  Synchrotron self-Compton emission, often invoked to explain X-ray hotspots in local radio galaxies (e.g., Harris et al. 2000), cannot produce X-rays over a large area because the synchrotron photon density can only exceed that of the CMB in compact regions.  IC scattering of optical/IR photons from the central AGN (Brunetti et al., 2001) can not account for emission so far from the core.  Shock heating of an ambient medium by the expanding radio source --- a proposed mechanism for the extended X-ray emission in PKS1138-262 (Carilli et al., 2002) --- would require an enormous reservoir of hot, high filling-factor gas which is not detected in the off-axis regions.  Finally, thermal emission from hot intracluster gas cannot be ruled out, but the surface brightness of the observed emission is much higher than expected and a contrived containment model would be needed to explain an extremely elongated morphology aligned with the radio lobes.  We believe that in 4C 23.56, as in most previously reported HzRGs with X-ray-bright radio lobes, IC scattering of CMB photons is the most likely emission mechanism for the majority, if not the entirety, of the extended X-ray flux.  Fourteen $z>1$ HzRGs currently in the literature (see references in \\S1) have X-ray lobes on scales ranging from a few tens to a few hundreds of kpc with typical luminosities of a few times $10^{44}$ erg s$^{-1}$.  The X-ray lobes in 4C 23.56 are the second largest discovered to date\\footnote{The giant radio galaxy 6C 0905+3955 at $z=1.88$ exhibits X-ray emission associated with hotspots and lobes which are separated by $\\sim 865$ kpc (Blundell et al. 2006).}, and the first $\\gae 0.5$ Mpc-scale lobes discovered at the $z>2$ epoch suspected to be the predominant era of massive galaxy formation.   \\subsection{Energetics of the X-ray-detected lobes}  We now briefly consider the energetics of the X-ray detected lobes, assuming the entirety of the extended X-ray emission arises in IC scattering of CMB photons.  Particles with a large Lorentz factor, $\\gamma_{e}$, will upscatter CMB photons of frequency $\\nu_{CMB}$ to higher frequencies $\\nu_{scat}$ as  \\begin{equation} {\\nu_{scat} \\over \\nu_{CMB}} \\simeq {4 \\over 3} \\gamma_{e}^2. \\end{equation}  \\noindent Scattering to X-ray energies therefore requires electrons with a Lorentz factor $\\gamma_{e}\\sim 10^3$.  If we assume that all electrons responsible for the X-ray lobes have this Lorentz factor, we can use the observed X-ray luminosity and the known CMB photon energy density to estimate the energy stored in the IC-emitting electron population.  The energy flux in scattered X-rays is  \\begin{equation} L_{X}=N_{e} U_{rad} {4 \\over 3} \\gamma_{e}^2 \\sigma_{T} c, \\end{equation}  \\noindent while the energy of the relativistic electrons is  \\begin{equation} E_{e}=N_{e} \\gamma_{e} m_{e} c^2, \\end{equation}  \\noindent where $N_{e}$ is the number density of relativistic electrons, $\\sigma_{T}$ is the Thompson cross-section, $c$ is the speed of light, and $m_{e}$ is the rest mass of the electron.  $U_{rad}$ is the energy density of the CMB, which increases from the local value with redshift as $(1+z)^4$.  The energy in the scattering electron population is then given by  \\begin{equation} E_{e}={3 \\over 4} {L_{X} m_{e} c \\over U_{rad} \\gamma_{e} \\sigma_{T}} \\simeq {3 \\over 4} {L_{44} \\over \\gamma_{e}(1+z)^4} 10^{64}\\, \\mathrm{ erg,} \\end{equation}  \\noindent where $L_{44}$ is the 2--10 keV X-ray luminosity in units of $10^{44}$ erg s$^{-1}$.  As noted by Erlund et al. (2006), this calculation provides only a lower limit on the energy in the scattering population.  Assuming a more realistic power-law distribution of Lorentz factors reveals dependencies of the total electron energy on the minimum and maximum Lorentz factors and on the X-ray photon index.  These parameters are poorly constrained; adopting typical values increases the estimated energy by a factor of a few.  In addition, this estimate considers only the energy in relativistic electrons.  A substantial population of relativistic protons (e.g., Celotti \\& Fabian, 1993; Sikora et al., 1995) could further increase the implied lobe energies by a factor of up to $10^3$.  Converting our measured 0.5--8 keV luminosity to 2--10 keV with $\\Gamma=3$, we find a lower limit to the energy stored in relativistic electrons in the lobes of 4C 23.56 of $E_{e}\\simeq8.0\\times10^{58}$ erg; adopting a more typical $\\Gamma=2$, $E_{e}\\simeq2.2\\times10^{59}$ erg.  Comparing these estimates with those of Erlund et al. (2006) for three $z\\sim2$ HzRGs we note that while the X-ray lobes in 4C 23.56 are at least twice as large as the lower-redshift sources, the minimum energy stored in IC-emitting electrons appears to be roughly comparable.  Considering only losses through IC emission, the radiative lifetime of the 4C 23.56 X-ray lobes is $\\sim10^{7}$ yr.   \\subsection{Potential for environmental feedback}  An enormous reservoir of energy --- equivalent to $\\sim 10^8$ supernovae --- is stored in the IC-emitting electron population within the radio lobes of 4C 23.56. As the cooling timescale of this population through IC-scattering to X-ray wavelengths is relatively long, a large fraction of this energy could potentially be injected into the surroundings during a given epoch of AGN activity. The vast extent of the X-ray luminous lobes suggests a mechanism by which significant feedback could take place at radii of up to 250 kpc from the central nucleus.  It is not clear, however, if the energy available in the IC scattering electron population is sufficient to be a major source of feedback influencing the evolution of either the host galaxy or the forming cluster environment.  Assuming a velocity dispersion of 300 km s$^{-1}$, the lobe energy estimates derived above are 10--20 times lower than the binding energy of a $10^{12}$ M$_{\\odot}$ elliptical and roughly 500 times lower than the $5\\times10^{49}$ erg M$_{\\odot}^{-1}$ estimated by Benson et al. (2003) to be required to inhibit the excessive growth of such a galaxy at high redshift.  In terms of heating of the intracluster medium, Wu, Fabian, \\& Nelson (2000) estimate non-gravitational heating of 1--2 keV per particle is needed to explain the observed deviation of galaxy cluster properties from the theoretical scaling relations.  For a cluster with a gas mass of $10^{13}$ M$_{\\odot}$, this is equivalent to $\\sim5\\times10^{61}$ erg or roughly 300 times our minimum energy estimate for the IC-scattering population in the lobes of 4C 23.56.  Unless the assumption of a single Lorentz factor, $\\gamma_{e}=10^3$, is dramatically wrong or there is a significant proton component to the jet, the energy contained in the IC-scattering population in the lobes appears to fall short of the required feedback energies for the host galaxy and cluster environment by at least two orders of magnitude, even assuming perfectly efficient energy transfer acting over the required scales.  In addition, the observation of significant cooling in IC-scattering of CMB photons is evidence that a significant fraction of this energy does not get injected into the local environment.  Nevertheless, IC scattering could provide a very efficient mechanism for producing hard, ionising photons on scales far from the central AGN.  Given reasonable assumptions for the halo gas density, X-ray luminosity at the level observed in 4C 23.56 could potentially ionize a large volume of halo gas over scales of several hundred kpc, as outlined in Scharf et al. (2003).  Moreover, several additional mechanisms could increase the feedback potential of the X-ray lobes dramatically.  Energy stored within the magnetic field, found to be of order $E_{e}$ in the Erlund et al. (2006) sources, may increase the total energy of the jet.  More significant may be the $pdV$ work associated with the expansion of lobes into the surrounding medium, which has been estimated to provide energies of up to $10^{61}$ erg in lower redshift radio galaxies and observed to significantly heat the surrounding intracluster medium (e.g., B{\\^i}rzan et al., 2004; McNamara et al., 2005).  Finally, bursts of AGN activity may occur several times over the course of a Gyr, re-accelerating the radio lobes and driving subsequent periods of environmental feedback."
    },
    "0608/astro-ph0608183_arXiv.txt": {
        "abstract": "We present observations of the unusually bright and long $\\gamma$-ray burst GRB\\,050820A, one of the best-sampled broadband data sets in the \\Swift\\ era.  The $\\gamma$-ray light curve is marked by a soft precursor pulse some 200 s before the main event; the lack of any intervening emission suggests that it is due to a physical mechanism distinct from the GRB itself. The large time lag between the precursor and the main emission enabled simultaneous observations in the $\\gamma$-ray, X-ray, and optical band-passes, something only achieved for a handful of events to date.  While the contemporaneous X-rays are the low-energy tail of the prompt emission, the optical does not directly track the $\\gamma$-ray flux.  Instead, the early-time optical data appear mostly consistent with the forward shock synchrotron peak passing through the optical, and are therefore likely the beginning of the afterglow.  On hour time scales after the burst, the X-ray and optical light curves are inconsistent with an adiabatic expansion of the shock into the surrounding region, but rather indicate that there is a period of energy injection.  Observations at late times allow us to constrain the collimation angle of the relativistic outflow to $6.8^{\\circ} \\lesssim \\theta \\lesssim 9.3^{\\circ}$.  Our estimates of both the kinetic energy of the afterglow ($E_{\\mathrm{KE}} = 5.2^{+7.9}_{-4.1} \\times 10^{51}$ ergs) and the prompt $\\gamma$-ray energy release ($E_{\\gamma} = 7.5^{+6.7}_{-2.4} \\times 10^{51}$ ergs) make GRB\\,050820A one of the most energetic events for which such values could be determined. ",
        "introduction": "\\label{sec:intro}  With the discovery of the cosmological nature of $\\gamma$-ray bursts (GRBs) in 1997 \\citep{mdk+97}, astronomers were suddenly forced to explain the enormous isotropic energy release of these distant explosions. Some of the most energetic events, such as GRB\\,990123, seemingly released enough energy in the prompt $\\gamma$-rays ($E_{\\gamma,\\mathrm{iso}} = 1.2 \\times 10^{54}$ ergs; \\citealt{bbk+99}) to rival the rest mass of a neutron star.  Furthermore, broadband modeling of the best sampled events has shown that a comparable amount of energy remains in the shock, powering the long-lived X-ray, optical, and radio afterglow (see e.g.~\\citealt{pk01,yhs+03}).  The hypothesis that GRBs are a-spherical explosions \\citep{r99}, supported by the appearance of achromatic ``jet'' breaks in a large number of afterglow light curves \\citep{sph99}, proved to be a turning point.  With typical opening angles of a few degrees, the true energy release from most GRBs is $\\sim 10^{51}$ ergs, on par with that of a supernova (SN).  This realization enabled the discovery of a standard energy reservoir for the collimation-corrected prompt energy \\citep{fks+01} and kinetic energy of the afterglow \\citep{bkf03}.  GRBs are now considered promising standard candle candidates, with the hope of Hubble diagrams out to $z \\approx 6$ offering complementary constraints to Type Ia SNe on the cosmology of our universe (\\citealt{fag+06,dlx04}; c.f.~\\citealt{fb05}).  Launched in November 2004, the \\Swift\\ Gamma-Ray Burst Explorer \\citep{gcg+04} was designed to position GRBs, disseminate accurate coordinates to ground-based observatories in real-time, and follow the UV and X-ray afterglows from minutes to days after the event.  In only a year of full operation, \\Swift\\ has brought about a number of fundamental advances in the GRB field, including the discovery of the first X-ray (GRB\\,050509b; \\citealt{gso+05}) and near-infrared (GRB\\,050724; \\citealt{bpc+05}) afterglows of a short-hard burst, the detection of the high-redshift ($z = 6.3$) GRB\\,050904 \\citep{hnr+06,cmc+06}, and the ability to measure broadband light curves starting shortly after, and in a few cases even during the $\\gamma$-ray event itself.  Despite these advances, measuring the bolometric fluences of \\Swift\\ events has proved challenging, for a number of reasons.  First, the limited energy range of the \\Swift\\ Burst Alert Telescope (BAT; \\citealt{bbc+05}) means that \\Swift\\ can accurately characterize only the softest GRB spectra. Second, few \\Swift\\ events have shown conclusive signs of a jet break, leaving geometric corrections highly uncertain.  Finally, the X-ray light curves of \\Swift\\ afterglows have shown both bright flares \\citep{brf+05} and slow decays \\citep{nkg+06}.  Both behaviors have been attributed to late-time ($t >> t_{\\mathrm{GRB}}$) energy injection, and at times have rivaled the energy release of the prompt $\\gamma$-ray emission (see e.g.~\\citealt{fbl+06}).  On 20 August 2005 UT, the BAT detected and localized the unusually bright and long GRB\\,050820A, a truly rare burst in the \\Swift\\ sample.  The $\\gamma$-ray light curve is marked by a soft pulse of emission preceding the main event by over 200 s.  The main emission was bright enough to be detected by the \\KW\\ instrument, providing a $\\gamma$-ray spectrum extending beyond 1 MeV as well as continuous coverage over the entire $\\sim$ 600 s burst duration.  Since \\Swift\\ triggered on the precursor, both space- and ground-based facilities were able to image the transient during the bulk of the prompt emission.  Such contemporaneous multi-wavelength observations have only been achieved for a handful of bursts to date.  The bright X-ray ($F_{\\nu} \\sim 0.7$ mJy) and optical ($R \\sim 14.5$ mag) afterglows made it possible to study the evolution of the afterglow for weeks after the burst, providing one of the most detailed broadband light curves in the \\Swift\\ era.  Finally, late-time \\HST\\ observations allowed us to constrain the jet-break time, and hence the geometry of the outflow.  Even after applying the collimation correction, we find GRB\\,050820A is an exceptionally energetic event.  Our work proceeds as follows: in \\S\\ref{sec:obs}, we outline our broadband observations of the GRB\\,050820A, beginning with the high-energy prompt emission and followed by the X-ray, optical, and radio afterglow.  We find the afterglow data are incompatible with the standard model of synchrotron radiation from a single, highly-relativistic shock expanding adiabatically into the surrounding medium \\citep{spn98}.  Instead in \\S\\ref{sec:anal} we use power-law fits ($F_{\\nu} \\propto t^{-\\alpha} \\nu^{-\\beta}$) to model the afterglow, dividing the burst into segments based on noticeable temporal breaks in the X-ray and optical light curves.  This analysis allows us to investigate the early broadband light curve (\\S\\ref{sec:pre} and \\ref{sec:early}), late-time ($t >> t_{GRB}$) energy injection in the forward shock (\\S\\ref{sec:inject}), the structure of the circum-burst medium (\\S\\ref{sec:cbm}), and the geometry and energetics of the event (\\S\\ref{sec:energy}).  Throughout this work we adopt a standard cosmology with $H_{0} = 71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\mathrm{M}} = 0.73$, and $\\Omega_{\\Lambda} = 0.27$.  We also make use of the notation $Q_{X} \\equiv 10^{X} \\times Q$. ",
        "conclusions": "\\label{sec:conc}  GRB\\,050820A joins a select sample of events with simultaneous observations in the $\\gamma$-ray and optical bands, and an even smaller group with contemporaneous X-ray observations as well.  Such events have led to fundamental advances in our understanding of GRBs, including the discovery of a reverse-shock optical flash from GRB\\,990123 \\citep{abb+99} and possibly from GRB\\,050904 \\citep{bad+06}.  The early-time optical emission from GRB\\,041219A also showed a bright optical flash, but the rise was correlated with an accompanying peak of $\\gamma$-ray emission, suggesting a common origin for the two components \\citep{vww+05}.  The early behavior of GRB\\,050820A is unlike either of these events. \\citet{vww+06} have shown the contemporaneous optical emission is well described as the sum of two components: one proportional to the prompt $\\gamma$-ray emission and one smoothly varying forward shock term.  While the $\\gamma$-ray component is important for $t_{\\mathrm{BAT}} < 300$ s, the optical peak at $t_{\\mathrm{BAT}} \\approx 600$ s is dominated by emission from the external shock region. Furthermore, the post-peak decay rate is inconsistent with reverse-shock emission.  Instead we interpret this as the forward shock peak frequency passing through the optical bands.  This is not unlike what was seen in the optical for GRB\\,060124 \\citep{rcc+06}, although the time resolution in the prompt phase was much poorer than for this burst. The contemporaneous optical light curves of GRB\\,050319 \\citep{qry+06,wvw+05} and GRB\\,050401 \\citep{ryk+05} did not show this peak phase, but extrapolations to late-times were consistent with the adiabatic expansion of a forward shock. Another different behavior was seen in the early optical light curve of GRB\\,050801, which showed an extended plateau phase correlated with the X-ray emission, hinting at continued energy injection from a central engine refreshing the external shocks \\citep{rmy+06}.  The contemporaneous X-ray emission, on the other hand, is the low-energy tail of the prompt emission.  This behavior was also seen for GRB\\,060124 \\citep{rcc+06} and has been hinted at in the rapid decline in early X-ray light curves attributed to high-latitude emission \\citep{lzo+06}, as well as the bright X-ray flares seen in many XRT light curves \\citep{brf+05}. It is clear then, that, in marked contrast to the X-ray emission, contemporaneous optical emission exhibits a large diversity in behavior.  Unfortunately the physical mechanism behind this diversity remains to be explained.  The issue of burst geometry is a particularly interesting one in the \\Swift\\ era.  The steep post-break decay slope, seen simultaneously in multiple filters, makes GRB\\,050820A one of the most convincing examples of a beamed event in the \\Swift\\ sample. The X-ray afterglow, however, is too faint at late times to provide broadband confirmation.  In fact, very few \\Swift\\ bursts, including those, like GRB\\,050408 \\citep{fpp+06}, that have been followed for months, show signs of a jet break in the XRT light curve \\citep{nkg+06}.  Typical jet breaks in pre-\\Swift\\ bursts occurred on time-scales of several days \\citep{zkk06}.  \\citet{psf03} predicted that \\Swift\\ would detect bursts with wider opening angles than previous missions due to the increased sensitivity of the BAT.  However, not a single \\Swift\\ afterglow has shown a convincing jet-break-like transition in multiple band-passes (candidates include GRB\\,050525A and GRB\\,050801; see \\citealt{pmb+06} and references therein for a more thorough discussion).  While it may be that most jet breaks, like GRB\\,050820A, occur at late times, beyond the sensitivity of the XRT and most ground-based facilities, this is nonetheless troubling.  On the one hand, few if any X-ray jet breaks were seen in pre-\\Swift\\ bursts; all collimation angles were determined from optical (and sometimes radio) light curves.  Conversely, given the large number of well-sampled XRT light curves, and the fact that such fundamental results for GRB cosmology as the Ghirlanda relation rest on our picture of GRBs as a-spherical events, this is clearly a matter that merits further investigation.  One consequence of the large opening angle associated with GRB\\,050820A is a correspondingly large burst and afterglow energy.  In fact, of all the bursts compiled in the \\citet{ggl+04} sample, GRB\\,050820A has the largest prompt energy release.  And unlike GRB\\,990123 \\citep{pk01}, this large $\\gamma$-ray energy was accompanied by a correspondingly large kinetic energy imparted to the afterglow.  The only comparable event for which such energies could be determined was the high-redshift GRB\\,050904, which released a total energy of $\\sim 10^{52}$ ergs \\citep{tac+05,fck+06}. Given the large $\\gamma$-ray fluence, similar events should have been easily detected by both \\Swift\\ and previous GRB missions.  And given the bright optical afterglow and the late jet-break, such events are strongly favored for ground-based follow-up (i.e.~redshift determination).  The lack of a large sample of such events means they must be relatively rare in the universe.  Like many other \\Swift\\ GRBs, the X-ray light curve of GRB\\,050820A exhibits a phase of shallow decay incompatible with the standard forward shock model \\citep{nkg+06}.  GRB\\,050820A is relatively unique, however, in that this epoch is also well-sampled in the optical.  The seemingly simultaneous breaks in the optical light curve bolster the commonly-held belief that this phase is caused by some form of refreshed shocks \\citep{zfd+06}. Coupled with the large gap between the precursor and the bulk of the prompt emission, the late-time energy injection poses fundamental challenges to any central engine model.  Finally, we return to the question of the radio afterglow.  Radio observations typically probe low-Lorentz factor ejecta ($\\Gamma \\sim 2-3$) at large distances from the central engine ($r \\sim 10^{17}$ cm).  The forward shock peak frequency reaches the radio much later than the optical or X-ray bands.  Thus, radio emission is usually visible at later times than optical or X-ray emission, and is well-suited to study afterglows when the emission is isotropic (i.e.~after the jet break) or even in some cases when the ejecta has slowed to Newtonian expansion \\citep{bkf04}. For this reason, late-time radio observations are considered the most accurate method for model-independent calorimetry. For GRB\\,050820A, this paradigm has broken down.  The burst had a bright optical and X-ray afterglow, but weak emission in the radio. It is hoped that further studies of such energetic GRBs in the \\Swift\\ era will help to elucidate some of these puzzles."
    },
    "0608/astro-ph0608526_arXiv.txt": {
        "abstract": "s{ Galactic {\\em dark matter} is modelled by a scalar field in order to {\\em effectively} modify Kepler's law {\\em without} changing standard Newtonian gravity. In particular,  a {\\em solvable} toy model with a {\\em self-interaction} $U(\\Phi)$ borrowed from non-topological solitons produces already qualitatively correct {\\em rotation curves} and {\\em scaling relations}. Although relativistic effects in the halo are very small, we indicate corrections arising from the general relativistic formulation. Thereby, we can also probe the weak gravitational lensing of our soliton type halo. For cold scalar fields, it corresponds to a gravitationally confined Boson-Einstein condensate, but of galactic dimensions.\\\\ \\\\ } ",
        "introduction": "The existence of {\\em dark matter} was postulated for the first time by Zwicky\\cite{zwi} in 1933. He showed that the virial masses of clusters of galaxies exceed the mass of luminous matter by roughly one order of magnitude. In the 1970's spectrographs became sensitive enough that rotation curves of spiral galaxies could be observed in large numbers by Rubin and collaborators\\cite{RFT80,RBFT85}. They all showed a systematic tendency to stay flat as function of galactocentric radius, $v_\\varphi (r) \\simeq {\\rm const.}$, whereas the mass distributions of the bulges and the radially exponentially declining surface densities of the galactic disks  lead to the prediction of falling rotation curves. Evidence for an extra mass component became even more decisive when HI--rotation curves\\cite{ABBS} were observed, because the neutral hydrogen gas extends radially outwards many disk scale lengths. Even though there is a long debate\\cite{Bo03} to what degree the inner parts of galaxies are dominated by dark matter, the flat outer HI-rotation curves prove beyond doubt the existence of dark matter halos. The analysis of the kinematics of satellite galaxies\\cite{zar2} showed that the dark halos extend out to several hundred kpc with the cumulative mass still rising linearly with galactocentric radius, $M(r) \\propto r$, which corresponds to flat rotation curves.  Persic {\\em et al.}\\cite{PS96} investigated statistical correlations between structural parameters of visible and dark matter using a data set of almost 1000 galaxies and developed the model of the so called `universal rotation curve'. Accordingly, galaxies of low luminosity are more dominated by dark matter than bright galaxies. Dwarf irregular galaxies and low surface brightness galaxies seem to be dominated completely by dark matter. Burkert\\cite{Bu95,BS99} finds for such galaxies an empirical fit to the rotation curves which implies asymptotically $v_\\varphi^2 \\rightarrow \\ln r/r$, i.e. a logarithmic modification of the Kepler law.   For several classes of gravitational theories, it has been argued that the introduction of dark matter is necessary\\cite{nester}, although in the modified Newtonian dynamics (MOND) scheme, exact expressions were derived for the Newtonian potential and for the rotation curve\\cite{BM95}. However, recent observational constraints\\cite{AB01} make it more and more challenging to interpret galactic rotation in terms of new gravitational physics\\cite{Be88}.   An intriguing possibility is a {\\em scalar field} model of the halo first proposed by one of us\\cite{Sch94,schunck,sch97,Sc98,Sc99} using a complex massless field coupled to the Einstein equation, an idea which was later taken up by several other authors. Hence, the complex massless scalar field forming star-like galactic halos embodies another form of {\\em dark energy}.  Massive compact halo objects, so-called MACHOs\\cite{alc}, consisting of baryonic matter or, alternatively, of {\\em boson star}s (BSs)\\cite{MS00,MS02a,SM03} do not seem to be sufficient to resolve this problem completely. More recent approaches\\cite{SKM03,FM04} propose connections between scalar field matter and the cosmological constant.  On the other hand, the standard {\\em cold dark matter} (CDM) model, where dark matter particles interact only through gravity, may be in conflict with observational features on Mpc scale. In particular,  the CDM model of Navarro, Frenk and  White (NFW)\\cite{b1} tends to predict density profiles of dark halos which are cuspy at the center\\cite{a17}. Therefore, Spergel and Steinhardt\\cite{{SS00}} proposed that dark matter is self-interacting, an idea which found ramification in Refs.~\\refcite{RT00,a13}. The evolution of the {\\em self-interacting} dark matter  has been recently analyzed in more detail\\cite{BS02,BSI02}.   In the same vein,  we consider here a model\\cite{MS02,MSP02} where dark matter is modelled by a primordial scalar field with a {\\em self-interaction} $U(\\Phi)$. Since the observed rotation velocities are roughly bounded\\footnote{The maximum velocity within a rotation curve of almost 400 km/s is found in a Sa galaxy\\cite{RFT80}; the velocity of 367 km/s at the farest  measured point  belongs also to a Sa galaxy\\cite{RBFT85}.}  by $v_\\varphi \\leq 10^{-3}c$, {\\em i.e.} are non-relativistic, a Newtonian type approximation sufficies.  The formation of dark halos as a gigantic Bose-Einstein condensation in early Universe is discussed in Sec.~\\ref{BEC}. In particular, a boson star like toy model with a $\\Phi^6$ repulsive self-interaction is known\\cite{Mi78} to allow in the limiting case of {\\em flat} spacetime  an exact spherically symmetric {\\em soliton} solution of the corresponding nonlinear Klein-Gordon equation, as is shown in Sec.~\\ref{ntssol}. Simulating the halo by such a {\\em non-topological soliton} (NTS) for the {\\em positive} range of the potential, yields a Newtonian mass distribution which provides a qualitatively rather good fit to the rotation curve data of dwarf irregular and low surface brightness galaxies, see Sec.~\\ref{rotcur} and the comparison with observations in Sec.~\\ref{observ}. In Sec.~\\ref{outlook}, we indicate the extension to the general relativistic formulation and speculated on the superstring nature of the dark matter scalar. In Sec.~\\ref{lensing}, we display the gravitational lensing properties of the NTS halo. ",
        "conclusions": ""
    },
    "0608/astro-ph0608289_arXiv.txt": {
        "abstract": "s { The MareNostrum Universe is one of the biggest SPH cosmological simulations done so far. It contains more than 2 billion particles ($2\\times 1024^3$) in a 500 h$^{-1}$ Mpc cubic volume.  This simulation has been performed on the MareNostrum supercomputer at the Barcelona Supercomputer Center. We have obtained more than 0.5 million halos with masses greater than a typical Milky Way galaxy halo.  We report results about the halo mass function, the shapes of dark matter and gas distributions in halos, the baryonic fraction in galaxy clusters and groups, baryon oscillations in the dark matter and the halo power spectra as well as the distribution and evolution of the gas fraction at large scales.  } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608240_arXiv.txt": {
        "abstract": "{We present results of Very Long Baseline Array (VLBA) observations of PKS 0528+134 at five frequencies (2.3, 5.0, 8.4, 15.4, and 22.2 GHz). These quasi-simultaneous data enable us to study the spectral distribution of Very Long Baseline Interferometer (VLBI) components for the first time in this highly variable source, from which the central compact core is identified. Our observations indicate that there are two bendings for the jet motion at parsec scale. We provide an approximate spatial fit to the curved jet trajectory using the Steffen et al. (\\cite{Steffen95}) helical model. We further investigate the proper motions of three jet components, which all show superluminal motion. At high frequencies (15.4 and 22.2 GHz) we detected a new component, which is estimated to be related to a radio burst peaking at about 2000. ",
        "introduction": "As one of the brightest active galactic nuclei, \\object{PKS 0528+134} has been intensively studied at the radio, X-ray, and gamma-ray bands. It is one of the furthest blazars detected by EGRET above 100 Mev (Mukherjee et al. \\cite{Mukherjee97}) with red-shift z=2.07 (Hunter et al. \\cite{Hunter}). There are relatively less optical observations compared with the observations at other bands because of its location close to the Galactic plane (galactic latitude b=$-11\\degr0'43''$) and thus high Galactic extinction. Whiting et al. (\\cite{Whiting03}) obtained $B=17.04\\pm0.23$, $V=16.74\\pm0.15$ at 2001.66 after the extinction corrections. But these are different from the previous values B=20.0 (Condon, Hicks \\& Jauncey \\cite{Condon77}) and V=19.5 (Wall \\& Peacock \\cite{Wall85}), implying its optical variability. People have observed its variability on a timescale as short as a few days at gamma-ray (Mukherjee et al. \\cite{Mukherjee96}), as short as two days in the optical band (Ghosh et al. \\cite{Ghosh00}), and from several months to a few years in the radio band (e.g., Pohl et al. \\cite{Pohl95}; Peng et al. \\cite{Peng01}). There is a delay for the radio burst from high frequencies to low frequencies. Peng et al. (\\cite{Peng01}) showed in their light curves that the delay for the 1996 burst is about 0.7 yr between 10.7 and 5.0 GHz. Mukherjee et al. (\\cite{Mukherjee96}) showed that the delay between the gamma-ray flare and the radio burst is a few months.  Very Long Array (VLA) observations detected a diffuse emission 2$''$ in extent and centered 1$''$ north of the core with $<$ 0.2\\% and 1.2\\% of the peak brightness at 4.9 and 1.5 GHz, respectively (Perley \\cite{Perley82}). \\object{PKS 0528+134} is frequently observed as a calibrator for VLBI observations (e.g. Shen et al. \\cite{Shen05}). The first VLBI images of \\object{PKS 0528+134} (Charlot \\cite{Charlot90}) show a one-sided extension at 2.3 GHz and a double structure at 8.4 GHz, indicating a core-jet structure at P.A. (position angle) $\\sim-140\\degr$. Zhang et al. (\\cite{Zhang94}) studied in detail the geometry, radiation mechanism, and physical environment of the emission region in the source through the X-ray and VLBI observations, and found a bent jet extending toward the northeast on a parsec scale. Krichbaum et al. (\\cite{Krichbaum95}, hereafter K95) presented VLBI observations at 8, 22, 43 and 86 GHz, and also revealed a strongly bent one-sided core-jet structure with at least three superluminal components and two apparently stationary jet components. Their first 86 GHz VLBI observations of this source revealed an unresolved core ($\\le$50$\\mu$as) and a new jet component (N2) whose ejection was believed to relate to an outburst in the mm-regime and a preceding gamma-ray flare observed in early 1993. Britzen et al. (1999, hereafter B99) studied in detail the morphology, kinetics, and the relations between the morphological changes in the source structure and the increased activity in the radio and gamma-ray bands through the eight-year 8 GHz geodetic VLBI observations composed of 20 epochs. They found seven distinct jet components that all showed superluminal motions. The jet trajectory was bent twice, and different components had different paths. But these radio observations cannot give the spectral distribution of VLBI components, because the source is strongly variable, while VLBI observations made at different frequencies were not simultaneous.   In Sect. 2, we present the first quasi-simultaneous VLBI mapping observations of \\object{PKS 0528+134} at five frequencies (2.3, 5.0, 8.4, 15.4, and 22.2 GHz). Then in Sect. 3 we discuss the structure of the jet, give the spectra of the jet components, identify the central core component based on its flat spectrum characteristic, study the proper motions of the jet components, estimate the ejection epochs of two components (a and n2), and describe the jet bending by means of the helical model, followed by the discussion and conclusion in Sects. 4 and 5, respectively.  Throughout this paper, the radio spectral index $\\alpha$ is defined as $S_{\\nu}\\propto\\nu^{-\\alpha}$. By assuming $H_0=71$ kms$^{-1}$ Mpc$^{-1}$, $\\Omega_M=0.27$, and $\\Omega_\\Lambda=0.73$ (Spergel et al. \\cite{Spergel03}), we have a scale for \\object{PKS 0528+134} (z = 2.07) of 1 mas $=8.45 $ pc. ",
        "conclusions": "The five high-resolution VLBI maps made at frequencies 2.3, 5.0, 8.4, 15.4, and 22.2 GHz show that \\object{PKS 0528+134} has a jet structure extending in the north east direction, and the jet bending is clearly seen at about 0.3 and 1.4 mas from the core. Our registration of VLBI components a, f, and g is consistent with the one in B99. We also detect a new VLBI jet component n2, which is likely to be ejected during the radio burst at about 1999.5. We estimate the ejection epoch of component a to be 1991.94 by the linear back-extrapolation of its proper motion, but cannot determine which flux-density burst is truly related to the ejection of component a. No counter-jet is seen in our maps, but its existence cannot be ruled out. We also find that the position angles of the jet components in \\object{PKS 0528+134} are large when they are initially ejected from the core, but gradually decrease with the increasing separation from the core until about 25$\\degr$, which is referred to as the position angle regression.  This first quasi-simultaneous VLBA observations at five frequencies provide a good condition to calculate the spectra of the jet components without taking the variations of the flux densities with the time into account. We identify k component as the core mainly based on its flat spectral index of $0.30\\pm0.08$ and an inverted spectrum peaking at about 6.95 GHz. But we cannot determine which absorption mechanism (SSA, FFA, or others) is mainly responsible for the observed inverted spectrum; future low-frequency observation with sufficient spatial resolution may be useful for distinguishing these possible models. The brightness temperature $T_b$ of each component is given in Table~\\ref{tab:1}; the brightness temperature of core component k does not exceed the equipartition brightness temperature cutoff of $\\sim10^{11}$ K (Readhead \\cite{Readhead94}) if we take the Doppler boosting effect into account. The presence of fast superluminal motions indicates that the relativistic effects are important in this source.  We attempt to describe the trajectory of the jet in \\object{PKS 0528+134} using the helical jet model proposed by S95. The parameters of kinetics obtain from this helical model are consistent with ones from observations. We find from this helical model that there is a reacceleration process of the velocity component $\\beta_z$ along the jet axis and that the jet finally moves nearly along the jet axis, which can be used to explain the position-angle regression. Since the reacceleration process in this helical model cannot explain this big difference in the jet's velocities, we suggest further that a much higher apparent velocity of g component than its inner components' may be due to the relatively higher initial velocity of g component than its inner components. To understand the reacceleraton process well in this helical model needs a good time coverage of high-resolution VLBI observation just as the jet component ejects."
    },
    "0608/astro-ph0608076_arXiv.txt": {
        "abstract": "Observed properties of Wolf-Rayet stars at high metallicity are reviewed. Wolf-Rayet stars are more common at higher metallicity, as a result of stronger mass-loss during earlier evolutionary phases with late WC subtypes signatures of solar metallicity or higher. Similar numbers of early (WC4--7) and late (WC8--9) stars are observed in the Solar neighbourhood, whilst late subtypes dominate at higher metallicities, such as Westerlund~1 in the inner Milky Way and in M83. The observed trend to later WC subtype within metal-rich environments is intimately linked to a metallicity dependence of WR stars, in the sense that strong winds preferentially favour late subtypes. This has relevance to (a) the upper mass limit in metal-rich galaxies such as NGC~3049, due to softer ionizing fluxes from WR stars at high metallicity; (b) evolutionary models including a WR metallicity dependence provide a better match to the observed N(WC)/N(WN) ratio. The latter conclusion partially rests upon the assumption of constant line luminosities for WR stars, yet observations and theoretical atmospheric models reveal higher line fluxes at high metallicity. ",
        "introduction": "Wolf-Rayet (WR) stars represent the final phase in the evolution of very massive stars prior to core-collapse, in which the H-rich envelope has been stripped away via either stellar winds or close binary evolution, revealing products of H-burning (WN sequence) or He-burning (WC sequence) at their surfaces, i.e. He, N or C, O (Crowther 2007).  WR stellar winds are significantly denser than O stars, as illustrated in Fig.~\\ref{WRross}, so their visual spectra are dominated by broad emission lines, notably He\\,{\\sc ii} $\\lambda$4686 (WN stars) and C\\,{\\sc iii} $\\lambda$4647-51, C\\,{\\sc iii} $\\lambda$5696, C\\,{\\sc iv} $\\lambda$5801-12 (WC stars). The spectroscopic signature of WR stars may be seen individually in Local Group galaxies (e.g. Massey \\& Johnson 1998), within knots in local star forming galaxies (e.g. Hadfield \\& Crowther 2006) and in the average rest frame UV spectrum of Lyman Break Galaxies (Shapley et al. 2003).  In the case of a single massive star, the strength of stellar winds during the main sequence and blue supergiant phase scales with the metallicity (Vink et al. 2001). Consequently, one expects a higher threshold for the formation of WR stars at lower metallicity, and indeed the SMC shows a decreased number of WR to O stars than in the Solar Neighbourhood. Alternatively, the H-rich envelope may be removed during the Roche lobe overflow phase of close binary evolution, a process which is not expected to depend upon metallicity.  \\begin{figure} \\centerline{\\psfig{file=test.eps,width=3.0in}} \\caption{Comparisons between stellar radii at Rosseland optical depths of 20 (= $R_{\\ast}$, black) and 2/3 (= $R_{2/3}$, grey) for HD~66811 (O4\\,If), HD~96548 (WN8) and HD~164270 (WC9), shown to scale, together with the wind region corresponding to the primary optical wind line forming region, $10^{11} \\leq n_{e} \\leq 10^{12}$ cm$^{-3}$ (hatched) in each case, illustrating the highly extended winds of WR stars with respect to O stars (Crowther 2007).}\\label{WRross} \\end{figure}  WR stars represent the prime candidates for Type Ib/c core-collapse supernovae and long, soft Gamma Ray Bursts (GRBs).  This is due to their immediate progenitors being associated with young massive stellar populations, compact in nature and deficient in either hydrogen (Type Ib) or both hydrogen and helium (Type Ic).  For the case of GRBs, a number of which have been associated with Type Ic hypernovae (Galama et al. 1998; Hjorth et al. 2003), a rapidly rotating core is a requirement for the collapsar scenario in which the newly formed black hole accretes via an accretion disk (MacFadyen \\& Woosley 1999).  Indeed, WR populations have been observed within local GRB host galaxies (Hammer et al. 2006).  In this review article, the observed properties WR stars at high metallicity are presented and discussed.  \\begin{figure} \\centerline{\\psfig{file=wrpop.eps,width=4.8in}} \\caption{Subtype distribution of Milky Way ($<$3kpc), LMC and SMC WR stars, in which known binaries are shaded (Crowther 2007).}\\label{wrpop} \\end{figure} ",
        "conclusions": "Wolf-Rayet stars are more common at higher metallicity, as a result of stronger mass-loss during earlier evolutionary phases. Late WC subtypes appear to be signatures of solar metallicity or higher, as witnessed within the Westerlund~1 and Quintuplet clusters in the inner Milky Way, M31, M83, NGC~3049, Mrk~309. The observed trend to later WC subtype is intimately linked to a metallicity dependence of WR stars, in the sense that strong winds preferentially favour late subtypes (Crowther et al. 2002). This has relevance to (a) the upper mass limit in metal-rich galaxies, due to softer ionizing fluxes from WR stars at high metallicity (Gonzalez Delgado et al. 2002); (b) evolutionary models including a WR metallicity dependence provide a better match to the observed N(WC)/N(WN) ratio (Eldridge \\& Vink 2006). The latter item relies in part upon the assumption of constant line luminosities for WR stars (e.g. Schaerer \\& Vacca 1998), yet observations and theoretical atmospheric models reveal higher line fluxes at high metallicity (Crowther \\& Hadfield 2006)."
    },
    "0608/astro-ph0608306_arXiv.txt": {
        "abstract": "The trajectory of a primary $\\gamma$-ray detected with an array of at least four atmospheric Cherenkov imaging telescopes can be reconstructed from the shower image centroid positions and geometrical considerations independent of the primary energy. Using only the image centroid positions some cosmic-ray discrimination is also possible. This minimal approach opens the possibility of pushing the analysis threshold to lower values, close to the hardware threshold.  \\vspace{1pc} ",
        "introduction": "The rapid development of ground-based high-energy astronomy during the last ten years \\cite{weekes03} is closely  associated with the success of the Atmospheric Cherenkov Imaging technique \\cite{hillas96} developed with the Whipple Observatory 10m telescope \\cite{weekes89} and elsewhere. Cosmic rays and $\\gamma$-rays entering the atmosphere generate cascades of particles radiating Cherenkov light. At ground level, the Cherenkov light front extends over more than $\\rm \\sim 130m$ from the shower axis. Thus, a single imaging telescope achieves an effective collection area of more than $\\rm 5\\times10^4m^2$. The shape of the image can be used to preferentially select $\\gamma$-ray events over the much more frequent cosmic-ray events. Furthermore, the elongated image of a $\\gamma$-ray shower points back to the source on the sky, providing information on its arrival direction.  Most of the present generation experiments (CANGAROO, HESS, MAGIC and VERITAS) consist of arrays of imaging telescopes. Following earlier pioneering studies \\cite{krennrich98, chadwick96}, the power of the stereoscopic technique \\cite{grindlay75}, was plainly illustrated with the HEGRA experiment \\cite{kohnle96, pulhofer03} and more recently with the HESS experiment \\cite{hinton04} which achieved unprecedented levels of $\\gamma$-ray discrimination.  In these arrays, each telescope records an image of the same shower. For a $\\gamma$-ray primary, the major axis of each image points back to the source so that superposing the images produces the source location on the sky. Furthermore, the position of the image around the source corresponds to the position of the shower around the line from the telescope to the source. Combining this information from two or more telescopes provides the position and orientation of the shower axis. In other words, each telescope can be associated with a plane that contains the shower axis and the telescope. As the shower axis is the line of intersection of the planes associated with each of the array telescopes, only two telescopes are required to extract this axis. Since accurate reconstruction is difficult or impossible when the impact point is closely aligned with the two telescopes, a minimum of three telescopes is preferable.  The stereoscopic technique described above requires that each image contains a sufficient number of pixels for the major axis to be clearly identified. However, this is generally not the case for events close to the detection threshold and for which some or even all of the images can consist of only a few pixels significantly above noise fluctuations. Only the image positions are then available for a reconstructive analysis. In this paper, we study the possibility of reconstructing $\\gamma$-ray events using only the shower image centroid position from each telescope. Such a minimalist analysis might be particularly important for the lowest energy events. The next section describes the method and in the following section, we use Monte-Carlo simulations to present an illustration of the method. Finally, we discuss approaches for discriminating against cosmic-rays when the image positions are the only available significant information. ",
        "conclusions": "By using only the image centroid positions from imaging atmospheric telescopes, the axis of $\\gamma$-ray showers can be reconstructed for arrays of at least four telescopes. For a four-telescope array, we have verified that the analysis provides reasonable precision. Applied to events close to detector threshold, this method can provide a good starting point for more sophisticated and much more computation-demanding methods involving, for example, detailed models of the Cherenkov images \\cite{lebohec,deneroy03}.  We have identified criteria useful for discriminating against cosmic rays while still using only the first-order image parameters, the image centroids and brightnesses. These criteria could be combined with more sophisticated image analysis, such as producing a global image width from the light distribution across the line connecting the image centroid to the reconstructed source position.  Since the analysis presented here uses the minimal amount of information from the recorded Chenenkov images, the results must provide the lowest analysis energy threshold for a given experiment. This analysis, in comparison to analysis techniques using higher-order image parameters, could result in lower energy thresholds for modern air Cherenkov imaging telescope arrays. This technique could also be used in the design of future detectors and suggests that coarser pixels with perhaps smaller light collectors than in HESS and VERITAS could be cost saving and still produce accurate results in future very large arrays."
    },
    "0608/astro-ph0608124_arXiv.txt": {
        "abstract": "We present a study of two spectral lines, \\ion{Fe}{I} 6173 \\AA~and \\ion{Ni}{I} 6768 \\AA, that were candidates to be used in the Helioseismic and Magnetic Imager (HMI) for observing Doppler velocity and the vector magnetic field.   The line profiles were studied using the Mt. Wilson Observatory, the Advanced Stokes Polarimeter and the Kitt Peak McMath telescope and one meter Fourier transform spectrometer atlas.  Both \\ion{Fe}{I} and \\ion{Ni}{I} profiles have clean continua and no blends that threaten instrument performance.   The \\ion{Fe}{I} line is 2\\% deeper, 15\\% narrower and has  a 6\\% smaller equivalent width than the \\ion{Ni}{I} line.  The potential of each spectral line to recover pre-assigned solar conditions is tested using a least-squares minimization technique to fit Milne-Eddington models to tens of thousands of line profiles that have been sampled at five spectral positions across the line.   Overall, the \\ion{Fe}{I} line has a better performance than the \\ion{Ni}{I} line for vector magnetic field retrieval.   Specifically, the \\ion{Fe}{I} line is able to determine field strength, longitudinal and transverse flux four times more accurately than the \\ion{Ni}{I} line in active regions.    Inclination and azimuthal angles can be recovered to $\\approx$2$^\\circ$ above 600 Mx/cm$^2$ for \\ion{Fe}{I} and above 1000 Mx/cm$^2$ for \\ion{Ni}{I}.   Therefore, the \\ion{Fe}{I} line better determines the magnetic field orientation in plage, whereas both lines provide good orientation determination in penumbra and umbra.  We selected the \\ion{Fe}{I} spectral line for use in HMI due to its better performance for magnetic diagnostics while not sacrificing velocity information.  The one exception to the better performance of the \\ion{Fe}{I} line is when high field strengths combine with high velocities to move the spectral line beyond effective sampling range.   The higher $g_{eff}$ of \\ion{Fe}{I} means that its useful range of velocity values in regions of strong magnetic field is smaller than \\ion{Ni}{I}. ",
        "introduction": "\\label{intro}  The Helioseismic and Magnetic Imager (HMI) needs to accurately measure Doppler velocity and vector magnetic field with limited spectral information, sampling roughly half a dozen wavelengths across the spectral line.   HMI is one of a suite of instruments to be included on-board the Solar Dynamics Observatory (SDO) with launch currently scheduled for August of 2008.  The spacecraft is expected to have a large velocity range of $\\pm$4000 m/s.  The first requirement for choosing a spectral line is that it contain a clean continuum with no blends and no near-by lines.  Helioseismology requires that a spectral line be narrow and deep since a steep $dI/d\\lambda$ ensures greater sensitivity to small Doppler shifts.  The \\ion{Ni}{I} 6768 \\AA~line satisfied these requirements and was selected as the spectral line for the Michelson Doppler Imager (MDI) and Global Oscillation Network Group (GONG) instruments.  Vector magnetic field measurements benefit from a high Land\\'e factor, g$_{eff}$, and a simple Zeeman splitting geometry.  It is also important to minimize the number of blends that become apparent in sunspot umbrae where the lower temperatures allow for molecular absorption.  An understanding of the center-to-limb variations of the selected spectral line is important because the accuracy of ``look-up\" algorithms (such as used by MDI) are based on line depth and wing-slope parameters.  In addition, it is desirable for a line to be relatively insensitive to thermodynamic changes so that heights of formation do not drastically change due to moderate perturbations of temperature and density.  Of course, optimum instrumental transmission at the selected wavelength is also a requirement.  The potential of each spectral line to recover pre-assigned solar conditions is tested using a least-squares minimization technique to fit Milne-Eddington models to tens of thousands of line profiles that have been sampled at five spectral positions across the line.  See Figure \\ref{JGFILTERS} for an example of filter profiles applied at five positions across the \\ion{Ni}{I} line.  Similar research was carried out by Graham {\\it et al.}, (2002) to determine the minimum number of spectral filters necessary to recover the vector magnetic field parameters.   The work contained in this paper builds on that premise insofar as we examine the differences inherent in two spectral lines in their capacity to recover the vector magnetic field with filter-polarimetric measurements. ",
        "conclusions": "Both \\ion{Fe}{I} 6173 \\AA~and \\ion{Ni}{I} 6768 \\AA~profiles have clean continuum and no blends which threaten performance (see Figures~\\ref{6173w} and~\\ref{6173asp}).  The \\ion{Ni}{I} parameters of line width, line depth and wing slope show smooth and predictable changes as a function of center-to-limb angle (see Figure~\\ref{cl6768}). \\ion{Ni}{I} has been carefully studied by Bruls (1990) who concluded the line was a good choice for helioseismology because the line profile is stable and not very sensitive to variation of temperature and temperature gradient in the photosphere.  Extensive studies of this nature, including effects of granulation on the line formation, are yet to be carried out for \\ion{Fe}{I}.  Not surprisingly, the higher $g_{eff}$ \\ion{Fe}{I} serves as a much better diagnostic tool than \\ion{Ni}{I} to ascertain the inclination angle, field strength and filling factor of the magnetic field.  This is a robust result  as it is obtained repeatedly using a variety of analysis techniques.  First, the ASP observations of \\ion{Fe}{I} and \\ion{Ni}{I} compared to the ASP observations of the benchmark  \\ion{Fe}{I} 6301 \\AA~/ 6301 \\AA~ line pair show \\ion{Ni}{I} to have much greater uncertainty in recovering the field strength, filling factor and inclination  (see Figure~\\ref{aspff}).  Second, the simulated line profile data shows \\ion{Fe}{I} to perform nearly a factor of two better than the \\ion{Ni}{I} line for field strength and inclination determination (see Figure~\\ref{JGSIM1}).  Third, results using observed profiles show that \\ion{Fe}{I} is able to determine field strength, longitudinal and transverse flux four times more accurately than \\ion{Ni}{I} in active regions.    Inversions using observed profiles show inclination and azimuthal errors to be recovered to $\\approx$2$^\\circ$ degrees above 600 Mx/cm$^2$ for \\ion{Fe}{I} and above 1000 Mx/cm$^2$ for \\ion{Ni}{I} (see Figures~\\ref{JGASP1} and \\ref{JGASP3}).  This means, using the \\ion{Fe}{I} line allows better magnetic field orientation determination in plage, whereas both lines will provide a good orientation determination in penumbra and umbra.    From the above-mentioned results, we predict acceptable performance for \\ion{Fe}{I} 6173 \\AA~for vector magnetic field determination.  In recovering velocities, the performance of the spectral lines from inversion methods shows both the \\ion{Fe}{I} and \\ion{Ni}{I} lines perform well for low field strengths and low velocities, with errors of roughly 15 m/s.  However, the higher $g_{eff}$ of \\ion{Fe}{I} means that its operational range of velocity values in regions of strong magnetic field is smaller than \\ion{Ni}{I}.  \\ion{Fe}{I} performs poorly at high field strengths and high velocities where the line has moved beyond the filter sampling range.  This implies that in order to have accurate velocity measurements in umbrae using the \\ion{Fe}{I} line, the line sampling must be increased by either increasing the number of filter positions or chasing the line by changing the filter positions as a function of spacecraft velocity.  Simulations using only the Stokes $I$ and $V$ profiles shows that the velocity errors increase for both spectral lines in umbra if Stokes $Q$ and $U$ are not considered (see Figure~\\ref{JGSIMIV}  as compared to Figure~\\ref{JGSIM1}) .  This result implies that active region helioseismology accuracy will improve if the full Stokes vector is used for analysis.  Using the observed profile performance for \\ion{Fe}{I} as a good indicator for its performance on-board the HMI spacecraft, we conclude that it will perform well for $F_L$ and $F_T$ in general, for field directions in all regions with more than $600$~Mx/cm$^2$, and for magnetic field strength in penumbra and umbra."
    },
    "0608/astro-ph0608548_arXiv.txt": {
        "abstract": "We present the results of a search for new members of the Taurus star-forming region using the Infrared Array Camera (IRAC) aboard the {\\it Spitzer Space Telescope}. With IRAC images of 29.7~deg$^2$ of Taurus at 3.6, 4.5, 5.8, and 8.0~\\micron, we have identified sources with red mid-infrared colors indicative of disk-bearing objects and have obtained optical and infrared spectra of 23 of these candidate members. Through this work, we have discovered 13 new members of Taurus, two of which have spectral types later than M6 and thus are likely to be brown dwarfs according to the theoretical evolutionary models of Chabrier and Baraffe. This survey indicates that the previous census of Taurus has a completeness of $\\sim80$\\% for members with disks. The new members that we have found do not significantly modify the previously measured distributions of Taurus members as a function of position, mass, and extinction. For instance, we find no evidence for a population of highly reddened brown dwarfs ($A_K\\sim2$) that has been missed by previous optical and near-infrared surveys, which suggests that brown dwarf disks are not significantly more flared than disks around stars. In addition to the new members, we also present IRAC photometry for the 149 previously known members that appear within this survey, which includes 27 objects later than M6. ",
        "introduction": "\\label{sec:intro}  As one of the nearest sites of active star formation ($d=140$~pc), the Taurus complex of dark clouds has long been a popular laboratory for studies of the birth of stars and planets. A fundamental prerequisite for most observations of Taurus is a census of its stellar content. Over the decades, a variety of methods have been employed to identify members of Taurus, including objective prism spectroscopy of H$\\alpha$ \\citep{joy46,joy49,bri93,bri99}, X-ray imaging with the {\\it Einstein Observatory} \\citep{fei87,wal88} and the {\\it R\\\"ontgen Satellite} \\citep{ss94,neu95,wic96,bri99}, proper motion measurements \\citep{jh79,har91,gom92}, near- and mid-infrared (IR) photometry from ground-based telescopes \\citep{gom94,itg96,lr98} and the {\\it Infrared Astronomical Satellite (IRAS)} \\citep{bei86,ken90,ken94}, and wide-field optical and near-IR imaging \\citep{bri98,bri02,luh00tau,luh04tau,luh06tau,luh03tau,mar01,gui06}. The census of known members of Taurus has steadily grown \\citep{ck79,hb88,kh95} and now contains more than 300 stars and brown dwarfs.   The {\\it Spitzer Space Telescope} \\citep{wer04} offers a combination of sensitivity and large field of view that is unprecedented for a mid-IR telescope, making it uniquely suited for finding disk-bearing objects down to low masses ($M\\sim0.01$~$M_\\odot$), through high levels of extinction ($A_V\\sim100$), and across large areas of sky ($\\theta\\sim10\\arcdeg$). To take advantage of this capability, we have used archival {\\it Spitzer} images encompassing 29.7~deg$^2$ of Taurus to perform a search for new members of the region. In this paper, we present the collection and analysis of the {\\it Spitzer} data (\\S~\\ref{sec:irac}), the identification of candidate members with these data (\\S~\\ref{sec:cand}), the spectroscopy and classification of these candidates (\\S~\\ref{sec:spec}), an evaluation of the completeness of this survey and previous ones (\\S~\\ref{sec:complete}), and the implications of our updated census for the distribution of Taurus members as a function of position, mass, and extinction and the implications of our mid-IR photometric catalog for the frequency of disks in Taurus (\\S~\\ref{sec:imp}). ",
        "conclusions": "We have presented the results of a mid-IR imaging survey of 29.7~deg$^2$ of the Taurus star-forming region that was performed with IRAC aboard the {\\it Spitzer Space Telescope}. We used these data to search for new members of Taurus by identifying objects with colors indicative of circumstellar disks. Through spectroscopy of 23 of these candidates, we have discovered 13 new members. One of these objects appears to be a class~I source based on its IRAC colors, while two of the new members have spectral types later than M6 and thus are likely to be brown dwarfs. The number of new members found in this survey indicates that the previous census of Taurus exhibits a completeness of $\\sim80$\\% for class~I and II sources. Meanwhile, by considering the number of previously known members of Taurus that we have recovered, we estimate that our survey has a completeness of $\\sim94$\\% for class~I and II members. This completeness extends to high levels of extinction, ranging from $A_K\\sim2$-10 for $M\\sim0.02$-2~$M_\\odot$. If brown dwarf disks are more flared than disks around stars, than the average extinction toward brown dwarfs should be higher than toward stars, but we have found no evidence of this in our survey. The new members that we have found do not significantly alter the distributions of Taurus members as a function of position, mass, and extinction that have been measured in previous membership surveys.  In addition to performing a search for new members of Taurus, we have also measured mid-IR photometry for the 149 previously known members of Taurus within the IRAC survey, which comprises nearly half of the known membership. With these data, we have measured disk fractions of $61\\pm8$\\% and $40\\pm13$\\% for the stellar ($\\leq$M6) and substellar ($>$M6) members of Taurus.  The disk fraction for stars in Taurus is higher than the fraction for stellar members of IC~348 and Chamaeleon~I, which is consistent with the slightly younger age of Taurus implied by Hertzsprung-Russell diagrams of these populations. Meanwhile, the difference in disk fractions for stars and brown dwarfs in Taurus has only marginal statistical significance and may be a reflection of the incompleteness for class~I brown dwarfs in surveys of this region to date. Indeed, $\\sim500$ candidate members of Taurus from this IRAC survey remain unobserved with spectroscopy, most of which would be class~I brown dwarfs if confirmed, although a vast majority are probably red galaxies."
    },
    "0608/astro-ph0608581_arXiv.txt": {
        "abstract": "{This paper presents new five passbands ($UBVRI$) optical wide-field imaging data accumulated as part of the DEEP Public Survey (DPS) carried out as a public survey by the ESO Imaging Survey (EIS) project. Out of the 3~square degrees originally proposed, the survey covers {\\tt 2.75} square degrees, in at least one band ( normally $R$), and 1.00 square degrees in five passbands. The median seeing, as measured in the final stacked images, is $0\\farcs97$, ranging from $0\\farcs75$ {\\tt to} $2\\farcs0$. The median limiting magnitudes (AB system, 2\\arcsec aperture, $5\\sigma$ detection limit) are $U_{AB}=25.65$, $B_{AB}=$25.54, $V_{AB}=$25.18, $R_{AB} =$ 24.8 and $I_{AB} =$24.12~mag, consistent with those proposed in the original survey design. The paper describes the observations and data reduction using the EIS Data Reduction System and its associated EIS/MVM library. The quality of the individual images were inspected, bad images discarded and the remaining used to produce final image stacks in each passband, from which sources have been extracted. Finally, the scientific quality of these final images and associated catalogs was assessed qualitatively by visual inspection and quantitatively by comparison of statistical measures derived from these data with those of other authors as well as model predictions, and from direct comparison with the results obtained from the reduction of the same dataset using an independent (hands-on) software system. Finally to illustrate one application of this survey, the results of a preliminary effort to identify sub-mJy radio sources are reported. To the limiting magnitude reached in the $R$ and $I$ passbands the success rate ranges from 66 to 81\\% (depending on the fields).  These data are publicly available at CDS\\thanks{ Available at CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/}. ",
        "introduction": "\\label{sec:introduction}  Deep multi-wavelength observations are a pre-condition for the selection of suitable galaxy samples for spectroscopic follow-up observation using the multiplex capability of spectrographs mounted on large-aperture telescopes. Foreseeing the need for such samples for the scientific exploitation of multi-object spectrographs such as FORS and VIMOS, ESO's Survey Working Group (SWG) recommended the ESO Imaging Survey (EIS, \\citet{RenzinidaCosta97}) project to undertake a series of deep, optical/infrared observations. Based on ideas submitted by the ESO community and evaluated by the SWG, the DEEP Public Survey (DPS) was designed to provide a data set that could be used to study the large-scale distribution of galaxies at $z\\sim1$ and to identify a large number of Ly-$\\alpha$ galaxies and,  combined with near-infrared data, extremely red objects (EROs) for follow-up work using VLT.  In its entirety this survey comprises three optical and infrared strategies. The optical part, the focus of the present paper, consists of a deep survey covering three regions of 1~square degree on the sky, in the $U$-, $B$-, $V$-, $R$- and $I$ passbands, using the wide-field imager (WFI) mounted on the ESO/MPG 2.2m telescope at La Silla. Each of the three DEEP regions --~1,~2,~3~-- is covered by four WFI pointings (in five passbands) --~a, b, c, d. The regions were selected both to enable observations year-round and because they overlapped with regions of other scientific interest. For instance, DEEP1 was chosen to complement  the sub-mJy ATESP radio survey \\citep{Prandoni2000a,Prandoni2000b,Prandoni2005} carried out with the Australia Telescope Compact Array (ATCA) covering the region surveyed by the ESO Slice Project \\citep{Vettolani1997}, while DEEP2 included the CDF-S field. Finally, DEEP3, was chosen in the northern galactic hemisphere, thus providing an almost year-round coverage\\footnote{The text for the original proposal can be found at http://www.eso.org/science/eis/documents.html}. The location and the characteristics of the surveyed regions as well as the planned limiting magnitudes in each passband can be found in the DPS strategy page\\footnote{http://www.eso.org/science/eis/surveys/strategy\\_DPS.html} available in the EIS web-pages.  This paper is part of a series describing the data releases made by the EIS project. In \\cite{Dietrich2005} (hereafter Paper~I) the data reduction system and the general procedures adopted are described. Here the observations, reduction, and science verification of the optical data taken as part of the DPS, publicly released late 2004, are reviewed. Furthermore, in this paper changes to the dataset that have occurred since the original release, stemming from additional evaluation of the data, are also described.  Complementing the optical results reported in the present paper, the results obtained from infrared observations covering these fields are presented in \\cite{Olsen2006}. Sect.~\\ref{sec:observations} describes optical observations while the reduction and calibration of the data are presented in Sect.~\\ref{sec:reduction}. Final survey products such as stacked images and science-grade catalogs extracted from them are presented in Sect.~\\ref{sec:products}. The quality of these products is evaluated in Sect.~\\ref{sec:discussion} by comparing statistical measures obtained from these data with those of other authors. In this section the results of a preliminary assessment of the optical identification of sub-mJy radio sources from the ATESP 1.4 GHz radio survey are also discussed. Finally, in Sect.~\\ref{sec:summary} a brief summary of the paper is presented. ",
        "conclusions": "\\subsection{Comparison of counts and colors}  A key element in public surveys is to provide potential users with information regarding the quality of the products released. To this end a number of checks of the data are carried out and several diagnostics plots summarizing the results are automatically produced by the EIS Survey System. They are an integral part of the product logs available from the survey release page. Due to the large number of plots produced in the verification process these are not reproduced here.  Instead, to illustrate the results only a small set of these plots are shown.  A relatively simple statistics that can be used to check the catalogs and the star/galaxy separation criteria is to compare the star and galaxy number counts derived from the data to that of other authors and/or to model predictions. For stars we use Vega magnitudes and for galaxies AB magnitydes, while for both we use the MAG\\_AUTO magnitudes from the catalog. Figs.~\\ref{fig:gal_deep1}--\\ref{fig:gal_deep3} in appendix~\\ref{app:galcounts} show for each region a mosaic where the rows represent the passbands from $U$(top) to $I$ (bottom) and the columns the different fields available from a (left) to d (right). Here objects with CLASS\\_STAR$<$0.95 were considered to be galaxies.  Note that the number counts shown in the figure take into account the effective covered area, which is available in the {\\tt FIELDS} table (see appendix~D of Paper~I). Inspection of all figures shows that the galaxy counts obtained from the final produced stacks are, in general, in excellent agreement with those obtained by previous authors, \\citep{Arnouts2001,Metcalfe2001}, regardless of the passband. The figures also highlight some of the challenges in dealing with mosaics, as the completeness of the counts can be seen to vary from field to field.  A complementary test is to compare the stellar counts to those predicted by models such as the galactic model of \\cite{Girardi2005}. Figs.~\\ref{fig:starcounts1}--\\ref{fig:starcounts3} in appendix~\\ref{app:starcounts} show these comparisons, in the same format as the galaxy counts, for each of the observed regions. The good agreement between the data and the model predictions is indeed remarkable, especially considering the uncertainties in the zero-point and the model. These results give further support to the photometric calibration and the scientific quality of the science-grade data produced automatically by the EIS system.  While useful to detect gross errors, number counts are not sufficiently sensitive to identify more subtle differences. The comparison of expected colors of stars with theoretical models provides a better test of the accuracy of the photometric calibration in the different bands. Using the model of \\cite{Girardi2005}, the predicted colors of stars can be obtained and compared to the data to evaluate at least the rough consistency of the photometric calibration in different passbands.  Such comparisons were made for each of the four fields covered in all five passbands and are shown in Figs.~\\ref{fig:starcol1}--\\ref{fig:starcol4} presented in appendix~\\ref{app:colorcolor}. From these figures one finds that at least for the fields with complete color information there is a fairly good agreement between the locus of stars and those predicted by the model in four different color-color diagrams, indicating a good ($\\lesssim 0.1~mag$) consistency among the photometric calibration obtained in different passbands.   \\void{ Overall, the offsets from the predicted colors are of the order that one can expect for automatic calibrations as those done here and using the standard calibration plan of ESO.\\\\ The $U$- and $B$-bands are the ones for which the largest offsets, up to 0.2~mag in the worst case (DEEP2b), are observed. The redder bands, in particular $R$ and $I$, have much smaller offsets which in all cases are less than 0.1~mag and in many cases negligible. These results are similar to those obtained for the XMM fields \\citep{Dietrich2005} and seem to indicate that the automatic calibration and the standard calibration plan are better adapted to the red than the blue bands. The previous comparison of the expected and measured colors of the stars, points out that in the case of long-term projects like those carried out by EIS, it is important to have a careful calibration plan if a high-precision absolute calibration of the data is required. The most efficient way would probably be to assure, for each field and passband combination, calibration data taken in a perfectly photometric night. Such observations had to be accompanied by a thorough set of standard star observations with a good coverage in both airmass and color in order to obtain precise photometric calibrations. }  \\void{ In the current project such data were not obtained, as reflected by the small zero-point offsets.}   \\subsection{Comparison with other reductions}  The data presented here have been reduced using the same method and parameter settings as those discussed in Paper~I. In that paper the results of the reduction through the EIS data reduction system and that using the GaBoDS pipeline \\citep{Schirmer2003,Erben2005} developed in Bonn were compared. Extensive and repeated comparisons were made between SExtractor-produced catalogs from the images generated by these two independent systems.  Initial discrepancies, resulting from the different techniques used (e.g. cosmic ray removal, gain-harmonization) were resolved, leading at the end to results in excellent agreement.  A similar comparison has been partially carried out for the present dataset and it has shown an important short-coming of the current implementation of the EIS/MVM library. Even though the chip-to-chip gain correction is not mandatory, the log file created by the library provides no information of whether or not it has been applied.  In fact, comparison of a subset of frames showed that in the original reductions made public late 2004 the images for DEEP2 were not homogenized. Since in the final, sky-subtracted image, residual differences are not seen this was only detected from the comparison with the independent reduction conducted in Bonn (see also \\citealt{Hildebrandt}) .  It is important to emphasize that the images for DEEP2 have been re-reduced and have replaced those released in October 2004. The same is true for the associated science-grade catalogs.   \\void{ One of the important implications of this is that the log file produced by the EIS/MVM programs needs to be improved. For instance, from the comparison it became evident that in one of the regions (DEEP2) the correction for the chip-to-chip variations had not been applied. However, the log file does not report whether or not this correction is applied, and without proper reporting it becomes difficult after sky subtraction to devise a procedure to check for this in a systematic way.}    \\subsection {Radio/Optical correlation} \\label{sec:atesp}  \\begin{figure}[t] \\resizebox{\\hsize}{!} {\\includegraphics{5030fig2.eps}} \\caption{ Distribution of separations bewtween radio sources and the nearest $R$-detected sources for DEEP1a (upper panel), DEEP1b (middle panel) and DEEP1c (lower panel).  The dashed histograms show those considered as actual counterparts.} \\label{fig:atid} \\end{figure}   \\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics{5030fig3.eps}} \\caption{Radio-optical offset computed for all the optical counterparts in the fields  DEEP1 a,b,c.  The 2 arcsec circle defines the boundaries of those considered to be ``good'' identifications. Also reported are the values of the median of the offset distribution in right ascension and declination.} \\label{fig:offset} \\end{figure}   As mentioned earlier, the DEEP1 region was chosen to complement the sub-mJy ATESP 1.4 GHz radio survey carried out with the Australia Telescope Compact Array (ATCA). The ATESP survey covers a strip of $26\\times 1$ square~degrees, at declination $\\delta\\sim-40^\\circ$, and has provided a catalogue of $\\sim3000$ radio sources to a flux limit of $S_{1.4 GHz}\\sim0.4$ mJy \\citep{Prandoni2000a,Prandoni2000b}.  The aim of the ATESP radio survey is to study the evolutionary properties of the objects belonging to the mJy and sub-mJy populations, by directly measuring the evolution of star-forming radio sources, in the redshift range $0<z<1$, and of low-luminosity AGNs up to higher redshifts.  In the $4\\times0.25$~square degrees of the DEEP1 region covered by DPS there are 109 ATESP 1.4 GHz radio sources - 25 in DEEP1a, 26 in DEEP1b, 29 in DEEP1c and 29 in DEEP1d. Since the latter field was not covered by the present survey, the analysis below refers to only 80 radio sources. To identify optical counterparts of these radio sources the nearest optical object was searched for. All the available single passband catalogs for the fields DEEP1a, b and c were treated separately.  Fig.~\\ref{fig:atid} shows the radio-optical distance distribution of the $R$-selected optical counterparts of the ATESP radio sources for DEEP1a, b and c, respectively. As can be seen most optical counterparts are found within 2~arcsec. Beyond this distance, the distribution is flat, as expected for spurious identifications. Therefore, in this preliminary analysis this value is taken to be the maximum distance from a radio source for an R-band object to be considered its optical counterpart. The distribution of offsets for all the identified radio source-optical pairs is shown in Fig.~\\ref{fig:offset}. It is worth noticing that systematic offsets in the relative position between the radio and optical sources are only $-0\\farcs073$ in declination and $-0\\farcs213$ in right ascension (see median values in Fig.~\\ref{fig:offset}), and have therefore been neglected. Also note that matches obtained in the $R$-band are the most complete, since all the radio sources identified in any of the $U$-, $B$-, $V$- or $I$-bands are also identified in $R$-band.   The ATESP source identification rate varies from 66\\% to 81\\%, depending on the field, while the estimated contamination  is 10-14\\%. These identification rates are consistent with the value of 74\\% reported for the VVDS-VLA sample \\citep{Ciliegi2005}, where a radio/optical analysis of a sample of $\\mu$Jy sources was performed down to a similar optical depth. It is also interesting to compare these results to those obtained using a shallower survey, such as the $I-$band EIS-WIDE survey \\citep[Patch A, ][]{Nonino1999}. In the 3~square~degrees region that this patch covers, a radio/optical analysis was carried out, yielding the identification of $\\sim57\\%$ ATESP radio sources down to $I=22.5$, the limiting magnitude of the optical imaging \\citep{Prandoni2001}. By comparing this identification rate to the one obtained with the DPS data, one finds that there is a significant increase, demonstrating the need for deep follow-up surveys for properly identifying the mJy/sub-mJy population.   The multi-band DPS data was used to analyze the color properties of the ATESP radio sources. Fig.~\\ref{fig:atcol} shows in the ($R-I$)$\\times$ ($B-R$) color diagram the location of the optical objects associated to the radio sources (filled circles), compared to the colors of field galaxies (dots). Also plotted are the evolutionary tracks of various optical galaxy populations (from early-type to spiral to starburst galaxies) in the redshift range $0<z<3$.  From this plot it is clear that while the majority of the ATESP sources follows the early-type galaxy tracks (dotted and dashed lines) up to redshift $\\sim2$,  some ATESP sources have colors consistent with being spiral or starburst galaxies (long-dashed, dot-dashed and solid lines) up to z$\\sim0.5-1$. These results are consistent with those of previous work \\citep[e.g. ][]{Prandoni2001} and the most likely interpretation is that classical radio galaxies, usually associated to early-type objects (ellipticals and/or S0) are, in general, more powerful and can be seen to larger redshifts than radio sources associated to star-forming galaxies.  To dramatically increase the identification rate of the ATESP sources, one requires deep near-infrared ($J$, $H$, $K$) imaging. In this respect, the infrared part of DPS \\citep[see ][]{Olsen2006} is very useful, since it provides $J$ and $K$ images down to $K_{AB}=21.3$ for the DEEP1a and b fields. Such data should allow the identification of the most distant radio galaxies (reddened by redshift effects) and possible ``rare'' intrinsically red objects (like EROs, ULIRGs, etc), probably reddened by dust.  The preliminary radio/optical analysis described above will be further developed in a future paper (Mignano {\\it et~al.} in prep.), where the multi-color information will be used to determine spectral types and (photometric) redshifts for the identified radio sources. We notice that, to fully exploit the capabilities of the optical data described above, the entire DEEP1 $2^{\\circ}\\times 0.5^{\\circ}$ region has been imaged also at 5~GHz, again using the ATCA. The 5~GHz data \\citep{Prandoni2005} give the possibility of measuring the radio spectral index (between 1.4 GHz and 5 GHz) of the radio sources. This additional information may help constraining the origin (star formation or nuclear activity) of the radio emission in mJy and sub-mJy radio sources and in understanding whether it is linked to the optical light.   \\begin{figure*}[t] \\resizebox{12cm}{!}{\\includegraphics{5030fig4.eps}} \\hfill \\parbox[b]{55mm}{ \\caption{Color-Color Diagram. In the figure is plotted the $(R-I) \\times (B-R)$ colors for the whole optical population (dots) and for the optical counterparts of ATESP radio sources (filled circles). Dot, dashed, long-dashed, dot-dashed and solid lines represent the evolutionary tracks of bursts (single burst of star formation followed by passive evolution), ellipticals, Sb and Sc spirals, and starburst galaxies, respectively.  The evolutionary tracks were computed at intervals of 0.2 in redshift using the GISSEL libraries \\citep{BruzualCharlot93}. The model tracks span the redshift range $0.0<z<3.0$.} \\label{fig:atcol}} \\end{figure*}    This paper presents the survey products prepared by the EIS team for the optical part of the Deep Public Survey (DPS) conducted over 4 years by the team.  The survey was designed to cover 3 independent strips of the sky with an area of 1~square degree each in five passbands. The paper describes the observations and presents a comprehensive discussion of the set of products prepared and validated by the EIS Data Reduction System.  These products are publicly available including fully calibrated reduced images taken on a nightly basis, final stacked images and extracted source catalogs. It should be noted that the present optical data are complemented by observations in infrared reported  by \\cite{Olsen2006}."
    },
    "0608/astro-ph0608254_arXiv.txt": {
        "abstract": "{} {We present near ultraviolet (NUV:1750 - 2800\\AA) and far ultraviolet (FUV: 1350 - 1750\\AA) light-curves for flares on 4 nearby dMe-type stars (GJ 3685A, CR Dra, AF Psc and SDSS J084425.9+513830.5) observed with the $\\it GALEX$ satellite.} {Taking advantage of the time-tagged events recorded with the $\\it GALEX$ photon counting detectors, we present high temporal resolution ($<$ 0.01 sec) analysis of these UV flare data.} {A statistical analysis of 700 seconds of pre-flare quiescence data for both CR Dra and SDSS J084425.9+513830.5 failed to reveal the presence of significant micro-flare activity in time bins of 0.2, 1 and 10 second intervals. Using an appropriate differential emission measure for both the quiescent and flaring state, it is possible to reproduce the observed FUV:NUV flux ratios. A major determinant in reproducing this flux ratio is found to be the value of plasma electron density during the flare. We also searched the count rate data recorded during each of the four flare events for periodicity associated with magneto-hydrodynamic oscillations in the active region coronal loops. Significant oscillations were detected during the flare events observed on all 4 stars, with periodicities found in the 30 to 40 second range. Flare oscillations with this periodicity can be explained as acoustic waves in a coronal loop of length of $\\approx 10^{9}$ cm for an assumed plasma temperature of $5-20 \\times 10^{6}$K. This suggests a loop length for these M-dwarf flares of less than $1/10^{th}$ of the stellar radii. We believe that this is the first detection of non-solar coronal loop flare oscillations observed at ultraviolet wavelengths.} {} ",
        "introduction": "Although flare stars were detected early in the last century, it was not until the late 1940's that flare star research attracted wide-spread attention with the observation by Joy $\\&$ Humason (\\cite{joy49}) of a four magnitude eruptive event. Since then there have been many hundreds of publications concerning these active stellar sources. The first coincidence between optical and radio activity was reported by Lovell $\\&$Solomon (\\cite{Lovell66}), with the first X-ray flare being observed by Heise et al. (\\cite{Heise75}). Surveys show that low mass M-dwarf (dMe) stars account for more than 75\\% of the stellar population in the solar neighborhood and, due to the strong magnetic fields that cover most of their stellar disks, many of them exhibit significant levels of coronal activity. In a recent survey of 8000 late-type dwarfs contained in the Sloan Digital Sky Survey (SDSS) catalog, West et al. (\\cite{west04}) have found that $>$ 50\\% of stars of spectral class M4 to M9 have high levels of magnetic activity, as inferred by the presence of the H$\\alpha$ emission line in their stellar spectra. The most easily observable manifestation of this activity is that of stellar flare eruptions, as typified by the star UV Ceti, which involves random outbursts in which the short-term stellar brightness increases significantly on time-scales of seconds to hours at X-ray, UV, visible and radio wavelengths (Schmitt et al. \\cite{schmitt93}; Phillips et al. \\cite{phillips88}; Stepanov et al. \\cite{stepanov95}). Although the detailed physics associated with the formation of stellar flares is still uncertain (Haisch, Strong $\\&$ Redono  \\cite{haisch91}), it is clear that these eruptions are linked to the magnetic heating processes occurring in stellar coronae. We note that the soft X-ray emission from flares (which is indicative of gas with temperatures of $>$ 10$^{6}$K) is dominated by coronal line-emission, while the UV radiation during flare events originates in the chromosphere and transition region and is characterized by emission from a gas with a temperature of $\\sim$ $10^{5}$K. There is still much debate as to the relative contributions from emission line and/or UV continuum flux to the total heating budget in flare stars (Robinson et al. \\cite{robinson01}; Hawley et al.\\cite{hawley03}; G\\\"{u}del et al.\\cite{gudel03}). \\begin{table*} \\begin{center} \\caption{Summary of $\\it GALEX$ flare star observational parameters}  \\begin{tabular}{lllccc} \\hline \\hline Star Name&R.A.(J2000)&Dec.(J2000)& Survey Mode&MAST Identifier&Exposure (sec)\\\\ \\hline GJ 3685A&11:47:41&+00:15:12&MIS&MISDR1-13062-0283-0003&1244\\\\ AF Psc&23:31:44&-02:45:12&GI*&GI1-067026-PGC71626-0002&1705\\\\ CR Dra&16:17:05&+55:16:09&DIS&ELAISN1-07-0004&1650\\\\ SDSS J084425&08:44:26&+51:38:31&MIS&MISDR1-03333-0447-0001&1669\\\\ \\hline \\multicolumn{6}{l}{* pointed Guest Investigator observation} \\\\ \\hline \\hline \\end{tabular} \\end{center} \\end{table*} Recent EUV and X-ray observations suggest that flares statistically contribute a significant fraction of the overall (X-ray) heating of the coronae on all magnetically active stars (G\\\"{u}del et al. \\cite{gudel03}; Arzner $\\&$ G\\\"{u}del \\cite{arzner04}). In particular, the recent $\\it XMM-Newton$ satellite observations of dMe star flares by Mitra-Kraev et al. (\\cite{mitra05a}) have noted a correlation between the observed ultraviolet (UV) flare energy and the corresponding increase in X-ray luminosity, with the UV flares temporally preceding the peaks in X-ray flux. These observations would seem to support the predictions for the Neupert Effect (Neupert \\cite{neupert68}), as outlined in the chromospheric evaporation model for solar flares (Antonucci, Gabriel $\\&$ Dennis \\cite{anton84}). In this picture, during the impulsive phase of a stellar flare, the UV/optical emission due to accelerated electrons that gyrate downward along the magnetic field lines and impact the chromosphere, should precede the more slowly evolving X-rays that are emitted by the heated plasma.  Although flare events have been routinely recorded on many M-dwarfs at radio(Jackson, Kundu $\\&$ White \\cite{jackson89}), optical (Gunn et al. \\cite{gunn94}), Ultra-violet (Robinson et al. \\cite{robinson01}), Extreme Ultra-violet (Audard et al. \\cite{audard00}) and X-ray (Marino, Micela $\\&$ Peres \\cite{marino00}) wavelengths, such observations have generally suffered from both poor temporal-resolution and the inability to track a large event throughout the entire pre-flare to post-flare period. However, Beskin et al. (\\cite{beskin88}) recorded emission structure on time-scales of 0.3 -- 0.8 seconds in high-temporal (0.3 milli-second) U-band visible observations of flare events on several bright UV Ceti-type stars. Using the scale height of the atmosphere and the velocity of the shock led them to the conclusion that optical flares on M-dwarfs were due to thermal phenomena. We note that U-band observations of both solar and stellar flares have been widely used as a proxy for non-thermal hard X-ray flare emission (Neidig \\cite{neidig89}, Hawley et al. \\cite{hawley95}).  High ($<$ 1 sec) temporal resolution observations of ultraviolet flares can provide detailed information on microflaring activity and also determine whether wave activity is present. Both of these are thought to be important in the heating of coronae (Robinson, Carpenter $\\&$ Percival \\cite{robinson99};  Roberts \\cite{roberts04}). Furthermore, the ability to record emission throughout an entire energetic flare event can provide insights into the nature of the pre-flare and post-flare quiescent phases and how their statistical small-scale flux variations may be related to the onset of a major flare event (G\\\"{u}del et al. \\cite{gudel04}). The NASA $\\it GALEX$ satellite (Martin et al. \\cite{mar05}), launched in 2003, has recently been shown to be an excellent observational platform for the serendipitous detection and high time-resolved observation of stellar flares in two ultraviolet (UV) photometric bands. For example, an extremely large flare event was observed on the dM4e star GJ 3685A on April 20th, 2004 that revealed an increase in stellar brightness of $>$ 10 UV magnitudes in a period of $<$ 200 seconds (Robsinson et al. \\cite{robinson05}). These UV observations provided compelling evidence for two distinct classes of flares during this eruption, each characterized by significantly different ratios of their measured near UV (NUV) and far UV (FUV) fluxes. In this Paper we re-visit those data (which were analyzed under the assumption of a blackbody energy distribution) and in addition, we report on three further flare events on the stars CR Dra, AF Psc and SDSS J084425.9+513830.5 (hereafter referred to as SDSS J084425) that were serendipitously recorded with a high time resolution of $\\sim$ 0.01 sec by the $\\it GALEX$ satellite during its all-sky survey observations during the period 2004 - 2005. \\begin{figure*}[hbt] \\vspace{18cm} \\special{psfile=Fig1a.ps hoffset=5 voffset=530 hscale=32 vscale=32 angle=-90} \\special{psfile=Fig1b.ps hoffset=5 voffset=355 hscale=32 vscale=32 angle=-90} \\special{psfile=Fig1c.ps hoffset=5 voffset=180 hscale=32 vscale=32 angle=-90} \\special{psfile=Fig1d.ps hoffset=240 voffset=530 hscale=32 vscale=32 angle=-90} \\special{psfile=Fig1e.ps hoffset=240 voffset=355 hscale=32 vscale=32 angle=-90} \\special{psfile=Fig1f.ps hoffset=240 voffset=180 hscale=32 vscale=32 angle=-90}  \\vspace*{-0.5cm}  \\caption{The calibrated flare light-curves for AF Psc, CR Dra, GJ 3685A and SDSS J084425 for both the FUV and NUV bands.} \\label{flares} \\end{figure*}   \\begin{figure*}[hbt] \\vspace{7cm} \\special{psfile=Fig2a.ps hoffset=-10 voffset=210 hscale=34 vscale=34 angle=-90} \\special{psfile=Fig2b.ps hoffset=240 voffset=210 hscale=34 vscale=34 angle=-90}  \\vspace*{-1cm}  \\caption{The $\\it GALEX$ FUV:NUV flux ratios for the two flare stars AF Psc and GJ 3685A.} \\label{flares} \\end{figure*}  \\begin{figure*}[hbt] \\vspace{17cm} \\special{psfile=Fig3a.eps hoffset=-20 voffset=60 hscale=50 vscale=60 angle=0} \\special{psfile=Fig3b.eps hoffset=240 voffset=60 hscale=50 vscale=60 angle=0}  \\vspace*{-2.5cm}  \\caption{Pre-flare quiescent photon data for CR Dra (left) and SDSS J0084425 (right) binned in 0.2, 1.0 and 10.0 second intervals, shown together with the 3-$\\sigma$ statistical error limits (dashed lines). No evidence is found for micro-flare activity.} \\label{flares} \\end{figure*} ",
        "conclusions": "We have presented near and far ultraviolet light-curves for stellar flares observed with the NASA $\\it GALEX$ satellite on the four nearby dMe-type stars GJ 3685A, CR Dra, AF Psc and SDSS J084425. These data were recorded as time-tagged UV photon events with a temporal resolution of $<$ 0.01 seconds, thus enabling a detailed time-analysis of both their quiescent and flaring states. During 700 seconds of quiescence, prior to the flare outbursts on both CR Dra and SDSS J084425, no statistical evidence was found for micro-flare outbursts in the NUV with data binned in 0.2, 1.0 and 10.0 second time-intervals.  We have used a modified differential emission model (DEM) curve of the flare star EV Lac in conjunction with the CHIANTI atomic database to produce theoretical  DEM curves for a typical dMe star in both quiescence and a flaring state. From the quiescent DEM distribution we were able to reproduce FUV:NUV flux ratios greater than unity, consistent with the measurements during the flares. A critical input parameter in determining the FUV:NUV ratio is the value of plasma electron density during the flare event. It was also found that the major single emitter in the FUV filter is C~{\\sc iv}, although weaker emission lines and the continuum provide $\\sim$50\\% of the flux. In the NUV filter, a significant contribution comes from upper chromosphere/lower transition region emission lines, in addition to continuum emission. Several coronal lines are present in both filters, with $\\sim$10\\% contribution in the NUV filter, whose main contributer is from continuum processes.  We have also searched for periodicities in the observed UV count-rate data recorded during the four flare events to search for coronal loop oscillations. Using both the wavelet and Fourier transform analysis methods on these flare data, we have found compelling statistically significant evidence for oscillations in both of the NUV and FUV data for flares on AF Psc and GJ 3685A, and in the NUV data for the flares on CR Dra and SDSS J084425 with periodicities in 30-40 second range. Assuming typical values of electron pressure during a flare, these periodicities may be due to acoustic waves in coronal loops of length $\\sim$ $10{^9}$ cm for a plasma with a temperature range of $5-20 \\times 10^6$ K. These values suggest a loop length of less than $1/10^{th}$ of the M-dwarf stellar radii. Finally we note that although oscillations have been previously observed on the Sun and in X-ray observations of the M-type dwarf AT Mic, we believe that this is the first detection of non-solar coronal loop flare oscillations at ultraviolet wavelengths."
    },
    "0608/astro-ph0608704_arXiv.txt": {
        "abstract": "% We have obtained high resolution (R = 10,000) K-band spectra of candidate young massive stars deeply embedded in high-mass star-forming regions. These objects were selected from a near-infrared survey of 44 regions of high-mass star-formation (\\cite{Kaper}). In these clusters, 38 OB stars are identified whose K-band spectra are dominated by photospheric emission.  In almost all those stars, the K-band spectra are indistinguishable from field stars.  However, in some stars the profile of the Br$\\gamma$ line is different (less deep, or absent) from those of the O field stars. One of the explanations of these profiles might be an enhanced mass-loss. ",
        "introduction": "The advent of high-quality near-infrared instrumentation has opened up a new window on the birth sites of massive stars. The formation time scales of a massive star is short, of the order of 100,000 years. This means that the newly born stars are still deeply embedded in their parental molecular cloud. Not much is known about the onset of the stellar wind or the initial rotation properties of young massive stars. Occurring behind large amounts of extinction, inside their natal molecular cloud, these evolutionary stages can only be observed at near-infrared wavelengths.   We have identified 38 OB type stars inside massive star forming regions (\\cite{Ostarpaper}) using K-band spectroscopy. They are classified based on the classification scheme of \\cite{hanson96}. Three O3-O4 stars are detected, while only a few O3 stars are known in our galaxy. The K-band spectra of the OB stars are in general similar to those of the (more evolved) OB stars located in OB associations and the field. This suggest that the K-band spectral properties of the very young OB stars are already indistinguishable from those of the more evolved stars. ",
        "conclusions": ""
    },
    "0608/astro-ph0608197_arXiv.txt": {
        "abstract": "We have obtained spectra of 102 red giant branch (RGB) stars in the Leo~I dwarf spheroidal galaxy (dSph) and 74 RGB stars in the Leo~II dSph using the Low Resolution Spectrograph (LRIS) on the Keck~I 10-meter Telescope.  We report on the calcium abundances [Ca/H] derived from the strengths of the Ca~II triplet absorption lines at 8498, 8542, 8662~\\AA~in the stellar spectra using a new Ca~II triplet calibration to [Ca/H].  The two galaxies have different average [Ca/H] values of $-1.34 \\pm 0.02$ for Leo~I and $-1.65 \\pm 0.02$ for Leo~II with intrinsic abundance dispersions of 1.2 and 1.0~dex, respectively. The typical random and total errors in derived abundances are 0.10 and 0.17~dex per star. For comparison to existing literature, we also converted our Ca~II measurements to [Fe/H] on the scale of Carretta and Gratton (1997) though we discuss why this may not be the best determinant of metallicity; Leo~I has a mean [Fe/H] $\\rm = -1.34$ and Leo~II has a mean [Fe/H] $\\rm = -1.59$. The metallicity distribution function of Leo~I is approximately Gaussian in shape with an excess at the metal-rich end, while that of Leo~II shows an abrupt cutoff at the metal-rich end. The lower mean metallicity of Leo~II is consistent with the fact that it has a lower luminosity, hence lower total mass, than Leo~I; thus the evolution of Leo II may have been affected more by mass lost in galactic winds. We compare our observed metallicity distribution functions to those inferred by Dolphin (2002) from modelling of deep color-magnitude diagrams (CMDs) obtained with the Hubble Space Telescope (HST). The models predict a median metallicity that is systematically more metal-rich than observed, by 0.36~dex in the Leo~I dSph, and by 0.49~dex in the Leo~II dSph. We show that differences this large are not likely caused by non-solar [Ca/Fe] in the dSph stars.  Our direct and independent measurement of the metallicity distributions in these dSph will allow a more accurate star-formation histories to be derived from future analysis of their CMDs. ",
        "introduction": "The goal of this project is to map the chemical abundance distribution of the stellar populations of the Leo~I and Leo~II dwarf spheroidal (dSph) galaxies in order to constrain the physical processes that regulate their evolution. Table~\\ref{table-dsphs} lists the properties of each galaxy for reference.  Dwarf spheroidal galaxies (dSphs) are low luminosity ($M_V \\le -14$ mag), low surface brightness ($\\Sigma_V \\simlt 22$ mag/arcsec$^2$), and low mass ($M_{tot} \\simlt 10^7 {\\rm M}_\\odot$) systems.  They have little or no interstellar gas ($M_{HI} \\simlt 10^5 {\\rm M}_\\odot$), and are dark-matter dominated.  For recent reviews, see Mateo (1998), Smecker-Hane \\& McWilliam (1999), and Grebel (2000).  Their stellar velocity dispersions are typically $\\sim 10$ km/sec, and therefore galactic winds from the first generation of supernovae are expected to efficiently rid them of gas and prevent extended or multiple epochs of star formation.  Although they seem to lack the mass and material to sustain star formation for a significant fraction of the Hubble time, the CMDs of the dSph satellites of the Milky Way have shown that most dSphs have {\\it very complex} star formation histories (SFHs) and chemical evolution. Cold dark matter hierarchical galaxy formation models have difficulty explaining extended epochs of star formation in dSphs. They are predicted to begin forming stars at early epochs, but lose a significant amount of gas due to photoevaporation during reionization, which curtails their star formation activity (Barkana \\& Loeb 1999). Klypin et al.~(1999) also points out that CDM models vastly over predict the number of low mass halos; the Milky Way and M31 have only one-tenth the number of low mass satellites. Are these ``missing\" satellites dark because they never accreted gas, or were the dark matter halos stripped of their baryons through supernovae-driven winds? Determining the age of the oldest stars in the outer halo satellites such as the Leo~I and Leo~II dSphs, whose Galactocentric distances are 270 and 204~kpc, respectively, is important because CDM models predict that these should form later than the Milk Way's inner halo globular clusters and dSphs.  Almost every dSph in the Local Group appears to have had a unique SFH.  From analysis of a CMD obtained with WFPC2 aboard the HST, Mighell \\& Rich (1996) have shown that the Leo~II dSph formed essentially all of its stars from $\\sim 7$ to 14~Gyr ago with 50\\% of its stars formed in a 4~Gyr period. In striking contrast, most of the stars in the Leo~I dSph are younger than this; 87\\% of its star formation occurred from 1 to 7~Gyr ago and only 12\\% occurred $\\simgt 10$~Gyr ago (Gallart et al.~1999a and 1999b).  More recent photometry from Held \\etal (2000), which includes a much larger radial coverage than the HST data, also shows an extended blue horizontal-branch in Leo~I indicating the existence of an old, metal-poor population.  Held \\etal (2001) also report the discovery of candidate RR~Lyrae stars in the system, further strengthening the case for an old population in Leo~I.  Defining the chemical evolution of the dSphs will be key to understanding their complex SFHs because the chemical abundances constrain the inflow and outflow of gas from the systems. Did some dSphs accrete fresh gas that fuelled subsequent star formation? Were outflows driven by supernovae so extensive that today's dSphs have only a small fraction of their former mass? The chemical abundances of dSph stars -- the mean metallicity as well as the detailed distribution in metallicity  -- can be determined from spectroscopy of individual red giant stars in the dSphs. Measuring the metallicity distribution of individual stars will greatly improve the SFHs derived from analysis of CMDs.  In order to understand the true spread in stellar chemical abundances more precise abundance indicators than can be inferred from CMD analysis are needed.  We, therefore, measured the abundance distribution of the Leo~I and Leo~II dSphs from spectroscopy of individual red giant stars using the Ca~II absorption lines in the near infrared (8498, 8542, and 8662~\\AA). We utilized the new calibration of Ca~II to [Ca/H] from Bosler (2004, hereafter B04 \\footnote{ftp://sculptor.ps.uci.edu/pub/smecker/bosler\\_thesis.pdf}) to derive abundances with random and total errors of 0.10 and 0.17~dex, respectively. Additionally, we calculated the [Fe/H] values of the stars based on earlier empirical calibrations between Ca~II and [Fe/H] simply for comparison with existing work using this method; however, we discuss the pitfalls of using the Ca~II triplet to determine [Fe/H]. The new Ca~II calibration to [Ca/H] is completely independent of star to star variations in [Ca/Fe] unlike previously utilized Ca~II abundance calibrations. ",
        "conclusions": "We have determined calcium abundances for 102 RGB stars in the Leo~I dSph galaxy and 74 RGB stars in the Leo~II dSph galaxy using the equivalent widths of the near infrared Ca~II triplet lines. The metallicity scale used to convert the Ca~II measurements is based on our calibration of Ca~II to [Ca/H] from Galactic cluster stars with ages from $\\approx$ 2 to 14~Gyr and $-2.2$ \\ltsima [Ca/H] \\ltsima $+0.2$. Unlike the calibration of Ca~II to [Fe/H] developed by R97b, the Ca~II to [Ca/H] is independent of star-to-stars variations in [Ca/Fe] with [Fe/H]. Defying the tradition of using ``metallicity\" interchangeably with iron abundances, [Ca/H] was taken as the relevant metallicity in this investigation.  The Leo~I dSph has a average metallicity of [Ca/H] = $-1.34 \\pm 0.02~ (\\sigma = 0.21)$ and an intrinsic spread in metallicity of 1.2~dex. The average metallicity of the Leo~II dSph is more metal-poor with [Ca/H] = $-1.65 \\pm 0.02~ (\\sigma = 0.17)$ and has a 1.0~dex spread in metallicity.  The metallicity distribution function is different for the two dSphs. Leo~I has an approximately Gaussian distribution with a slight excess at the metal-rich end, while Leo II has a steeper cutoff at the metal-rich end at [Ca/H] $\\approx -1.4$. The differences in the shape of the distributions may be caused by greater mass loss in the less massive Leo~II dSph, whose star formation ended approximately 7~Gyr ago, compared to the more massive Leo~I dSph, whose star formation ended only $\\sim 1$ to 2~Gyr ago. Significant differences, on the order of $-$0.5~dex, exist between the observed metallicity distribution functions and those derived by sophisticated modelling of Hess diagrams in CMDs obtained with HST. However, the assumption of a Galactic range of [Ca/Fe] values show our results to be consistent with the [Fe/H] values derived in other spectroscopic and photometric studies."
    },
    "0608/astro-ph0608474_arXiv.txt": {
        "abstract": "We analyze the fundamental plane projections of elliptical galaxies as a function of luminosity, using a sample of $\\approx 80,000$ galaxies drawn from Data Release 4 (DR4) of the Sloan Digital Sky Survey (SDSS).  We separate brightest cluster galaxies (BCGs) from our main sample and reanalyze their photometry due to a problem with the default pipeline sky subtraction for BCGs.  The observables we consider are effective radius ($R_e$), velocity dispersion ($\\sigma$), dynamical mass ($M_{\\rm dyn} \\propto R_e\\sigma^2$), effective density ($\\sigma^2/R_{\\!e}^2$), and effective surface brightness ($\\mu_e$).  With the exception of the $L$-$M_{\\rm dyn}$ correlation, we find evidence of variations in the slope (i.e. the power-law index) of the fundamental plane projections with luminosity for our normal elliptical galaxy population. In particular, the radius-luminosity and Faber-Jackson relations are steeper at high luminosity relative to low luminosity, and the more luminous ellipticals become progressively less dense and have lower surface brightnesses than lower luminosity ellipticals. These variations can be understood as arising from differing formation histories, with more luminous galaxies having less dissipation. Data from the literature and our reanalysis of BCGs show that BCGs have radius-luminosity and Faber-Jackson relations steeper than the brightest non-BCG ellipticals in our sample, consistent with significant growth of BCGs via dissipationless mergers.  The variations in slope we find in the Faber-Jackson relation of non-BCGs are qualitatively similar to that reported in the black hole mass-velocity dispersion ($M_{\\rm BH}$-$\\sigma$) correlation.  This similarity is consistent with a roughly constant value of $M_{\\rm BH}/M_\\star$ over a wide range of early type galaxies, where $M_\\star$ is the stellar mass. ",
        "introduction": "Early-type galaxies are observed to populate a tight plane -- the fundamental plane (FP) -- in the space defined by their effective radii, velocity dispersions, and surface brightnesses (\\citealt{Dressler87,DD87,BBF92,Burstein97}). Such a correlation is expected theoretically if early-type galaxies are in virial equilibrium.  The observed plane is, however, tilted with respect to the simplest virial theorem expectation, which must be accounted for in galaxy formation models. This tilt could be caused by a variation in the dynamical mass-to-light ratio for elliptical galaxies as a result of a varying dark matter fraction (e.g., \\citealt{Padmanabhan04,BMQ05}) or stellar population variations (e.g., \\citealt{Gerhard01}). Another explanation for the FP tilt could be non-homology in the surface brightness profiles of elliptical galaxies (e.g., \\citealt{GC97,TBB04}).  Projections of the FP contain additional information about the properties of early type galaxies. These include, e.g., the Faber-Jackson (FJ) relation \\citep{FJ76} between velocity dispersion and luminosity and the radius-luminosity relation.  Although the FP itself is quite tight over a wide range of early type galaxies (e.g., \\citealt{Bernardi03c}), there are a number of indications of variations in the FP projections themselves with galaxy luminosity (which we refer to as ``variations in slope'' or ``curvature'' in the FP projections). For example, the slope of the Faber-Jackson relation decreases in lower luminosity ellipticals \\citep{Tonry81,Davies83,MG05}.  In addition, \\citet{OH91} find that while brightest cluster galaxies (BCGs) lie on the same FP as lower mass ellipticals, their central velocity dispersions are roughly constant with increasing luminosity and their effective radii increase steeply with luminosity, significantly more so than for normal ellipticals. \\citet{Lauer06a} and \\citet{Bernardi06} have confirmed these differences in the FP projections of BCGs with larger data sets. These trends can be understood if dissipationless (gas-poor) mergers become increasingly important for luminous ellipticals \\citep{BMQ06}.  In addition to global correlations among the galaxy properties themselves, the masses of central black holes in early-type galaxies correlate well with the velocity dispersions \\citep{FM00,Gebhardt00,Ferrarese01,Tremaine02}, bulge concentrations \\citep{Graham01}, and stellar masses \\citep{Magorrian98,MH03,HR04} of their host galaxies. \\citet{Wyithe06} and \\citet{GH06} have argued that the $M_{\\rm BH}-\\sigma$ relation steepens at high black hole masses and flattens at low masses, respectively. This is similar to the luminosity dependence in the FJ relation noted above. If supported by larger data sets, the similar luminosity dependence of the $M_{\\rm BH}-\\sigma$ and FJ relations could have important implications for understanding the relationship between the formation of galaxies and their central massive black holes.  The purpose of the present paper is to quantify the variation in the FP projections with luminosity using the exquisite statistics made possible by the Sloan Digital Sky Survey (SDSS) Data Release 4 (DR4; \\citealt{sdss_dr4}). These variations have, until now, been largely overlooked by previous FP studies, because simple power-law assumptions were often made when fitting the data. This can be seen, for instance, in the sample of $\\sim9000$ galaxies from \\citet{Bernardi03b}, where hints of the variations we discuss are quite noticeable in the plots. In the next section we describe the data used in our analysis (\\S \\ref{sec:data}). We also describe a problem in the standard pipeline SDSS photometry of BCGs that leads us to reanalyze the BCGs separately from the bulk of the ``normal'' elliptical galaxy population. Having defined our sample, in \\S \\ref{sec:FP} we examine the observed FP projections and quantify curvature in these projections as a function of luminosity. We also use published photometry of nearby (non-SDSS) BCGs to discuss the relationship between our results and the properties of BCGs.  In \\S \\ref{sec:disc} we discuss our results and their connection to galaxy-BH correlations. Throughout this paper, we use a standard cosmology with $(\\Omega_M,\\Omega_\\Lambda,h)=(0.3,0.7,0.7)$. ",
        "conclusions": "\\label{sec:disc}  We have analyzed the fundamental plane projections of elliptical galaxies as a function of luminosity, using a sample of $\\approx 80,000$ galaxies drawn from DR4 of the SDSS.  We have separated BCGs from our main elliptical galaxy sample and reanalyzed their photometry. The observables we consider are effective radius ($R_e$), velocity dispersion ($\\sigma$), dynamical mass ($M_{\\rm dyn} \\propto R_e\\sigma^2$), effective density ($\\sigma^2/R_{\\!e}^2$), and effective surface brightness ($\\mu_e$). With the exception of the $L$-$M_{\\rm dyn}$ correlation, we find clear evidence of variations in the slope of the FP projections with luminosity (Figs. \\ref{fig:allslopes} \\& \\ref{fig:quad}) in the normal elliptical galaxy sample. The trends we find are that the radius-luminosity and Faber-Jackson relations are steeper at high luminosity relative to low luminosity and that more luminous ellipticals become progressively less dense and have lower surface brightnesses than lower luminosity ellipticals.  These trends are consistent with existing results in the literature, though our results have much better statistics.  It is interesting to note that the results we find appear to continue to even lower luminosity spheroidal systems. In particular, for a sample of dwarf elliptical (dE) and dwarf spheroidal (dSph) galaxies in the {\\it B} band, the radius-luminosity power-law slope $\\alpha \\approx$ 0.28--0.55 and the Faber-Jackson power-law slope $\\beta \\approx$ 1.5--2.5 \\citep{deRijcke05}, broadly consistent with an extrapolation of our results to even lower luminosities. These trends are also seen in the dE sample of \\citet{GG03}. We note, however, that there are likely structural differences between luminous ellipticals and dEs that complicate this simple extrapolation.  The variations in the FP projections we find at the luminous end are consistent with less and less dissipation during the formation of elliptical galaxies with increasing luminosity.  Less dissipation would result in increasingly larger, less dense, lower surface brightness, and lower $\\sigma$ galaxies, relative to fiducial power-law scalings for the elliptical galaxy population.  This interpretation is consistent with a variety of other observational evidence for less dissipation in the formation of luminous ellipticals (e.g., \\citealt{KB96}).  Luminosity (or galaxy-mass) dependent variations in the FP projections should in principle provide a strong constraint on galaxy formation models.  For example, numerical simulations of gas-free (``dry'') merger remnants find that the FP projections have $\\alpha \\sim 0.7$ and $\\beta \\sim 4$ for very wide orbits, while lower angular momentum and higher energy orbits have less energy exchange between stars and dark matter and thus the FP projections steepen to $\\alpha \\sim 1$ and $\\beta \\sim 6$ \\citep{BMQ06}.  These values of $\\alpha$ and $\\beta$ are consistent with the trends we find in the data at the luminous end, suggesting an increasing importance of dry mergers for luminous ellipticals.  For lower-luminosity ellipticals, however, dry mergers are less likely to be important in determining the properties of elliptical galaxies.  Current simulations of elliptical galaxy formation via mergers of gas-rich disks find that the remnants lie roughly on the FP and its projections, with the gas fraction of the progenitor galaxies strongly affecting the properties of the resulting spheroidal merger remnants \\citep{Robertson06}.  These simulations also find more dissipation in the formation of lower luminosity ellipticals, which is consistent with our data.  It remains to be seen whether simulations of gas-rich mergers can also account for the systematic variations in the FP projections with luminosity (e.g., with a varying gas fraction).  We have also reanalyzed the photometry of SDSS BCGs using the C4 cluster sample of \\citet{Miller05}; the standard pipeline photometry underestimates the luminosities and radii for luminous BCGs (Figure~\\ref{fig:RevsRe}) due to an overestimate of sky background.  With our corrected photometry, we find that the radius-luminosity relation of BCGs is noticeably steeper than that of normal (non-BCG) ellipticals and the FJ relation for BCGs flattens at the luminous end. The same trends are also seen in local samples of BCGs (eg., \\citealt{Lauer06a}).  These observational results are consistent with significant growth of BCGs by dissipationless mergers.  \\subsection{Massive Black Holes}  For the bulk of the elliptical galaxy population, both the $M_{\\rm BH}$-$\\sigma$ and Faber-Jackson relations scale roughly as $\\sigma^4$. As noted in \\S 1, however, \\citet{Wyithe06} argued that the $M_{\\rm BH}$-$\\sigma$ relation steepens at high black hole masses and flattens at low masses, albeit with only modest statistical significance.  \\citet{GH06} see a similar flattening at low black hole masses and \\citet{Lauer06a} argue for a similar steepening at high black hole masses based on the properties of the surface brightness cores in luminous elliptical galaxies.  These luminosity-dependent variations in the $M_{\\rm BH}$-$\\sigma$ relation are qualitatively similar to the luminosity dependent variation in the FJ relation. Together with a constant $M_\\star/L$, these results imply a roughly constant value of $M_{\\rm BH}/M_\\star$ across a range of early type galaxies (see, e.g., \\citealt{HR04} for a direct compilation of $M_{\\rm BH}/M_\\star$ that bears this out).  For massive galaxies, the approximate constancy of $M_{\\rm BH}/M_\\star$ may be a consequence of dissipationless mergers, which preserve $M_{\\rm BH}/M_\\star$ in the absence of energy loss by gravitational waves during black hole coalescence \\citep{BMQ06}.  For lower luminosity ellipticals, however, gas dynamics is important and a constant value of $M_{\\rm BH}/M_\\star$ implies a constant relative efficiency for forming stars and massive black holes in a given gravitational potential well."
    },
    "0608/astro-ph0608642_arXiv.txt": {
        "abstract": " a population of $10^{2.5} M_\\odot$ PBHs making up $0.1$\\% of the dark matter and growing through accretion to incorporate $10$\\% of the dark matter by the present day (we mark the mass of the seed PBH only, going on the conservative assumption that the PBH's dark matter halo is not directly accreted).  Both these points lie in a region of acceptable parameter space for the explanation of ULXs with IMBHs, but the exact extent of the region is difficult to define given the uncertainties mentioned above.  \\subsection{PBHs as SMBH Seeds}  The question of the origin of supermassive black holes at high redshifts has attracted a great deal of attention since quasars have been discovered at redshifts $z>6$, implying that black holes as massive as $\\sim 10^9 M_\\odot$ \\citep{Fan2003, Barthetal2003,Willottetal2003} exist when the universe is less than 1 Gyr old.  It has proven difficult to find a mechanism that can create such massive black holes so quickly.  Most proposals require smaller black holes to act as seeds for the build-up of SMBHs.  In some cases, these seeds form directly from the collapse of halos \\citep{LodatoNatarajan2006,SpaansSilk2006}.  In others, the seeds are the remnants of Population III stars \\citep{Shapiro2005,VR}, or they might be primordial. The question of whether or not SMBHs can be grown from the merging of stellar mass black holes has also been discussed in recent work \\citep[see][and references therein]{Lietal2006}. While in each scenario a case may be made for the ability of these seeds to result in the SMBHs we observe as quasars, sometimes requiring the invocation of self-interacting dark matter \\citep{Ostriker2000,Huetal2006} or the accretion of scalar fields \\citep{BeanMagueijo2002}, there is as yet no consensus on the matter.  We suggest that PBHs may be a viable SMBH seed candidate because of their ability to build up large dark matter halos that may assist in further growth through baryon accretion later on. Furthermore, forming earlier than Pop III stars, they have more time to grow in the epoch before quasars are observed.   \\subsection{Conclusions} \\label{sec:concl}  We have shown that primordial black holes can grow significantly after formation through the accretion of a dark matter halo.  In the radiation era this can lead to an increase of total mass on order unity, while during matter domination, the mass can grow by roughly two orders of magnitude.  Although a dark matter halo may not significantly increase the mass of the seed black hole itself due to the lack of a mechanism to dissipate angular momentum, the accumulation of a halo cannot be ignored when considering the ability of a PBH to accrete gas in later eras (in the case of collisional dark matter \\citep{Ostriker2000}, the PBH can directly accrete significant amounts of dark matter).  We have also shown that the parameter space available to PBHs as components of dark matter components is sufficient for them to play an interesting role in galaxies.  PBHs may be viable candidates for the seeds of intermediate mass black holes, possibly responsible for ultra-luminous x-ray sources, or they may play a role in seeding supermassive black holes currently found in the centers of galaxies.",
        "introduction": "\\label{sec:intro}  It is well-established that supermassive black holes (SMBHs) with masses in the range $m_{BH} \\sim 10^6 - 10^{9.5} M_{\\odot}$ reside in the centers of spheroidal systems \\citep{Bernardi2003}.  One can make a convincing case that these have grown largely through accretion, with the consequent energy emission observed in electromagnetic output and jets at an efficiency of $\\epsilon \\sim 0.1$ \\citep{Soltan,YuTremaine}. Observations of distant quasars have shown us that these SMBHs are already in place by redshifts of 6 and greater, but the mechanism of their formation remains a mystery.  Motivated by this question, we examine the potential for primordial black holes (PBHs) to grow through accretion to become seed masses for SMBHs.  Primordial black holes, defined as black holes forming in the early universe without stellar progenitors, were first proposed by \\citet{ZN1967} and \\citet{Hawking1971} as a possible consequence of the extremely high densities achieved in the Big Bang model. If they do indeed form in the early universe and can avoid evaporation~\\citep{Hawking1975} up to the present day, they must still exist and they may be important.  We use a combination of analytical and numerical methods to follow PBH growth through the radiation- and matter-dominated eras and show that a PBH can multiply its mass by up to two orders of magnitude through the accretion of a dark matter halo.  If PBHs can grow sufficiently by accretion (or if they are very large at birth) they may account for intermediate-mass black holes (IMBHs) in the mass range $10^2 M_\\odot \\lesssim M \\lesssim 10^4 M_\\odot$, which have been suggested as the engines behind ultra-luminous x-ray sources (ULXs) recently discovered in nearby galactic disks \\citep{Dewangan2005,Madhusudhanetal2006,Miller2004,Mushotzky2004}. Based on the observed ULX luminosities of $\\sim 10^{39} \\textrm{erg s}^{-1}$, stellar mass black holes are ruled out unless the emission is highly beamed.  IMBHs have the appropriate Eddington luminosity to explain ULXs, but since there is currently no easy way to produce black holes of this mass from stellar collapse at the abundances observed \\citep{FryerKalogera2001}, their origins are highly debated.  We suggest that PBHs may be able to grow to sufficient masses through the accretion of dark matter halos to account for a pervasive population of IMBHs.  These IMBHs may also be important in the build-up of SMBHs, as suggested by recent numerical studies \\citep{MicicAbelSigurdsson2005}, a scenario that may be testable with the Laser Interferometer Space Antenna (LISA) in coming years \\citep{Micic2,Clifford2004,Fregeau2006}.  Aside from observational indications, there have been physics-based inquiries indicating the plausibility of the production of PBHs through a variety of mechanisms in the early universe \\citep{Carr2005}, as briefly reviewed in \\S \\ref{sec:PBH}.  One may wish to ask at this point what constraints on PBH production exist given current observational limits.  This important question is addressed in \\S~\\ref{sec:PBHobslims} and Figure~\\ref{constraints}.  We will show that there is ample room for a population of PBHs that is both permitted and interesting.  Our study of PBH accretion is organized as follows. In \\S \\ref{sec:PBH} we discuss theories of PBH formation.  In \\S \\ref{sec:acc} we outline our accretion model. Sections \\ref{sec:rad} and \\ref{sec:matt} describe the accretion calculations in the radiation and matter eras, respectively, while in \\S \\ref{sec:comb} we do a combined calculation for both eras.  In \\S \\ref{sec:results} we give our results for the total accretion possible for a PBH, and in \\S \\ref{sec:discussion} we discuss the implications of our findings for the possible importance of PBHs in the present-day universe.  Appendices \\ref{sec:apprad} and \\ref{sec:appmatt} present details of our accretion calculations. ",
        "conclusions": "\\label{sec:discussion}  Our results suggest that PBHs could grow significantly after their formation by acquiring a dark matter halo and the resulting clothed black holes could make up an interesting fraction of the dark matter.  In the following discussion, we show that current observations are not in conflict with this conclusion, and in fact there is ample room both observationally and theoretically for PBHs to play a role in the universe today.  \\subsection{Observational Limits on PBHs} \\label{sec:PBHobslims}  We include in Figure~\\ref{constraints} a plot of the current observational limits on PBHs over a wide range of masses and dark matter fractions.  Here, we describe the limits illustrated in the plot.  Most limits on PBHs in the present-day universe are derived from considerations of PBHs as dark matter candidates.  PBHs massive enough to escape evaporation would certainly qualify as ``cold'' and ``dark'' matter; however, their existence could have noticeable effects on processes from nucleosynthesis to galaxy formation.  If PBHs form at very early times, and thus with very low masses ($\\lesssim10^{10}$ g), they could interfere with nucleosynthesis by emitting particles during their evaporation \\citep{Kohri2000}.  The abundance of PBHs forming after nucleosynthesis is constrained by measurements of the baryon fraction of the universe \\citep{Novikov1979}.  Massive PBHs can also be constrained by dynamical considerations in the low-redshift universe.  Some of the strongest current constraints are derived from the wealth of data from microlensing searches in Galaxy \\citep{Alcock2000, Alcock2001, Afonso2003}.  The frequency and character of observed microlensing events constrains black holes in the mass range 0.1 - 1 $M_\\odot$ to make up less than $\\sim 20$\\% of the dark matter in the Galactic halo \\citep{Alcock2000,Gould2005}.  The limits from microlensing are shown on Figure~\\ref{constraints} labeled ``MACHO'' and ``EROS.''  For larger masses, constraints on PBH dark matter can be found by examining the effect of PBHs on the matter power spectrum. Calculating the excess in the power spectrum that Poisson-distributed PBHs would contribute, \\citet{AMS2003} find an upper limit on present-day PBH mass of a few times $10^4 M_\\odot$, which is consistent with but tighter than the previous constraint at $10^6 M_\\odot$ based on the heating of the Galactic disk \\citep{LaceyOstriker}.  However, we note that this limit assumes all the dark matter is in PBHs: $\\Omega_{BH} = \\Omega_{DM}$.  If this assumption is relaxed, we find that a wide range of masses can be accommodated at lower density parameters.  Specifically, we find that the product of the initial density parameter and the initial mass are constrained by \\be \\label{eqn:AMS} m_{BH} \\Omega_{BH} < \\frac{m_{AMS}}{x} \\ee where $x$ is the factor by which the PBHs increase in mass via accretion and $m_{AMS}$ is $\\sim 10^4 M_\\odot$.  We include this limit in the constraint plot (Figure~\\ref{constraints}) with the label ``AMS.''  Other limits can be placed by considering the effects of compact objects along the line of sight lensing more distant sources \\citep{Wambsganss2002}.  \\citet{Dalcantonetal1994} search for the slight amplification in the continuum emission of quasars that would be expected were black holes to cross the line of sight during an observation.  With a large sample of observations, they are able to place limits on black holes in the range $\\sim 10^{-3} M_\\odot$ to $\\sim 300 M_\\odot$.  This limit is included in the constraints plot (Figure~\\ref{constraints}) and labeled ``QSO.'' \\citet{Wilkinsonetal2001} place limits on the abundance of massive black holes in the universe based on their predicted effect of creating multiple images of compact radio sources.  Studying a sample of 300 sources, they find a null result and from that can place a constraint on the density of black holes along the line of sight.  Their constraint is included in Figure~\\ref{constraints} with the label ``RADIO.''  Coincident with the compact radio source study, another group found a similar constraint by searching for the same lensing effect in gamma ray burst light curves \\citep{Nemiroffetal2001}.  The results of the two studies are consistent with each other, so for simplicity we include only the \\citet{Wilkinsonetal2001} result in Figure~\\ref{constraints}.  Finally, a limit on black holes with masses of $\\sim 10 M_\\odot$ and up can be placed by observing widely orbiting binary systems in the Galaxy \\citep{YooChanameGould2004}.  If a compact object passes between the two companion stars in a binary, the orbits of the stars will be perturbed.  \\citet{YooChanameGould2004} use this to estimate how many compact objects with the ability to disturb a binary system exist in the halo.  This constraint is included in Figure~\\ref{constraints} and labeled ``WB.''  We point out that none of the above observations significantly constrain black holes making up less than $\\sim$ 10\\% of the dark matter for a wide range of masses.  \\begin{figure}[htb] \\begin{center} \\resizebox{\\columnwidth}{!}{\\includegraphics{f5.eps}} \\end{center} \\caption{\\label{constraints}Observational constraints on black holes in the Galactic halo (see \\S \\ref{sec:PBHobslims}) from microlensing experiments (EROS and MACHO), quasar variability studies (QSO), compact radio source lensing (RADIO), the stability of wide binaries (WB), the high-wavenumber matter power spectrum (AMS), and the heating of the Galactic disk (DISK).  The region labeled ``PopIII remnants'' represents a rough estimate of the region of parameter space relevant to a scenario in which Population III star remnant IMBHs are responsible for ULX observations.  The star symbol gives the position of the example of a PBH ULX used in the abstract (see \\S \\ref{sec:ULX}).} \\end{figure}  \\subsection{PBHs as ULXs}\\label{sec:ULX}  This part of parameter space is consistent with an interpretation of the recent ULX observations as accreting, intermediate-mass black holes in nearby galaxies.  In this scenario, ULXs occur when IMBHs residing in the galaxy's halo pass through molecular clouds in the disk \\citep{Mapellietal2006,Miller2004,Winteretal2006}. The enhanced density in a molecular cloud is sufficient to trigger gas accretion, which causes the IMBH to emit X-ray radiation.  These sources would be transient, and in any given galaxy the number of sources detected at any given time would depend on the number and distribution of the IMBH population and the fraction of the disk made up of molecular clouds.  Several recent papers explore the possibilities for making ULXs with IMBHs.  \\citet{MiiTotani2005} estimate the number of ULXs expected if they are the result of IMBHs passing through molecular clouds.  Using the mass and dark matter fraction estimated for IMBHs if they are the compact remnants of Population III stars \\citep{MadauSilk2005}, $M_{IMBH} \\sim 10^2 - 10^3$ and $\\Omega_{IMBH}/\\Omega_b \\sim 0.1$, they find that the estimated number of ULXs is consistent with observations.  These results do not depend on the nature of the IMBHs -- if the IMBHs were primordial in nature rather than Pop III star remnants, they would also be capable of producing the observed sources.  Although these results are encouraging, the present lack of understanding of the nature of ULXs and the many uncertainties that go into predictions of the consequences of a halo population of IMBHs mean that the issue is far from resolved.  A recent paper drawing on the Mii \\& Totani result uses an ensemble of N-body simulations to place constraints on the number of IMBHs in the Milky Way halo by drawing on the fact that we have not observed any ULXs in our galaxy \\citep{Mapellietal2006}.  These authors find in one simulation that for a halo population of $\\sim 10^5$ IMBHs incorporating a fraction $0.1$\\% of the baryons and distributed in an NFW profile \\citep{NFW}, the predicted number of ULXs per galaxy is on order 1; however, the number of lower-luminosity X-ray sources is overproduced. The constraint found by this method may be applicable to our own galaxy, but it has not as yet been extended to other galaxies, where ULXs are observed.  Since the predictions depend strongly on not only the distribution and number of IMBHs, but also on the properties of the gas in the galactic disk and the efficiency of the black hole accretion (which in turn depends on whether or not an accretion disk is formed), it is difficult to generalize them to other systems.  Other authors have suggested that ULXs may be due to IMBHs accreting from captured stellar companions rather than molecular clouds \\citep{PooleyRappaport2005,Patrunoetal2006, Madhusudhanetal2006}.  In this case, the ULXs would ``turn on'' when residing in dense star clusters.  More detailed observations and simulations are required to answer the ULX question.  Here, we merely point out that the case for IMBH ULXs is an interesting one, easily consistent with current constraints on the dark matter fraction in black holes, and as we show, PBHs may account for or grow into IMBHs by the present era.  Thus PBHs should be considered viable candidates to explain these mysterious sources.  In Figure~\\ref{constraints}, we include as a region of interest the area in parameter space explored in the Mii \\& Totani \\citep{MiiTotani2005} paper (labeled ``PopIII remnants'').  We also mark with an asterisk"
    },
    "0608/astro-ph0608368_arXiv.txt": {
        "abstract": "{} {We investigate the environmental dependence of galaxy population properties in a complete volume-limited sample of 91566 galaxies in the redshift range $0.05 \\le z \\le 0.095$ and with $M_r \\le -20.0$ (that is $M^* + 1.45$), selected from the Sloan  Digital Sky Survey (SDSS) Data Release 4 (DR4). Our aim is to search for systematic variations in the properties of galaxies with the local galaxy density. In particular, we analize how the ($u - r$) color index and the morphological type of galaxies (the latter evaluated through the SDSS {\\em Eclass} and {\\em FracDev} parameters) are related to the environment and to the luminosity of galaxies, in order to find hints that can be related to the presence of a ''void'' galaxy population.} {\"Void\" galaxies are identified through a highly selective criterion, which takes also into account redshift and allows us to exclude from the sample all the galaxies that are not really close to the center of underdense regions. We study the ($u - r$) color distribution for galaxies in different luminosity bins, and we look for correlations of color with the environment, the luminosity, and the morphological type of the galaxies.} {We find that galaxies in underdense regions (voids) have lower luminosity ($M_r>-21$) and are bluer than cluster galaxies. The transition from overdense to underdense environments is smooth, the fraction of late-type galaxies decreases while the fraction of early-type galaxies increases smoothly from underdense to dense environments.} {We do not find any sudden transition in the galaxy properties with density, which, according to a suggestion by Peebles (2001), should mark the transition to a population of \"void\" galaxies in $\\Lambda$CDM models. On the contrary, our results suggest a continuity of galaxy properties, from voids to clusters.} ",
        "introduction": "A self-consistent theory of galaxy formation and evolution within the gravitational collapse paradigm has yet to come, although quite significant advancements have been recently done (e.g. \\citet{2004AIPC..743...33S} for an up-to-date review). Such a theory should explain the observed properties of galaxy populations in very different environments, ranging from very dense clusters to large voids. However, the mere existence of galaxies within voids is a challenge for any gravitational bottom-up scenario of the Large Scale Structure (LSS) formation in the Universe, because in this scenario voids form out of the expansion of underdense ($\\delta < 0$) initial density fluctuations. Theoretical models and numerical simulations \\citep[e.g.][]{1979ApJ...233..395O,2005MNRAS.362..213B} suggest that dark matter flows out of voids, and partly concentrates in some thin boundary shells. The latter, however, have not been detected, although there is clear evidence that the galaxy density profiles inside voids become steeper near their borders \\citep{1996A&A...314....1L,2004ApJ...607..751H,2005ApJ...621..643G}. Observationally, there is evidence that voids are surrounded by superclusters of galaxies, as is vividly shown by recent analyses of the Las Campanas and SDSS surveys \\citep{2003A&A...410..425E, 2003A&A...405..425E, 2005A&A...439...45E}. It is interesing to note that the void models of \\citet{1979ApJ...233..395O} and \\citet{2005AAS...206.1002H} are only concerned with the global dynamics of dark matter (DM): much less work has been devoted to understand the nonlinear collapse of DM halos and the formation of galaxies {\\em within} voids.  \\citet{1986ApJ...303...39D} suggested that voids could be populated by a population of faint dwarf galaxies, where star formation is suppressed by strong photoionization occurred at $z > 6$ \\citep{2005astro.ph..1304H}.\\\\ \\noindent The search for galaxies in voids started with the discovery of the voids themselves (i.e., \\citet{1996A&A...314....1L} for a full review in this field until 1996). \\citet{1996A&AS..116...43P, 1997A&A...325..881P} looked for emission-line galaxies in nearby voids, concluding that most of them are found close to the borders, and far from the deep internal regions. There have also been (failed) attempts at detecting a population of dwarf galaxies more homogeneously distributed within voids \\citep{2003MNRAS.341..981S,2005MNRAS.357..819S, 2004MNRAS.352..478R}, but the lack of an evident faint end of the dwarf galaxies luminosity function does not allow to make firm quantitative predictions about their abundance.\\\\ Since the seminal papers by \\citet{1999AJ....118.2561G,2000AJ....119...32G}, some efforts have been dedicated to detect and to study the properties of bright galaxies in nearby voids. The availability of large samples of galaxies from surveys like the 2dF and the Sloan Digital Sky Survey (SDSS) has made possible more complete analysis of the dependence of galaxy properties on the environment. \\citet{2005MNRAS.356.1155C} have shown that the luminosity function in voids tends to have a steeper slope at high luminosity, but not in the fainter range. The dependence of ($u - r$) color distribution on the environment has also been recently studied by \\citet{2004ApJ...615L.101B}, who found that its shape is bimodal and can be well described by a linear combination of two gaussian distributions for a very wide range of local densities. Surprisingly, they also find a small quantity of {\\em red}, early-type galaxies even in the lowest density environments, thus suggesting that there is not a rigid morphological or color segregation with environment. However, the environmental dependence of galaxy properties is only detected for galaxies in the luminosity range $-20.4  < M_{r} < -19.4$ \\citep{2004AJ....128.2677T}: for galaxies brighter than $M_{r} = -20.4$ there are very little variations of their properties with the environment. \\citet{2001ApJ...557..495P} suggested that galaxies lying within very underdense regions are representative of a distinct population. In this case, we should observe a transition in one or more galaxy properties moving from clusters to voids.\\\\ In order to verify the previous hypothesis, we will investigate the variation of the {\\em color distribution}, {\\em morphology}, and {\\em luminosity} with the local density. This kind of analysis has been made possible by the availability of a large sample of galaxies extracted from the SDSS - Data Release 4 \\citep[SDSS-DR4:][]{2000AJ....120.1579Y,2002AJ....123..485S,2004AJ....128..502A}. Using this huge data sample we confirm some previous results \\citep[e.g.][]{2004ApJ...617...50R}, and we are able to extract statistically significant subsamples for different relative underdensities. We will study the correlation between ($u - r$) colors and the environment for {\\em bright} and {\\em faint} galaxies, also looking for sudden changes which should mark the emergence of a distinct void galaxy populations, as suggested by Peebles (2001). We will use a criterion to estimate the density which is different from the $\\Sigma_{5}$ criterion used by, e.g, \\citet{2004AJ....128.2677T}; on the contrary, it is more similar to the one adopted by \\citet{2004ApJ...617...50R}, because we want to be sure that galaxies in voids are correctly identified and included in the sample.\\\\ The plan of the paper is as follows: after the description of the galaxy sample (sect. 2), in section 3 we will discuss the adopted criteria to construct our subsamples. In section 4 we will present our results concerning the ($u - r$) color distribution, and its dependence on the local density and luminosity. In section 5 we will analize the results of the previous section taking also into account the relationship between the ($u - r$) color distribution and the morphological type of galaxies, then discussing the relationship with the morphology-density relation (Dressler 1980). Finally, in section 6 and in section 7 we will present a general discussion and our conclusions.\\\\ Throughout, we adopt standard present day values of the cosmological parameters to compute comoving distances from redshift: a density parameter $\\Omega_m = 0.3$, a cosmological constant $\\Omega_{\\Lambda} = 0.7$, and a Hubble`s constant value $H_0 = 75~km~s^{-1} Mpc^{-1}$. ",
        "conclusions": "In this paper we analized the environmental dependence of galaxy population properties, looking for those features which could hint at the presence of a \"void\" galaxy population. The photometric properties we used in our analysis can be regarded as complementary respect to those used by other authors \\citep{2004ApJ...617...50R, 2005ApJ...620..618H}. In particular, our density estimator allows us to exploit redshift informations, then  finding galaxies in very underdense regions.\\\\ We can summarize our findings in a few points: \\begin{enumerate} \\item Luminosity is a sensitive parameter to characterize the environmental properties of galaxies. \\item On average, a very underdense environment ($N_{neigh}\\le4$) is populated by galaxies which are bluer than galaxies that are in a very dense enviroments, ($N_{neigh}>79$, see Fig.~\\ref{fig:ur1}). \\item On average, faint galaxies ($M_{r}>-21.0$) are bluer in very underdense environments than in dense environments and  the transition from blue to red, moving from under- to overdense environments, is more pronounced for faint galaxies ($M_{r}>-21.0$, Fig.~\\ref{fig:ur_luminosity}). \\item Although the transition from voids to clusters is more pronounced for faint galaxies, it does not seem to be a discontinuity (Fig.~\\ref{fig:ur_percentile_diff}). \\item Changes in the ($u - r$) color distribution are related to the environment, to the luminosity, and to the morphology of the galaxies. In fact, on average, fainter galaxies ($M_r>-21.0$) are bluer, late-type and in underdense regions (voids)  than brighter galaxies ($M_r<-21.0$), which are redder, early-type and in overdense regions (clusters). \\item Galaxies classified as early-type/late-type and PGs/SFGs, according to their photometric/spectroscopic parameters and to their class of star-forming activity, have a similar ($u - r$) color distribution (Fig.~\\ref{fig:ur_morphology}). Then, the environmental properties of late-type galaxies can be related to the {\\it morphology-density relation} \\citep{1980ApJ...236..351D} not only as a consequence of their morphological type, but also as a consequence of their star-forming activity. \\item We don't find any sudden transition in the properties of \"void\" galaxies, respect to cluster galaxies, as suggested by Peebles (2001). On the contrary, our results show a continuity in the properties of the galaxies, from voids to clusters. \\end{enumerate}"
    },
    "0608/hep-th0608018_arXiv.txt": {
        "abstract": "{We investigate a curved brane-world, inspired by a noncommutative $D_3$-brane, in a type IIB string theory. We obtain, an axially symmetric and a spherically symmetric, (anti) de Sitter black holes in $4D$. The event horizons of these black holes possess a constant curvature and may be seen to be governed by different topologies. The extremal geometries are explored, using the noncommutative scaling in the theory, to reassure the attractor behavior at the black hole event horizon. The emerging two dimensional, semi-classical, black hole is analyzed to provide evidence for the extra dimensions in a curved brane-world. It is argued that the gauge nonlinearity in the theory may be redefined by a potential in a moduli space. As a result, $D=11$ and $D=12$ dimensional geometries may be obtained at the stable extrema of the potential.}     \\preprint  \\begin{document} ",
        "introduction": " ",
        "conclusions": "To summarize, we have obtained generic de Sitter and Anti de Sitter static black hole geometries in $4D$ in a noncommutative brane-world. In particular, a constant moduli case was investigated to yield an axially symmetric dS and AdS black hole geometries with different horizon topologies. In addition, a case with an arbitrary moduli was analyzed for the spherically symmetric black holes in dS and AdS spaces. These black holes were shown to be characterized by their generalized ADM mass $M_{\\rm eff}$, generalized electric or magnetic charge ${\\hat Q}_{\\rm eff}$, in addition to ($k,\\Lambda$). The generalized mass and charge were shown to incorporate all orders in noncommutative parameters $\\Theta$. However, the ${\\cal O}(G_N)$ was argued to be sensistive to the symmetry of a black hole solution in the frame-work. For instance, it was shown that an axially symmetric black hole becomes significant to ${\\cal O}(G^2_N)$, while the spherically symmetry was restored to ${\\cal O}(G_N)$. It was argued that the gauge nonlinearity in the theory enforces back reaction into the metric. The multiple gauge fields in the frame-work were shown to incorporate  a shrink in the radius of event horizon. The extremal dS and AdS with different horizon geometries with a constant curvature were obtained. It was shown that the axially symmetric solution retains its spherical symmetry in the extremal limit. The moduli in the theory were shown to be governed by the electric (magnetic) charges at the event horizon of the dS and AdS black holes, which in turn reconfirms the attractor behavior of event horizon in the frame-work. In addition, the emerging notion of two dimensional de Sitter monopole black hole in the curved brane frame-work is significant. Thus, the $dS_2$ vacuum appears to be a potential candidate to explore some of the quantum gravity aspects in the frame-work.  \\par \\sp It was shown that the event horizon of the extremal AdS coincides with a curvature singularity in the dS space, leading to a single geometry in $2D$. The effective gravitational potential associated with the reduced mass of extremal dS black hole was analyzed to confirm the presence of three extra large dimensions in the regime. The Hagedron transitions in the near horizon geometry of $dS_5$ monopole black hole was analyzed to decouple the $\\Theta$-terms. The new $dS_5$ vacuum with a nontrivial cosmological potential possibly describes our brane-world, $i.e.$ a $D_3$-brane, at its boundary. The potential was argued to be at its local minima on the brane-world  and describes a small positive constant $\\Lambda$. Different space-time dimensions may be guided by different vacua at its extremum in the moduli space.  \\par \\sp It may be interesting to extend the curved brane-world, inspired by a noncommutative $D_5$-brane in type IIB string theory or in a $D=12$ construction for low energy gravity theory. With an appropriate choice of nonlinear EM-field, it may be possible to obtain $dS_4$ and $AdS_4$ black hole geometries using the noncommutative scaling in the theory. Then, the emerging large $4D$ black hole may provide some concrete clue to the tunneling between $dS$ and $AdS$ vacua in the low energy string theory."
    },
    "0608/astro-ph0608333_arXiv.txt": {
        "abstract": "We present critical, long-wavelength observations of $\\eta$ Carinae in the submillimetre using SCUBA on the JCMT at 850 and 450 $\\mu$m to confirm the presence of a large mass of warm dust around the central star.  We fit a two-component blackbody to the IR-submm spectral energy distribution (SED) and estimate between $0.3-0.7\\rm \\, M_{\\odot}$ of dust exists in the nebula depending on the dust absorption properties and the extent of contamination from free-free emission at the SCUBA wavelengths.  These results provide further evidence that $\\eta$ Carinae's circumstellar nebula contains $\\rm > 10\\,M_{\\odot}$ of gas, although this may have been ejected on a longer timescale than previously thought. ",
        "introduction": "\\label{intro}  $\\eta$ Carinae is one of the most luminous objects in our Galaxy ($\\rm 10^6\\, L_{\\odot}$) and is well known for dramatic outbursts in which material is ejected outwards from the star.  Its most famous mass ejection is thought to have occurred in the early 19th Century, where a few solar masses of gas was expelled in a few decades.  This feature is clearly seen in observations as a bipolar nebula, known as the Homunculus (Gaviola 1946).  The origin of the shape of the Homunculus is unknown, but could occur due to ejection of mass in a non-uniform interstellar medium i.e. if the ambient medium was denser in the equatorial plane, the nebula would preferentially expand along the poles.  Surrounding the Homunculus is a larger, less dense nebula, known as the `outer ejecta' (e.g. Weis 2005) which contains numerous filaments and dense condensates not dissimilar to those seen around the young supernova remnant Cassiopeia A (e.g. Smith \\& Morse 2004, Fesen et al 2001).  The central star is hidden from view by this structure but is thought to be a luminous blue variable (LBV) star with mass $\\rm 100\\, M_{\\odot}$ (Pittard 1999; Smith et al. 2003). LBV stars are well known for large luminosities, violent instabilities and periods of large mass loss.  The presence of a binary companion (e.g. Pittard 2003) would provide an explanation for the catastrophic mass ejections and the strong variabilities seen in the radio and X-ray (Cox et al. 1995; Duncan \\& White 2003).  The cycle of periodicity suggests that any exisiting binary would have a highly eccentric orbit, leading to strong wind-wind collisions such as those seen in WR binaries (Marchenko et al. 2002; Pittard 1999).  The evolutionary phase of an LBV star can last more than $10^4$\\,yrs and hence in this phase alone, a single star could lose roughly $\\rm 1\\,M_{\\odot}$ of gas.  Episodic eruptions will increase the mass loss further.  It is therefore not unreasonable to assume that the dust mass loss rate could be $\\rm \\ge 10^{-6}\\,M_{\\odot}\\,yr^{-1}$.  For the lifetime of the average LBV phase, we could expect LBV stars to inject $\\ge 0.01\\rm \\,M_{\\odot}$ of dust into the interstellar medium and could be responsible for the dust emission seen around the young supernova remnants Cas A (Dunne et al. 2003, Wilson \\& Bartla 2005), Kepler (Morgan et al. 2003), SN1998J (Pozzo et al. 2004), the Crab (Green, Tuffs \\& Pospescu 2004), SN2002 (Barlow et al. 2005) and more recently SN2003gd (Sugerman et al. 2006).  The structure of $\\eta$ Car is revealed in the many multiwavelength observations, yet there are issues surrounding the origin of the Homunculus.  IR observations at 12 - 17\\,$\\mu$m revealed a small structure within the Homunculus which appeared to be concentrated around the `waist' of the nebula (Morris et al. 1999; hereafter M99). They found a small bright core a few arcsecs across which they interpreted as a disk/torus of dust with two components at 200\\,K and 110\\,K with $\\rm \\sim 0.15\\, M_{\\odot}$ of dust.  They suggested (with a normal gas-to-dust ratio of 100) that the gas within this small torus could therefore be $\\rm \\sim 15\\, M_{\\odot}$.  This argument is flawed since estimating the 110\\,K dust mass relied on observations at the longer infrared wavelengths where resolution is limited to greater than 1 arcminute, so the precise location of the 110K dust component cannot be determined from the mid-IR images alone.  Davidson \\& Smith (2000) also argued against the IR emission in a small torus, proposing instead that this was too small to account for the large luminosities radiated at 100\\,$\\mu$m.  The visual image provides further evidence for this argument since there is no evidence of an opaque dusty disk on such a small scale.  Hony et al. (2001) then used high resolution images at 8 - 20\\,$\\mu$m to suggest that the observed IR dust is not warm dust at 110\\,K, but actually hot dust at 200\\,K with a strange overlapping ring structure in the middle of the nebula. They argue that these rings are similar to those seen in SN1987A and are at different polar axes to each other.  If this scenario is correct, the total dust mass in $\\eta$ Car decreases to $\\sim \\rm 10^{-4}\\,M_{\\odot}$ due to the strong dependence of dust mass with temperature; the gas mass ejected in the Homonculus would then be $< \\rm 10^{-2}\\,M_{\\odot}$ and the evolution of $\\eta$ Car would not be as violent as currently thought.  This claim has been disputed in Smith et al. (2000; 2003) who propose that the strange smooth loop structure seen in the Hony et al. observations are simply due to a disrupted torus.  They point to compact knots and hotspots in the IR with filaments and arcs which could be due to a toroidal distribution of dust which has been disrupted by subsequent mass loss. They also use temperature maps from ISO data to show that cooler `hidden' dust may exist in the outer polar lobes of the Homunculus and not in a compact torus.  M99, Davidson \\& Smith (2000) \\& Smith et al. (2002) have all argued that there is missing mass which has never been accounted for by the observations, and the cooler dust mass which has been found has not been resolved, so its precise location is unknown.  If there is missing ejected mass, this will have important implications for the energy budget and estimated parameters of LBV stars as well as its future evolution. To determine if colder dust could exist further out from the centre of the star, wavelength measurements are needed beyond the mid-IR range (Davidson \\& Smith 2000) which will dominate the dust mass.  We therefore use archived SCUBA observations of $\\eta$ Car to investigate the distribution of the dust and the ejected mass in one of the most studied objects in the nearby Universe. ",
        "conclusions": ""
    },
    "0608/astro-ph0608619_arXiv.txt": {
        "abstract": "We report measurements of time delays of up to 8 minutes in the centimeter wavelength variability patterns of the intra-hour scintillating quasar PKS~1257$-$326 as observed between the VLA and the ATCA on three separate epochs. These time delays confirm interstellar scintillation as the mechanism responsible for the rapid variability, at the same time effectively ruling out the coexistence of intrinsic intra-hour variability in this source. The time delays are combined with measurements of the annual variation in variability timescale exhibited by this source to determine  the characteristic length scale and anisotropy of the quasar's intensity scintillation pattern, as well as attempting to fit for the bulk velocity of the scattering plasma responsible for the  scintillation. We find evidence for anisotropic scattering and highly elongated scintillation patterns at both 4.9 and 8.5 GHz, with an axial ratio $>10:1$, extended in a northwest direction on the sky. The characteristic scale of the scintillation pattern along its minor axis is well determined, but the high anisotropy leads to degenerate solutions for the scintillation velocity. The decorrelation of the pattern over the baseline gives an estimate of the major axis length scale of the scintillation pattern. We derive an upper limit on the distance to the scattering plasma of no more than 10\\,pc. ",
        "introduction": "Interstellar scintillation (ISS) of extragalactic radio sources at centimeter wavelengths can be used as a probe of both their microarcsecond-scale structure and the structure of the turbulent Galactic interstellar medium (ISM) \\cite[e.g.][]{mj2002}.   However, before scintillation can be utilized as an effective tool to probe both the radio source structure and the ISM, it is first necessary to establish the overall dimensions of the scintillation pattern,  and the distance to and velocity of the medium responsible for the scattering of the source.  For a few sources the variations are sufficiently rapid that it is possible to measure a time delay in the arrival times of the pattern of flux density variations between widely separated telescopes as the Earth drifts through the scintillation pattern. In principle such a time delay provides a direct determination of the instantaneous scintillation velocity and helps ascertain the overall scale of the scintillation pattern. When the scintillation pattern is highly anisotropic, however, the anisotropy introduces a bias into the observed delay, as noted by \\citet{ps55} in an analysis of ionospheric scintillations.  For many sources affected by ISS the Earth's orbital motion causes a substantial change in the relative velocity between the Earth and the scattering plasma which in turn changes the observed pattern time delay.  In this case, measurements of the time delay at various epochs during the year can be used to constrain the anisotropy and position angle of the scintillation pattern.  PKS~1257$-$326 is a flat spectrum, radio-loud quasar at $z=1.256$ which exhibits intra-hour flux density variability at frequencies of several GHz due to ISS \\citep[][ hereafter Paper I]{big2003}.  The source's rapid variability makes it possible to measure a precise time delay between two widely separated telescopes, and hence to determine its structure through measurement of its scintillation characteristics.  The recent MASIV Survey of 700 flat-spectrum radio sources \\citep{lov2003} showed that rapid variability of the type exhibited by PKS~1257$-$326 is extremely rare.  Only two other quasars are known that vary  on timescales short enough to allow time delay measurements to be made.  These are PKS~0405$-$385 \\citep{ked97} and J1819$+$3845 \\citep{dtdb2000}.  Simultaneous observations of  PKS~0405$-$385 showed a time delay of around 2 minutes between the Australia Telescope Compact Array (ATCA) and the Very Large Array (VLA) \\citep{iau205:jau2000}.  Observations of J1819$+$3845 at the Westerbork Synthesis Radio Telescope and the VLA  also showed a clear delay between the pattern arrival times at each telescope, which changed over the course of the observation due to the rotation of the Earth and hence of the projected baseline vector \\citep{dtdb2002,dtdb2003}. Similar time delay techniques have been used to study ISS of pulsars \\citep[e.g.][]{lr70}, and interplanetary scintillation of AGN \\citep[e.g.][]{ck78}.  In the regime of weak scattering, relevant to the scintillation of PKS~1257$-$326 above $\\sim 5\\,$GHz, the length scale of the scintillation pattern is related to the Fresnel scale, $r_{\\rm F} = \\sqrt{cL/(2\\pi\\nu)}$, where $\\nu$ is the observing frequency and $L$ is the distance to the scattering medium which is assumed to be confined to a plane.  For the case of isotropic scattering, a source of angular diameter $\\theta_{\\rm S}$ larger than the angular Fresnel scale, $\\theta_{\\rm F}=r_{\\rm F}/L$, increases the scale length of the scintillation pattern by a factor $\\sim \\theta_{\\rm S}/\\theta_{\\rm F}$ \\citep[e.g.][]{nar92}.  Thus if the angular size of PKS~1257$-$326 is several times larger than $\\theta_{\\rm F}$, then the scale of the scintillation pattern should be determined chiefly by source structure. Assuming the same scattering material is responsible for the scintillation at different frequencies, then the time delay is independent of observing frequency {\\em unless} source structure changes with frequency to produce anisotropy in the scintillation pattern which is significantly different at different frequencies. If there is significant anisotropy in the scattering medium, however, then the medium may have a more dominant influence on the structure of the scintillation pattern.  We present here the results of simultaneous ATCA and VLA observations of PKS~1257$-$326, carried out during 2002 and early 2003.  These are combined with measurements of the characteristic timescale of variability from ATCA data collected over a three-year period, in order to determine the scintillation velocity and dimensions of the pattern. ",
        "conclusions": "\\subsection{Anisotropy in the scintillation pattern}  Figure~\\ref{fig-ell} shows the scintillation velocity over the course of the year for PKS~1257$-$326 along with a contour of the fitted scintillation pattern projected on the plane of the sky for anisotropic solution (c) in Table~\\ref{tab-fit}. Note that in Paper I the geometry (Figure 4 in Paper I) was displayed as if looking down on the Earth whereas in the present paper we show the view looking up at the sky, corresponding to a reflection of the east-west axis. It is interesting to note that the long axis of the scintillation pattern is aligned in a north-west direction.  This is close to the direction of the main extended arcsecond-scale jet in the source, and also with the milliarcsecond-scale jet components observed in VLBI data (Bignall et al., in preparation).  \\citet{ric2002} argued that for the case of PKS~0405$-$385 there is strong evidence that the scintillations result from anisotropic scattering in a thin screen. Likewise for PKS~1257$-$326 we observe the same large negative ``overshoot'' in the light-curve autocorrelations as investigated by \\citet{ric2002} and shown to be due to anisotropic scattering. However, given the coincidence of position angle, it may be the case for PKS~1257$-$326 that the anisotropy in the scintillation pattern is also partly related to elongation of the source itself. It is possible that if the source were elongated in a different direction, then such large and rapid scintillation might not be observed through the anisotropic scattering screen. Further analysis is required before conclusions can be drawn as to whether the anisotropy is entirely due to the scattering screen, source structure, or a combination of both source and screen.  \\begin{figure} \\epsscale{1}\\plotone{f6.eps} \\caption{Scintillation velocity projected onto the plane of the sky, showing the annual variation of the Earth's velocity with respect to the scattering screen for anisotropic solution (c) in Table~\\ref{tab-fit}. Dotted lines show the velocity on the days of the time delay observations. The filled ellipse represents a contour of the best-fit scintillation pattern at 4.9\\,GHz. The 1-$\\sigma$ contour of the best fit solution for $v_{{\\rm ISS},\\alpha}$ versus $v_{{\\rm ISS},\\delta}$ is plotted around the best fit velocity. The component of velocity parallel to the long axis of the scintillation pattern is poorly constrained. If the scattering screen were moving with the local standard of rest, the center of the ellipse would be at position marked in the top right-hand corner.} \\label{fig-ell} \\end{figure}  The origin of the anisotropy may be resolved through a detailed analysis of the polarized flux density scintillations which is ongoing. The ATCA observations provide the full set of Stokes parameters, $I$, $Q$, $U$ and $V$. While no significant circular polarization has been detected for PKS 1257$-$326 ($3\\sigma$ upper limit $V = 0.2$\\,mJy, or 0.1\\% of $I$), we detect significant flux density in the linear polarization parameters $Q$ and $U$.  If anisotropy is a property of the scattering medium, then the same degree of anisotropy should be present in all Stokes parameters, while if a property of the source, then the linear polarization is not likely to have the same anisotropy as Stokes $I$.  Cross-correlations between the  various Stokes parameter pairs can also be used to reveal the microarcsecond-scale polarized structure of the source \\citep{mj2002,ric2002}.  More information on the properties of both source and scatterer can also be obtained through a detailed analysis of light curve structure functions and power spectra -- here we have used only the $1/e$ decorrelation timescale to characterise the scintillations.  \\subsection{Implications for the source}  For a scattering screen at $L=10$\\,pc, the angular scale of the scintillation pattern implies that the angular size of the scintillating component of the source cannot be much larger than $\\sim 300$ by 30 $\\mu$as at 4.9\\,GHz (a factor $\\sim 1.2$ smaller at 8.5\\,GHz).  Based on the flux density in the unresolved VLBI core, the observed modulation index and also the asymmetry coefficient analysis of \\citet{sst2005}, we estimate that the flux density of the scintillating component is close to 100\\,mJy at 4.9\\,GHz. Then the lower limit for the implied brightness temperature of this component is $T_b \\approx 1.4\\times 10^{12}$~K. If the anisotropy is caused purely by the scattering medium and the source is intrinsically circularly symmetric with angular diameter 30\\,$\\mu$as, then the implied brightness temperature becomes a factor of 10 higher. On the other hand, if the scattering screen is in fact closer than 10\\,pc, this implies a larger angular size and hence lower $T_b$ for the source.  $T_b \\sim 10^{13}$~K is still within the range inferred from VLBI observations of other flat-spectrum quasars \\citep[e.g.][]{hor2004}.  A Doppler factor of $\\sim 10$--20 is implied to reduce the intrinsic brightness temperature below the inverse Compton limit, however the relativistic beaming evidently does not correspond to an expansion of the emitting region. While there has been some slow evolution of the source flux density over three years of monitoring, we observed a repeated annual cycle, and rapid scintillation has been seen every time PKS~1257$-$326 has been observed since 1995, indicating that the compact component is long-lived and relatively stable over at least a decade.  \\subsection{Implications for the scattering plasma}  The MASIV Survey showed that such extreme intra-hour variability (IHV) as observed in PKS~1257$-$326, PKS~0405$-$385 and J1819+3845 is exhibited by $\\ll 1$\\% of compact flat-spectrum radio sources. Combined with the interpretation of the rapid variability as due to scattering in the very local ISM, this suggests that the covering fraction of such nearby scattering plasma is very small. Alternatively, it is possible that a large fraction of sources may be sufficiently scatter-broadened by the more distant ISM to quench the weak scintillations in more nearby scattering plasma. Further investigation of these possibilities and of the local ISM structure inferred from other observations is needed to fully understand the IHV phenomenon. For a transition frequency between strong and weak scattering of $\\sim 5$\\,GHz, as indicated by the observed peak of modulation close to this frequency, and a screen distance of 10\\,pc, the scattering measure ($SM$) in conventional units of kpc m$^{-20/3}$ is found to be $\\log SM \\sim -2.8$ using equation 17 from \\citet{cl2002}.  All three fast scintillators show evidence of highly anisotropic scintillation patterns \\citep{ric2002,dtdb2003}. Interestingly, highly anisotropic scattering has also been inferred in the interpretation of parabolic arcs observed in the secondary spectra of pulsars \\citep[e.g.][]{wal2004}. \\citet{rc2004} discuss the evidence for highly anisotropic plasma turbulence concentrated in relatively thin layers, suggesting filamentary density structures aligned by the local magnetic field. If the scattering occurred in a very extended medium, the magnetic field, and hence anisotropy in the electron density fluctuations, would tend to have different orientations through the scattering medium and thus the scintillation pattern would be expected to be more isotropic.  \\subsection{Conclusions}  We have presented a combined analysis of two-station pattern arrival time delays and annual cycles in the scintillation timescale for the quasar PKS~1257$-$326. This is the first time such a combined fitting method has been applied to data on a scintillating quasar. We conclude that: \\begin{itemize}  \\item the large, rapid flux density variations observed in PKS~1257$-$326 are entirely due to interstellar scintillation.  \\item the scattering occurs in a confined region within 10\\,pc of the solar system.  \\end{itemize}  The scintillation  velocity and geometry of the scintillation pattern for PKS 1257$-$326 could not be uniquely determined using the annual cycle data alone (Paper I), although none of the conclusions presented in Paper I are  significantly altered by the new results. Analysis of the combined annual cycle and time delay data shows that there is a high degree of anisotropy in the scintillation pattern. The characteristic linear scale of the pattern along its short axis is well determined, as is the position angle along which the scintillation pattern is elongated.  The large anisotropy results in a degeneracy in fitting for the scattering screen velocity, with the best fit velocities  lying along a line in the $(v_{\\alpha},v_{\\delta})$ plane. The scintillation scale on the long axis is constrained using measurements of the spatial decorrelation of the pattern between the VLA and ATCA.  An axial ratio of at least 12:1 is inferred.  The implied source brightness temperature is high, up to $10^{13}$\\,K, but not incompatible with brightness temperatures derived from VLBI observations of other compact quasars. Although a compact source is a requirement for ISS to be observed, it seems likely that it is the unusual scattering properties of the local ISM on the line-of-sight to PKS~1257$-$326, rather than an intrinsic property of the source itself, which give rise to the unusually rapid scintillation observed in this case.  Pattern arrival time delays between widely separated telescopes, combined with measurements of the scintillation timescale, provide a powerful method of determining scintillation parameters for quasars which scintillate rapidly. In order to measure such time delays however, a source must vary by a detectable amount (typically of order 1 mJy or more) over a period of order 1 minute.  Such variability is observed in very few sources. Another problem is that the scintillation behaviour for some sources is episodic, making it more difficult to coordinate observations involving multiple telescopes.  With more sensitive telescopes coming on line in the future, one would expect to find more low flux density rapid scintillators, thus probing scattering on many more lines of sight through the local ISM."
    },
    "0608/astro-ph0608105_arXiv.txt": {
        "abstract": "{With the Plateau de Bure interferometer, we have measured the C, N, O and S isotopic abundance ratios in the arm of a spiral galaxy with a redshift of 0.89. The galaxy is seen face-on according to HST images. Its bulge intercepts the line of sight to the radio-loud quasar \\object{PKS~1830-211}, giving rise at mm wavelengths to two Einstein images located each behind a spiral arm. The arms appear in absorption in the lines of several molecules, giving the opportunity to study the chemical composition of a galaxy only a few Gyr old. The isotopic ratios in this spiral galaxy differ markedly from those observed in the Milky Way. The $^{17}$O/$^{18}$O and $^{14}$N/$^{15}$N ratios are low, as one would expect from an object too young to let low mass stars play a major role in the regeneration of the gas. ",
        "introduction": "The C, N, O, S isotopic abundances are enlightening relics of the interstellar medium (ISM) past history. Not only are these elements the most abundant ones after H and He, but they are those whose synthesis in stars has been the best investigated. The CNO isotopes are tentatively classified into primary elements ($^{12}$C, $^{16}$O,...), formed directly by He burning, and secondary ones ($^{13}$C, $^{14}$N, $^{17}$O,...) formed at a later stage from primary elements, e.g. through the CNO cycles. In simple evolution models, the fraction of secondary elements increases with time. Actually, chemical evolution is not so straightforward and we know that isotopes like $^{14}$N and perhaps $^{13}$C are already in part produced in the first generation of stars.  Whereas isotopic abundances can be precisely measured in the Sun and on the Earth, their values in the Galactic ISM and, mostly, in the extragalactic ISM are hard to derive. In a pioneering work, Penzias and co-workers (Penzias et al. \\cite{pen80}, \\cite{pen81}) used the millimeter (mm) emission lines of molecular isotopologues to infer the isotopic ratios in a dozen molecular sources. They found significant differences between the Solar System, the local ISM, the Galactic Center and circumstellar envelopes that they interpreted in terms of degree of nuclear processing, the Solar System representing the state of the local ISM some 4.5 Gyr ago. This picture of the Galactic ISM remains essentially valid today, although it has been recognized that the molecular ratios involving D and $^{13}$C are affected by isotopic fractionation or selective photodissociation and may not reflect the elemental isotopic abundance ratios (Wilson \\& Matteucci \\cite{wil92}).  So far, only a few isotopic abundance ratios have been reported in external galaxies. Those are mostly the C, N, O and S isotopic ratios derived in the nuclei of four nearby starburst galaxies, M~82, IC~342, NGC~253 and NGC~4945, from the mm emission of CO, CS and HCN (Henkel \\& Mauersberger \\cite{hen93}, Henkel et al. \\cite{hen98}, Chin et al. \\cite{chi99b}, Wang et al. \\cite{wan04}, Mart\\'in et al. \\cite{mar05}). These ratios are very uncertain due to beam dilution, because the lines of the rare isotopologues are weak, and because the line opacity of the main isotopologues is unknown. Moreover, the studies deal with nuclear regions that may not reflect the disk ISM abundances, as is the case in the Milky Way.  Another powerful way of studying the composition of the ISM consists in observing the molecular lines in absorption. The method needs only a strong background source and is insensitive to distance and to primary beam dilution. It yields directly the line opacities, provided the background source is fully covered by the intercepting clouds (e.g. Lucas \\& Liszt \\cite{luc98}). As far as molecular isotopic ratios are concerned, the method is limited to the line of sight to the brightest radio galaxies and quasars intercepted by deep absorption line systems. Few such absorption systems have been identified so far at millimeter wavelengths. They arise in the Milky Way (Lucas \\& Liszt \\cite{luc96}, \\cite{luc98}) and in a few distant galaxies with redshifts in the range z = 0.2 -- 0.9 (Wiklind \\& Combes \\cite{wik94}, \\cite{wik95}, \\cite{wik96}, Combes \\& Wiklind \\cite{com98}, Combes \\cite{com99}).  In this article, we present high sensitivity absorption spectra arising from the most remote presently known system, located at z = 0.89. The spectra cover the lines of a dozen of molecular species, including the rare isotopomers HC$^{17}$O$^+$, H$^{15}$NC and C$^{34}$S. From these spectra, we derive the C, N, O and S isotopic ratios in the spiral arm of a young, seemingly normal galaxy.\\footnote{We adopt the usual convention to simply denote by C, N, O and S the abundant isotopes $^{12}$C, $^{14}$N, $^{16}$O and $^{32}$S.} ",
        "conclusions": "The presence of an intervening nearly face-on spiral galaxy at z = 0.89 in the line of sight of the radio loud quasar PKS~1830-211 offers an unique chance to study the chemical and isotopic compositions of the interstellar medium in the spiral arm of a galaxy 2 -- 3 times younger than the Milky Way. The use of the PdB interferometer has enabled us to combine very flat baselines and high sensitivity, making it possible to measure reliably the C, N, O and S isotopic abundance ratios in HCO$^+$, HCN, HNC, CS, and H$_2$S. The ratios derived from different species or velocity components are fully consistent, which supports the view that they are not affected by residual opacity effects and/or fractionation effects and that they reflect the elemental isotopic ratios.  A comparison of the isotopic ratios in the z = 0.89 galaxy with those observed in the local ISM, the Galactic Center, the circumstellar envelope IRC+10216 and the Solar System shows major differences with these sources. Mainly, the $^{17}$O/$^{18}$O, and $^{14}$N/$^{15}$N ratios, which are sensitive probes of the degree of processing of the gas by low mass stars and which are expected to increase with time are found to be smaller than in the ISM or the Solar System and much smaller than in the Galactic Center and in the fully processed low-mass star envelope IRC+10216. Most remarkably, the observed $^{16}$O/$^{18}$O ratio is a factor of 10 lower than in the ISM and the Solar System and the $^{32}$S/$^{34}$S ratio is twice lower than in the Galactic sources. These latter ratios result mainly from nucleosynthesis in high mass stars, it is likely that the stellar Initial Mass Function in the z = 0.89 galaxy is different from that of the Milky Way at the same age. The low $^{17}$O/$^{18}$O, and $^{14}$N/$^{15}$N ratios are consistent with the young age of the galaxy ($<$ 6 Gyr) which is smaller than the lifetime of stars with masses $<$ 1.5 M$_\\odot$ and severly limits the importance of low mass stars in the regeneration of the ISM.  The isotopic ratios observed towards PKS~1830-211, although more extreme, are not unsimilar to those derived in the starburst nuclei of nearby galaxies, such as NGC~253, M82 and NGC~4945 (see e.g. Mart\\'in et al. \\cite{mar06}). The later ratios, which result from the observation of weak and broad emission lines, are however far more uncertain.  The extension of this study to other remote galaxies is hampered by the small number of millimeter radio sources whose line-of-sight intercepts the plane of a remote galaxy. PKS~1830-211 is by far the brightest such source. The second best, at present, is the gravitational lens B0218+357 and several molecular lines have been observed in absorption at z = 0.68 toward this object (Wiklind \\& Combes \\cite{wik95}, Menten \\& Reid \\cite{men96}, Combes \\& Wiklind \\cite{com97}). Even in this object, the detection of less abundant species is difficult (Henkel et al. \\cite{hen05}). Fortunately, the flux of this source has been growing lately at millimeter wavelengths, giving hope that the rare carbon, oxygen and sulfur isotopes can be studied in the near future."
    },
    "0608/astro-ph0608275_arXiv.txt": {
        "abstract": "We perform numerical simulations of  black-hole binaries to study the exchange of spin and orbital angular momentum during the last, highly nonlinear, stages of the coalescence process. To calculate the transfer of angular momentum from orbital to spin, we start with two quasi-circular configurations, one with initially non-spinning black holes, the other with corotating black holes. In both cases the binaries complete almost two orbits before merging. We find that, during these last orbits, the specific spin ($a/m$) of each horizon increases by only $0.012$ for the initially non-spinning configuration, and by only $0.006$ for the initially corotating configuration. By contrast, the corotation value for the specific spin should increase from $0.1$ at the initial proper separation of $10M$ to $0.33$ when the proper separation is $5M$. Thus the spin-orbit coupling is far too weak to tidally lock the binary to a corotating state during the late-inspiral phase. We also study the converse transfer from spin into orbital motion. In this case, we start the simulations with parallel, highly-spinning non-boosted black holes. As the collision proceeds, the system acquires a non-head-on orbital motion, due to spin-orbit coupling, that leads to the radiation of angular momentum. We are able to accurately measure the energy and angular momentum losses and model their dependence on the initial spins. ",
        "introduction": "\\label{Sec:Intro}   With the advent of new methods to evolve black-hole binaries for several orbits~\\cite{Pretorius:2005gq,Campanelli:2005dd,Baker:2005vv,Campanelli:2006gf,Baker:2006yw,Pretorius:2006tp,Scheel:2006gg} we can now pose questions of physical interest about the details of their interactions and the generation of gravitational radiation, and find answers that allow us to better understand the astrophysical properties of black-hole binaries.  The coalescence of unequal-mass black-hole binaries with randomly oriented individual spins is expected to be one of the strongest astrophysical sources of gravitational radiation. The first numerical simulations of non-spinning unequal-mass black-hole binaries, along with the calculation of the merger kicks, were reported in Refs.~\\cite{Campanelli:2004zw, Herrmann:2006ks, Baker:2006vn}.  In a preceding paper~\\cite{Campanelli:2006uy} we performed the first fully-nonlinear simulations of highly spinning black-hole binaries and showed that the spin-orbit coupling leads to significant changes in the orbital dynamics, as well as gravitation radiation emitted, in the final stages of a black-hole binary. In this paper we study in more detail the mechanisms that transfer orbital angular momentum to spin and {\\it vice versa}.  It is of astrophysical interest to know if tidal effects could lock the spin of black-hole binaries to a corotating state, i.e.\\ a state where the spins of the individual horizons are aligned with orbital angular momentum, and the spin magnitudes are such that the  horizon frequency $\\Omega_H$ ($\\Omega_H$ is the angular speed of locally-non-rotating observers as they pass through the horizon, as seen by stationary observers at infinity) is nearly equal to the orbital frequency $\\Omega_O$. For a Kerr hole with mass $m$ and specific spin $a/m$ the horizon frequency takes the value \\begin{equation} \\label{eg:OmegaH} \\Omega_H=\\frac{1}{2m}\\left(\\frac{a/m}{1+\\sqrt{1-(a/m)^2}}\\right). \\end{equation} See Ref.~\\cite{Caudill:2006hw} for a first post-Newtonian (1PN) correction to the corotation condition $\\Omega_O = \\Omega_H$. Note that in this paper we consider both the spin (or more accurately, the spin angular momentum) $S$ of a black hole, which has units of $M^2$, and the specific-spin $a/m$ ($m$ is the horizon mass), which is dimensionless. The two are related by $S=a m$. We will also make use of the following relationships between the spin, specific spin, horizon frequency, and irreducible mass ($m_{\\mbox{\\tiny IR}}$): \\begin{eqnarray} \\label{eq:SOmegaH} S &=& \\frac{4 m_{\\mbox{\\tiny IR}}^3 \\Omega_H}{\\sqrt{1 - 4 (m_{\\mbox{\\tiny IR}} \\Omega_H)^2}}, \\\\ \\label{eq:aOmegaH} a/m &=& 4 m_{\\mbox{\\tiny IR}} \\Omega_H \\sqrt{1 - 4 (m_{\\mbox{\\tiny IR}} \\Omega_H)^2}, \\end{eqnarray} where $16 \\pi m_{\\tiny \\mbox{IR}}^2$ is the horizon surface area.     There are two competing scales in this problem: one is the time for tidal effects to act to spin up (or down) the black holes to the orbital frequency and the other is the rate of change of the orbital frequency due to gravitational radiation.   In the early stages of the orbital motion of black-hole binaries, tidal effects are small and can be described by Newtonian theory; these effects decay strongly with the binary separation $r$ (as $\\mu/r^3$, with $\\mu$ the reduced mass of the system). Tidal effects will be stronger as the holes get closer, but this will also increase the orbital frequency (roughly like $\\Omega_O\\sim\\sqrt{\\mu/r^3}$). In order to be correctly quantified, these effects need to be evaluated dynamically with a fully non-linear numerical simulation of the final orbital stage of black-hole binaries. Despite this, several computations of quasi-circular sequences of initial data for black-hole binaries have assumed corotating spins from larger separations down to the estimated location of the innermost stable circular orbit (ISCO) ~\\cite{Blanchet:2001id,Grandclement02,Caudill:2006hw}. Corotation has a notable effect on the location of the ISCO in all these calculations; fortunately, recent numerical evolutions of equal-mass black-hole binaries ~\\cite{Campanelli:2006gf,Baker:2006yw,Pretorius:2006tp} do not find any particular physical effect associated with the location of the ISCO.   Bildsten and Cutler \\cite{Bildsten92} (see also \\cite{Kochanek:1992wk}) studied neutron star (NS) and black hole (BH) systems.  They conclude that, for the NS-NS binary the viscosity is too low to tidally synchronize the stars as they inspiral, and that there is simply not enough time during the gravitational radiation driven inspiral of NS-BH binaries to tidally lock the spin of the neutron star to the orbital period. Additionally, when the NS-BH binaries are close enough for the tidal effects to be strong, the neutron stars are disrupted in a few orbits. Since these arguments do not strictly apply to black-hole binaries, it is interesting to explore these tidal interactions dynamically with fully non-linear numerical techniques. This spin-up of NS-NS binaries has also been analyzed in Ref.~\\cite{Marronetti:2003gk}.  We will study in Sec.~\\ref{Sec:spinup} the tidal spin-up of two equal-mass, initially non-spinning, black holes starting from a quasi-circular orbit (as determined by 3PN trajectories) with nearly two orbits left before merger, as well as a similar configuration with initially corotating black holes. These equal-mass configurations involving the very last stages of the inspiral should nearly maximize the tidal effects.   While the transfer of orbital angular momentum to the spin of the members of a binary system is of astrophysical relevance, we can also study the conversion of spin into orbital angular momentum and verify the importance of its dynamical effects. To this end, we evolve binaries consisting of non-boosted, equal-mass black holes, with the black holes aligned along the $y$-axis and the individual spins aligned along the $+z$-axis, and measure the radiated energy and angular momentum. This simple set-up clearly displays the transfer of spin to orbital motion. For non-spinning holes this set-up produces purely head-on collisions, leading to the radiation of energy but {\\it not} angular momentum.  For spinning holes, the spin-orbit coupling causes the holes to acquire non-head-on orbital motion, leading to the radiation of angular momentum. Thus, an accurate measure of the angular momentum radiated would allow us to estimate the correct magnitude of this effect.  The leading post-Newtonian spin-orbit and spin-spin corrections to the acceleration of two equal-mass black holes (each of mass $m$), with spins $\\vec{S}_1$ and $\\vec{S}_2$ perpendicular to the motion, in a nearly head-on collision are: \\begin{eqnarray} \\vec{a}_{SO}&=&-\\frac{16}{r^3}(\\vec{v}_r\\times(\\vec{S}_1+\\vec{S}_2)),\\\\ \\vec{a}_{SS}&=&-\\frac{6}{mr^4}(\\vec{S}_1\\cdot\\vec{S}_2)\\,\\hat{r}, \\end{eqnarray} where $\\vec r$ is the binary separation, $\\vec v_r$ is the relative velocity, and $SO$ and $SS$ denote spin-orbit and spin-spin coupling terms respectively. While the spin-spin acceleration adds an attractive term along the radial direction (for parallel spins; repulsive for anti-parallel), the spin-orbit acceleration is perpendicular to both the mostly-radial (along $\\hat{r}$) head-on velocity and the spin, and hence leads to a non-vanishing orbital component of the binary motion.   In Sec.~\\ref{Sec:spin-orbit} we will study this nonlinear coupling using fully non-linear numerical evolutions and show that our method is able to extract meaningful results for this numerically demanding relatively small effect. ",
        "conclusions": "\\label{Sec:discussion}  We have studied the spin-up of initially non-spinning and initially corotating black-hole binaries during the last stages of the inspiral orbital motion and found that it is much smaller than that needed to achieve tidal locking. This means that the time scale of tidal effects to spin up the black holes is much longer than the time scale for gravitational radiation to drive the merger (and hence increase the orbital frequency). This also implies that in an astrophysical scenario, black holes with low spin merge with essentially the same spin magnitude they had when the binary formed (note that Eq.~(\\ref{Sdot}) implies that $|\\dot{\\vec{S}}|=0$). We also note that it is this short timescale for gravitational radiation during the last few orbits that is responsible for the absence of an ISCO (innermost stable circular orbit) when the two holes have comparable masses. The concept of the ISCO arises naturally in the extreme mass ratio limit, where there is a cutoff between stable and unstable orbits. In this limit, the instability of orbits inside the ISCO is not driven by gravitational radiation, but by non-radiative terms in the particle's equation of motion. Thus the particle very slowly inspirals due to the emission of weak gravitational radiation, and then plunges when it reaches the ISCO. On the other hand, in the comparable mass regime, there is no such cutoff because gravitational radiation dominates the dynamics during the entire late-time inspiral and plunge phases.  In Fig.~\\ref{fig:tracks} we show the tracks of the horizon centroids in the S0 and SC configurations. Although coordinate dependent, the tracks help illustrate the lack of an ISCO in the comparable mass regime. Rather than showing an abrupt change near the last orbit, the tracks show a smooth increase in plunge velocity over the entire evolution. Note that the SC track is a tighter spiral than the S0 track, which is consistent with the results found previously by the authors~\\cite{Campanelli:2006uy}. Additional horizon tracks for multi-orbit non-spinning binaries can be found in Refs.~\\cite{Baker:2006yw, Campanelli:2006gf}, while tracks for highly-spinning binaries can be found in Ref.~\\cite{Campanelli:2006uy}.  \\begin{figure} \\begin{center} \\includegraphics[width=3.3in]{S0SC_track.eps} \\caption{The horizon tracks for the S0 and SC configurations with tick marks every $10M$ of evolution (we display only one of the SC tracks for clarity). Note that the S0 and SC tracks agree well for the first $100M$, and that after $100M$ the SC tracks show a tighter spiral and orbit longer before merging. The S0 configuration completes slightly less than 1.75 orbits, while the SC configuration completes slightly more.} \\label{fig:tracks} \\end{center} \\end{figure}  After the merger, the final remnant black hole formed from the S0 configuration has a specific spin of $a/M_{H}=0.688\\pm0.001$ (where $M_{H}$ is the remnant horizon mass) while remnant formed from the SC configuration has a specific spin of $a/M_{H}=0.717\\pm0.001$. Using these values in conjunction those found in Ref.~\\cite{Campanelli:2006uy} (see Table~\\ref{tab:spindat}), we find (from a least-squares fit) that the specific spin of the final remnant horizon $(a/M_{H})|_R$ varies with the initial specific spins of the two (equal-mass, equal-spin) black holes $(a/m)|_{I}$ according to: \\begin{eqnarray} (a/M_{H})|_{R} &=& 0.6879 + 0.2952 \\left((a/m)|_{I}\\right) \\nonumber \\\\ &-&0.0374 \\left((a/m)|_{I}\\right)^2, \\end{eqnarray} with an error of $\\pm0.001$ in the range $(-0.757 < (a/m)|_{I} < 0.757)$. From this fit we extrapolate that the remnant horizon will have a specific spin no smaller than $0.355$ (for the case of initially-maximal, anti-aligned spins) and no larger than $a/M_{H} = 0.946$ (for the case of initially-maximal, aligned spins), which supports our assertion in~\\cite{Campanelli:2006uy} that the cosmic censorship conjecture is not violated by the merger of two black holes. A least-squares fit of the energy radiated versus the initial individual spins yields \\begin{eqnarray} \\frac{E_{rad}}{M} &=& 0.0348 + 0.0297 \\left((a/m)|_{I}\\right) +\\nonumber \\\\ && 0.0170 \\left((a/m)|_{I}\\right)^2, \\end{eqnarray} with an error of $\\pm0.001$ in the range $(-0.757 < (a/m)|_{I} < 0.757)$. Based on this fit we extrapolate that a quasi-circular binary will radiated no more than $8.1\\%$ of its mass during the final few orbits and merger.  \\begin{table} \\caption{The measured remnant horizon specific spin $(a/M_{H})|_R$ and energy radiated $E_{rad}$, as well as the predicted remnant horizon specific spin $(a/M_{H})|_{pred}$ and energy radiated $E_{rad}/M|_{pred}$ (based on a least-squares fit) for quasi-circular, equal-mass, equal-spin binaries with initial specific spins $(a/m)|_I$.} \\begin{ruledtabular} \\begin{tabular}{lllll}\\label{tab:spindat} $(a/m)|_I$ & $(a/M_H)|_R$ & $(a/M_{H})|_{pred}$ & $E_{rad}/M$ & $E_{rad}/M|_{pred}$\\\\\\hline $-0.757$ & $0.443\\pm0.001$ & $0.4430$ & $(2.2\\pm.01)\\%$& $2.20\\%$\\\\ $0.000$& $0.688\\pm0.001$ & $0.6878$&$(3.5\\pm0.1)\\%$ &$3.48\\%$\\\\ $0.1001$& $0.717\\pm0.001$ & $0.7169$&$(3.8\\pm0.1)\\%$&$3.79\\%$\\\\ $0.757$ & $0.890\\pm0.001$ & $0.8900 $&$(6.7\\pm0.2)\\%$ &$6.70\\%$\\\\ \\hline $-1.0$ & $***$ & $0.355$& $***$ &$2.2\\%$\\\\ $+1.0$ & $***$ & $0.946$& $***$ &$8.1\\%$\\\\ \\end{tabular} \\end{ruledtabular} \\end{table}  We have also found that the tidal coupling in black-hole binaries responsible for the transfers of angular momentum from the individual spins to orbital is relatively small. Nevertheless, we have been able to compute this transfer with sufficient accuracy to model its dependence on the initial spin. We find that the radiated angular momentum, which is a measure of the angular momentum transferred from spin to orbit, has a linear-plus-cubic dependence on the initial spins; in qualitative agreement both with the post-Newtonian (see Eq.~(\\ref{JdotPNSO})) and close-limit (see Eq.~(\\ref{Jcl})) predictions. The radiated energy has a quadratic-plus-quartic dependence on the initial spin; again, in qualitative agreement with the post-Newtonian (see Eqs.~(\\ref{EdotPN})) and close-limit (see Eq.~(\\ref{Ecl})) predictions.   Unfortunately the speculations of Price and Whelan~\\cite{Price:2001un} suggesting this spin-up effect might be strong, and even dominant, are not confirmed by our computations. We see neither the possibility of a binary stalling by this effect nor the dominance of the $\\ell=3$ mode suggested in Ref.~\\cite{Price:2001un}.        \\appendix"
    },
    "0608/astro-ph0608043_arXiv.txt": {
        "abstract": "We explore semi-complete self-similar solutions for the polytropic gas dynamics involving self-gravity under spherical symmetry, examine behaviours of the sonic critical curve, and present new asymptotic collapse solutions that describe `quasi-static' asymptotic behaviours at small radii and large times. These new `quasi-static' solutions with divergent mass density approaching the core can have self-similar oscillations. Earlier known solutions are summarized. Various semi-complete self-similar solutions involving such novel asymptotic solutions are constructed, either with or without a shock. In contexts of stellar core collapse and supernova explosion, a hydrodynamic model of a rebound shock initiated around the stellar degenerate core of a massive progenitor star is presented. With this dynamic model framework, we attempt to relate progenitor stars and the corresponding remnant compact stars: neutron stars, black holes, and white dwarfs. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608569_arXiv.txt": {
        "abstract": "We explore the behavior of the blue galaxy fraction over the redshift range $0.75\\le z\\le 1.3$ in the DEEP2 Survey, both for field galaxies and for galaxies in groups. The primary aim is to determine the role that groups play in driving the evolution of galaxy colour at high $z$.  In pursuing this aim, it is essential to define a galaxy sample that does not suffer from redshift-dependent selection effects in colour-magnitude space.  We develop four such samples for this study: at all redshifts considered, each one is complete in colour-magnitude space, and the selection also accounts for evolution in the galaxy luminosity function.  These samples will also be useful for future evolutionary studies in DEEP2.  The colour segregation observed between local group and field samples is already in place at $z\\sim 1$: DEEP2 groups have a significantly lower blue fraction than the field.  At fixed $z$, there is also a correlation between blue fraction and galaxy magnitude, such that brighter galaxies are more likely to be red, both in groups and in the field.  In addition, there is a negative correlation between blue fraction and group richness.  In terms of evolution, the blue fraction in groups and the field remains roughly constant from $z=0.75$ to $z\\sim 1$, but beyond this redshift the blue fraction in groups rises rapidly with $z$, and the group and field blue fractions become indistinguishable at $z\\sim 1.3$.  Careful tests indicate that this effect does not arise from known systematic or selection effects.  To further ensure the robustness of this result, we build on previous mock DEEP2 catalogues to develop mock catalogues that reproduce the colour-overdensity relation observed in DEEP2 and use these to test our methods.  The convergence between the group and field blue fractions at $z\\sim 1.3$ implies that DEEP2 galaxy groups only became efficient at quenching star formation at $z\\sim 2$; this result is broadly consistent with other recent observations and with current models of galaxy evolution and hierarchical structure growth. ",
        "introduction": "\\label{sec:intro} One of the most striking characteristics of the galaxy population is the well-known environmental segregation of the two main galaxy types: red, early-type galaxies with little ongoing star formation preferentially occur in galaxy groups and clusters, while blue, late-type galaxies with significant star-formation activity avoid such systems and preferentially populate the ``field''\\footnote{Throughout this paper, we shall use the term \\emph{field} to refer to those galaxies that are not in groups (i.e., isolated galaxies), not the full galaxy population}.  This observation has been recognized as a key to understanding galaxy formation and evolution for more than fifty years \\citep{SB51}.%   There is now overwhelming evidence that the galaxy population in clusters has evolved significantly with redshift down to the present day.  \\citet{BO84} were the first to present evidence that the fraction of blue galaxies, $f_b$, in clusters increases strongly with increasing $z$---the so-called Butcher-Oemler (BO) effect.  This basic result---an increased incidence of star-forming galaxies in distant clusters---has been replicated in numerous later studies using a variety of star-formation indicators, including cluster blue fractions \\citep{RS95, MD00, KB01, Ellingson01, Margoniner01},  emission-line galaxy fractions \\citep{Poggianti06}, and morphological fractions \\citep{ODB97, Couch98, vanD00, Fasano00, LOP02, Goto03}.  Such studies have also been extended to less massive galaxy groups \\citep{AS93, Wilman05b, MOL06}, which also appear to have proportionally more star-forming galaxies back in time. To be sure, there remain strong reasons to question the veracity of the BO effect as it was \\emph{originally} presented for galaxies in the \\emph{cores} of \\emph{rich} clusters (\\emph{e.g.}, \\citealt{Koo88, Smail98, Andreon06b}).  But it is now indisputable that clusters, on the whole, had proportionally more blue, star-forming, and morphologically late-type galaxies---in short, more star formation---at $z\\sim 0.5$ than they do at present.  It is tempting to conclude that this evolution is responsible for the substantial growth that has been observed in the number density of red galaxies since $z\\sim 1$ \\citep{Bell04b, Willmer06, Faber06}. But it is also important to note that, in addition to evolution \\emph{within} the group and cluster environment, the build-up of the red galaxy population could also be driven by an \\emph{increasing number density} of groups and clusters from the growth of structure. We shall try in this paper to shed light on the relative importance in DEEP2 of these two channels for red-galaxy formation.  A blue fraction that declines with time in groups and clusters is a natural consequence of hierarchical schemes for galaxy formation \\citep{Kauffmann95a, BFC96, Diaferio01, Benson01} if one assumes that groups and clusters play a role in quenching star formation.  Since these systems form at relatively late times, their member galaxies at intermediate redshift will have had less time, on average, to ``feel'' the effects of group membership; thus, a smaller fraction of them will have ceased forming stars.  Indeed, it has been shown \\citep{KB01} that the evolution of $f_b$ in clusters out to $z\\sim 0.4$ can be entirely accounted for by considering the effects of hierarchical structure formation and a declining universal star-formation rate.  For exactness in what follows, then, we shall use the phrase ``Butcher-Oemler effect'' to refer just to this effect---a decline in the typical blue fraction in groups and clusters that results \\emph{simply} from the increasing age of these systems with time, and \\emph{not} from any evolution in the quenching efficiency of groups and clusters. Although this working definition differs somewhat from the effect originally reported by \\citet{BO84}, it reflects usage that has become common in recent literature.  It is also worth noting that, under this definition, the BO effect may be stronger in some types of systems than in others (in clusters than in groups, for example), depending on the accretion rates and quenching efficiencies involved.  It remains unclear, however, which specific physical mechanisms are responsible for quenching star formation, over what timescales they act, and even whether these mechanisms are peculiar to groups and clusters. A variety of mechanisms % has been  proposed, including (\\emph{1}) galaxy mergers (\\emph{e.g}, \\citealt{TT72}), which occur primarily in galaxy groups (\\emph{e.g.}, \\citealt{CCM92}) and can trigger AGN feedback that sweeps merger remnants of their gas (\\citealt{SDH05}); (\\emph{2}) close, high-velocity galaxy encounters (``harrassment''; \\citealt{Moore96}), which occur primarily in clusters; (\\emph{3}) ram-pressure stripping of galaxy gas by the hot intracluster medium (\\citealt{GG72}); and (\\emph{4}) the less violent process known as ``strangulation'' in which gas accretion onto galaxy discs is cut off, either by stripping of galaxies' gaseous haloes (\\citealt{LTC80}) or by feedback from low-luminosity AGN \\citep{Croton06}, and star formation ends when the remaining supply of gas has been exhausted.  There is evidence that each of these processes is at work on some level; it may be that no single mechanism bears primary responsibility for galaxy evolution. Disentangling how and to what degree each of the above processes shapes the galaxy population requires detailed observations of galaxies over a wide range of redshifts and environments.   The observed evolution is rather complicated in its details, however. For example, the well-known morphology-density relation for galaxies in clusters (\\emph{e.g.},  \\citealt{Dressler80}) and groups (\\emph{e.g.} \\citealt{PG84}) evolves strongly with increasing redshift \\citep{Dressler97}, and there is evidence that the evolution in intermediate-density environments has occured more recently than in high-density environments \\citep{Smith05}. It also appears that the morphological mix of cluster early-types has changed with time \\citep{Fasano00, Postman05}.  Complicating things further, cluster blue fractions display a substantial system-to-system scatter at all epochs \\citep{BO84, Smail98, MD00, Goto03, DeP04}.  This scatter can be ascribed in part to the observed trends between blue fraction and cluster mass, \\emph{i.e.}, velocity dispersion \\citep{NKB88}, inferred virial mass \\citep{Martinez02}, richness \\citep{Margoniner01, Goto03, TPA05}, or total optical luminosity \\citep{Weinmann06}; the correlation between $f_b$ and mass also appears to evolve with redshift \\citep{Poggianti06}.  The scatter may also arise in part from a correlation between $f_b$ and a cluster's degree of dynamical relaxation (\\emph{e.g.}, \\citealt{MRU00}).   The scatter in $f_b$ might mask or enhance the redshift evolution, depending on the cluster selection criteria used \\citep{NKB88, AE99}.  Indeed, \\citet{Smail98} observe no BO effect out to $z\\sim 0.25$ in a sample of luminous X-ray clusters (but see also \\citealt{Fairley02}), and many authors have commented on the presence of individual massive clusters at high redshift with low $f_b$ and well-formed red sequences (\\emph{e.g.} \\citealt{Koo81, Homeier05, Andreon06}).  It is clear that a robust study of galaxy evolution in groups and clusters requires a sample that contains a broad range of systems and is uniformly selected at all redshifts.  Until relatively recently, studies of galaxy evolution in clusters were limited to cluster samples pre-selected from optical imaging or X-ray maps; such samples are prone to selection effects that bias the galaxy sample as a function of $z$.  The advent of large, densely sampled galaxy redshift surveys like the 2-degree Field Galaxy Redshift Survey (2dFGRS) and the Sloan Digital Sky Survey (SDSS), however, has allowed selection of large, low-redshift cluster samples directly from the galaxy distribution in redshift space, where redshift-dependent selection effects can be well understood and accounted for.  Several authors have studied the properties of the cluster galaxy population in the 2dFGRS \\citep{Balogh04, DeP04} and SDSS \\citep{Goto05a, Goto05b, Weinmann06, Quintero06}.  Within SDSS, there have even been detections of an evolving $f_b$ in clusters \\citep{Goto03, MOL06} and of an evolving morphology-density relation \\citep{GYTO04}.  With the high-redshift DEEP2 Galaxy Redshift Survey (\\citealt{DGN04}, Faber et al. in prep.) nearing completion, it is now possible to extend these studies to $z\\sim 1$.  DEEP2 has now yielded a catalogue of several thousand groups and small clusters over a wide range of masses at $z\\sim 1$ \\citep{Gerke05}.  A study of the DEEP2 sample by \\citet{Cooper06a} has already revealed that, \\emph{without} explicitly considering groups,  the correlations between galaxy properties and the density of the nearby galaxy distribution are qualitatively similar at $z\\sim 1$ to what is observed locally (\\emph{e.g.},  \\citealt{Balogh04, Hogg04}), although the relations differ in detail.  Interestingly, recent studies have shown that these correlations evolve with redshift, becoming stronger over time \\citep{Nuijten05, Cucciati06, Cooper06b}, and that the color-density relation in the \\emph{local} Universe can be ascribed almost entirely to mechanisms acting in groups and clusters \\citep{BBH06}. The primary goal of this paper, then, is  to establish the role that groups and clusters play, if any, in driving the evolution of the DEEP2 galaxy population.   In particular, we shall explore the possibility that group and cluster environments are responsible for the build-up of red galaxies observed over the DEEP2 redshift range \\citep{Bell04b, Willmer06, Faber06}; by so doing, we will try to shed light on the physical mechanisms driving the evolution. Although this study will focus on the evolutionary effects of galaxies' membership in groups and clusters (that is to say, of their dark matter halo masses), there is important complementary information to be gained by studying the effects of the local density of galaxies.  Such a study is undertaken in a companion paper to this one by \\citet{Cooper06b}. For now, however, we will measure $f_b$  for galaxies in groups and clusters, as a function of redshift, and compare it to that found in the general field (\\emph{i.e.}, in the galaxy population outside of groups).  We will also investigate the relation between $f_b$ and cluster properties in order to gain insight into the processes at work and also to understand and control any systematic selection effects.    The thoughtful reader may find it perverse to use the \\emph{blue} fraction to study the evolution of \\emph{red} galaxies, but $f_b$ has significant historical precedent in the study of cluster galaxy populations, so we use it for consistency with earlier studies. In any case, the possibly more natural \\emph{red} fraction statistic is simply $f_r = 1-f_b$.  We have chosen to consider galaxy \\emph{colour} rather than other indicators of galaxy type like morphology or [O II] line emission primarily because it can be measured accurately for the largest number of DEEP2 galaxies, allowing for robust statistics.  The DEEP2 ground-based imaging lacks the resolution necessary for morphological classification of most galaxies at $z\\sim 1$, and there is \\emph{HST} imaging for only a small fraction of the sample. Emission line strength can be measured accurately only in spectra with sufficient signal-to-noise ratio; not all DEEP2 galaxies meet this criterion.  Also, it has recently been shown that [O II] emission may not give a clear indication of star formation, especially in red galaxies \\citep{Yan06}. Regardless, there is a strong and relatively tight correlation between [O II] equivalent width and galaxy colour% (\\emph{e.g.}, \\citealt{Weiner05, Cooper06a}), so the two statistics should give similar results.   We shall proceed as follows.  In \\S \\ref{sec:DEEP2}, we describe the DEEP2 survey and the DEEP2 group catalogue, and in \\S \\ref{sec:mocks} we discuss the construction of mock DEEP2 galaxy catalogues appropriate for this study.  We describe our galaxy selection criteria and incompleteness corrections and present the precise definition of $f_b$ used here in \\S \\ref{sec:selection}. Measurements of $f_b$ in groups and the field are presented in \\S \\ref{sec:results}, and we perform tests with the mock catalogues in \\S~\\ref{sec:mock_tests}.  We discuss the implications  of our results for galaxy evolution models in \\S \\ref{sec:discussion}, and in \\S \\ref{sec:conclusion} we summarize our results and conclusions.  Readers uninterested in the details of our selection methods and robustness checks should skip to \\S\\S~\\ref{sec:sampledef}, \\ref{sec:results} and \\ref{sec:discussion}. Throughout the paper we assume a flat, $\\Lambda$CDM cosmology with $\\Omega_M=0.3$. ",
        "conclusions": "\\label{sec:conclusion}  We have used the DEEP2 group catalogue of \\citet{Gerke05} to study the blue fraction of galaxies in groups and the field at $z\\sim 1$.  After creating four samples that are carefully designed to be free from colour-dependent and redshift-dependent selection effects, we probe the dependence of $f_b$ on a galaxy's environment (group or field), on galaxy luminosity, on group properties, and on redshift.  First, over the redshift range $0.75\\le z <1$, there is a significant difference in $f_b$ between the DEEP2 group and field populations, with a substantially lower $f_b$ in groups than in the field. The well-known colour segregation observed between nearby group and field galaxy populations is thus already in place at $z\\sim 1$.   In addition, there is a negative correlation between $f_b$ and group richness over this redshift range, echoing local correlations between galaxy colour and various proxies for halo mass.  A trend also exists at these redshifts between $f_b$ and the luminosity of the galaxies being considered, with fainter galaxies being more likely to be blue; the strength of this trend does not appear to depend on whether those galaxies are in groups or the field.  No siginificant evolution in $f_b$ occurs in the DEEP2 group or field galaxy populations over the range $0.75\\le z \\la 1$; however, at $z\\ga 1$, there is dramatic, statistically significant evolution in the group population, with $f_b$ rising rapidly as $z$ increases. Most notably, the value of $f_b$ in groups becomes indistinguishable from the field value at $z\\sim 1.3$.  To assess the robustness of these results, we have constructed mock catalogues that replicate the colour-environment correlations in the DEEP2 data.  Tests with these mock catalogues indicate that the evolution and convergence are not artifacts of the group-finding algorithm.  Moreover, these tests suggest that the observed evolution can be explained partially---but not entirely---by evolution in the distribution of galaxy environments, \\emph{i.e.} by the growth of large-scale structure, coupled with the known correlation between environment and galaxy properties.  However, to explain the strong evolution and convergence, some further mechanism is required.  Because the $f_b$ evolution appears to weaken at $z\\la 1$, we conclude that it is different in nature from the classical BO effect, in which $f_b$ in clusters rises dramatically from $z=0$ out to $z\\sim 0.5$.  Indeed, the BO effect is a natural outcome of hierarchical structure formation, \\emph{provided that} groups and clusters are already efficient at quenching star formation in the galaxies that fall into them.  But there is a \\emph{convergence} between  $f_b$ values in groups and in the field at $z\\sim 1.3$, which implies that DEEP2 groups were, at some epoch, \\emph{not} especially efficient at quenching infalling galaxies.  In particular, if it takes $\\sim 1$ Gyr for a quenched galaxy to join the red sequence, our results indicate that the groups in this study only became suitable environments for quenching at $z\\sim 2$.  This conclusion is broadly consistent with current semi-analytic models of galaxy formation \\citep{Croton06, Bower06, Cattaneo06, KJS06}. In those models, the central galaxies of dark matter haloes are quenched when low-level AGN feedback becomes strong enough to stop gas in the halo from cooling and condensing onto the galaxy to form stars; this process becomes efficient at a threshold halo mass of a few times $10^{12} M_\\odot$ (\\citealt{DB06, Cattaneo06}; Croton et al in prep.). Our mock catalogues imply that the groups observed in DEEP2 have masses in the range $5\\times 10^{12} M_\\odot \\la M \\la 5\\times 10^{13} M_\\odot$.  A typical halo in this mass range at $z\\sim 1.3$ will drop below the threshold mass for quenching $\\sim 1h^{-1}$ Gyr earlier, at $z\\sim 2$.  This is a potential explanation for the convergence in the group and field $f_b$ values at $z\\sim 1.3$.  Present-day SAMs may have difficulty explaining these observations in detail, however, particularly because they universally assume that \\emph{satellite} galaxies are quenched in all haloes regardless of mass or redshift.  In the future, it will be worthwhile to make detailed comparisons between our observations and the precise predictions of SAMs, in order to refine the physical description of galaxy formation that they provide."
    },
    "0608/astro-ph0608039_arXiv.txt": {
        "abstract": "Critical insights on galaxy evolution stem from the study of bars. With the advent of high redshift $HST$  surveys that trace bars in the rest-frame optical band out to $z\\sim$~1, it becomes increasingly important to provide a reference baseline for bars at $z\\sim$~0 in the optical band. We present results on bars at $z\\sim$~0 in the optical and  near-infrared bands, based on 180 spirals in the OSUBSGS survey. (1)  The deprojected  bar fraction at $z\\sim$~0 is $f_{\\rm NIR1}$~$\\sim$~60\\%~$\\pm$~6\\% in the near-infrared $H$ band, and $f_{\\rm optical1}$$\\sim$~44\\%~$\\pm$~6\\%  in the optical $B$-band images. The latter likely miss  bars obscured by dust and star formation. (2)~The results before and after deprojection are similar, which is encouraging for high redshift studies that forego deprojection. (3) Studies of bars at $z \\sim$~0.2--1.0  (lookback times of 3--8 Gyr) have reported an optical bar fraction of $f_{\\rm optical2}$$\\sim$~30\\% $\\pm$ 6\\%, after applying cutoffs in absolute magnitude  ($M_{\\rm V}<$-19.3), bar size ($a_{\\rm bar} \\ge$ 1.5 kpc), and bar ellipticity ($e_{\\rm bar} \\ge$~0.4)  in order to ensure a complete sample, adequate spatial resolution, and reliable bar identification out to $z\\sim$~1. Applying these exact cutoffs in magnitude,  bar size, and bar ellipticity to the  OSUBSGS data yields  a comparable  optical $B$-band bar fraction at $z\\sim$~0 of $f_{\\rm optical3}$$\\sim$~34\\%~$\\pm$~6\\%. This rules out scenarios where the optical bar fraction in bright disks declines strongly with redshift. (4) We investigate bar strengths at  $z \\sim$~0 using the maximum bar ellipticity ($e_{\\rm bar}$) as a guide. Most  ($\\sim$~70\\%) bars have moderate to high ellipticity (0.50~$ \\le  e_{\\rm bar} \\le$~0.75), and only a  small fraction (7\\%--10\\%)  have 0.25~$ \\le  e_{\\rm bar} \\le$~0.40. There is no bimodality in the distribution of $e_{\\rm bar}$. The $H$-band bar fraction and  $e_{\\rm bar}$ show no substantial variation across RC3 Hubble types Sa to Scd. (5)~RC3 bar types should be used with caution. Many  galaxies with RC3 types `AB' turn out to be unbarred and RC3 bar classes  `B' and `AB' have a significant overlap in  $e_{\\rm bar}$. (6)~Most  (68\\% in $B$ and 76\\% in $H$)  bars have sizes  below 5 kpc. Bar and disk sizes correlate, and the ratio ($a_{\\rm bar}$/$R_{\\rm 25}$) lies primarily in the range 0.1 to 0.5. This suggests that the growth of bars and disks is intimately tied. \\end {abstract} ",
        "introduction": "Stellar bars  are recognized as the  most important internal factor that  redistributes the angular momentum of the baryonic and dark matter  components of  disk galaxies (e.g., Weinberg 1985; Debattista \\& Sellwood 1998, 2000; Athanassoula 2002; Berentzen, Shlosman, \\& Jogee 2006), thereby driving their dynamical and secular evolution. Bars efficiently drive gas from the outer disk to the central few hundred parsecs  and  are observed to feed central starbursts in local galaxies (Elmegreen 1994; Knapen \\etal 1995; Hunt \\& Malakan 1999; Jogee \\etal 1999;  Jogee, Scoville, \\& Kenney 2005). It remains a matter of contention whether large-scale bars relate to AGN activity in galaxies, given the reduction by several orders of magnitude needed in the specific angular momentum of gas before it can feed a central black hole, and conflicting observational results (see review by Jogee 2006 and references therein;  also Mulchaey \\& Regan 1997; Knapen \\etal  2000; Laine \\etal 2002; Laurikainen \\etal 2004). In several galaxies, bar-driven gas inflows appear intimately tied to the formation of disky, high $v/\\sigma$  stellar components in the inner kpc, or `pseudobulges' (Kormendy 1993; Jogee 1999; review by Kormendy \\& Kennicutt 2004; Jogee, Scoville, \\& Kenney 2005; Athanassoula 2005). Furthermore, the orbital structure of bars can lead to the observed peanut-shaped and boxy bulges in inclined galaxies (Combes \\etal 1990; Pfenniger \\& Norman 1990;  Bureau \\& Athanassoula 2005; Athanassoula 2005; Martinez-Valpuesta \\etal 2006; Debattista \\etal 2006).   Earlier ${\\it Hubble~Space~Telescope~(HST)}$  studies at optical wavelengths (e.g., Abraham \\etal 1999) reported a paucity of stellar bars and a sharply declining optical bar fraction at intermediate redshifts $z >$~0.5. Studies at near-infrared (NIR) wavelengths also found a low bar fraction, but the authors rightly concluded that the large effective point spread functions (PSFs) of the NIR camera only allowed the detection of  large bars whose semi-major axes  exceeded  $0.9\\arcsec$, corresponding to 7.2 kpc\\footnote{We assume in this paper a flat cosmology with $\\Omega_M = 1 - \\Omega_{\\Lambda} = 0.3$ and $H_{\\rm 0}$ =70~km~s$^{-1}$~Mpc$^{-1}$.} at $z\\sim$~1.0 (Sheth \\etal 2003). Recent works based on large optical surveys have  now demonstrated the abundance of bars  at intermediate redshifts $z \\sim$~0.2--1.0, corresponding to lookback times of  3--8 Gyr (Elmegreen \\etal 2004; Jogee \\etal 2004;  Zheng \\etal  2005; Sheth \\etal in preparation). The fundamental issue of how robust bars are, and the associated implications for bar-driven  evolution in disks  over the last 10 Gyr, remains open (e.g., Jogee \\etal  2004; Shen \\& Sellwood 2004; Athanassoula, Lambert, \\& Dehnen  2005;  Bournaud \\etal 2005; Berentzen, Shlosman, \\& Jogee 2006; Berentzen \\& Shlosman 2006; Martinez-Valpuesta \\etal 2006; Debattista \\etal 2006).   In order to put bars in a cosmological context, it now behooves us to characterize the frequency  and impact of bars by applying the {\\it same quantitative} methods to large samples at $z\\sim$~0 and at higher redshifts. Spurred by these considerations, we characterize in this paper the frequency and structural properties of bars in the local Universe at optical and NIR  wavelengths, by ellipse-fitting the $B$ and $H$ images of the OSU Bright Spiral Galaxy Survey (OSUBSGS; Eskridge et al. 2002) of 180 spirals. The first goal of  this study is   to  provide {\\it quantitative} characterizations  of the bar fraction $f_{\\rm bar}$  (defined as the fraction of disk galaxies that are barred) and structural properties (sizes, ellipticities, etc.) of bars  at $z\\sim$~0,  as a function of wavelength, Hubble types, and host galaxy properties. Furthermore, with the advent of high redshift $HST$ surveys, such as the  Tadpole field (Tran \\etal 2003), the Galaxy Evolution  from Morphology and SEDs (GEMS; Rix \\etal 2004), the Great Observatories Origins Deep Survey (GOODS; Giavalisco \\etal 2004), and COSMOS (Scoville \\etal 2006), which trace bars in the rest-frame optical band out to $z\\sim$~1, it becomes increasingly important to provide a reference baseline for bars at $z\\sim$~0 in the optical band. Thus, a second  goal of our study is to provide {\\it a rest-frame optical }  $z \\sim$~0 point for bars  based on  ellipse fits, in order to directly compare  with studies of intermediate-redshift bars (Jogee \\etal  2004; Elmegreen \\etal 2004; Zheng \\etal  2005; Sheth et al. in preparation) that also use ellipse fits. In particular, we use in this paper the same procedure of ellipse fits ($\\S$~3.1) and the same quantitative  characterizations ($\\S$~3.3) of bars that were  applied by Jogee \\etal  (2004) to bars at intermediate redshifts ($z \\sim$~0.2--1.0) in  the GEMS survey.  Several studies have used the OSUBSGS to  gauge bars in the local Universe (e.g., Eskridge et al. 2000; Block et al. 2002; Whyte \\etal 2002; Buta \\etal 2005), but they  differ significantly from our study and cannot meet our two goals. Eskridge \\etal (2000) visually classified bars in the $H$ band, and in the $B$ band, they used the Third Reference Catalog of Bright Galaxies (de Vaucouleurs et al. 1991; hereafter RC3) visual  bar classes. Such visual classifications form an invaluable first step, but by definition, are subjective and difficult to compare with results from other studies. Block \\etal (2002) and later Buta \\etal (2005)  applied the gravitational torque $Q_{\\rm b}$ method, based on Fourier amplitudes, to $H$-band images of 163 and 147 OSUBSGS galaxies, respectively. This quantitative method is less subjective than visual classification, but the results of Block \\etal (2002) and Buta \\etal (2005) cannot be compared to intermediate redshift studies for two reasons. First, the latter studies were based on the {\\it HST} Advanced Camera for Surveys (ACS) data and trace the rest-frame optical properties of bars, while Block \\etal (2002) and Buta \\etal (2005) deal with the rest-frame NIR. Second, it is non-trivial to derive  $Q_{\\rm b}$ for intermediate-redshift galaxies because of resolution and signal-to-noise limitations.  Whyte \\etal (2002) fitted ellipses to $B$-band images of only 89 of the 180 OSUBSGS galaxies, and do not provide a distribution of bar properties as a function of Hubble type. Our present study complements these existing studies by ellipse fitting $B$-band and $H$-band images of  all  180 OSUBSGS galaxies, and performing a comprehensive, statistically significant analysis of  barred galaxies in the local Universe. It complements the ongoing analysis (Barazza, Jogee, \\& Marinova 2006) \\ of local bars based on a sample of 5000 galaxies in the Sloan Digitized Sky Survey (SDSS).    The outline of this paper is as follows. $\\S$~2 discusses the sample selection based on the OSUBSGS survey (Eskridge \\etal 2002). $\\S$~3 describes  the ellipse-fitting method, the criteria used for identifying bars, and  deprojection of images and profiles to face-on. In  $\\S$~4.1--4.4, we present results on the  bar fraction at $z\\sim$~0, its dependence on Hubble type, the distribution of bar sizes and strengths as characterized by ellipse-fitting, and the variation of bar properties along the Hubble sequence. Results  are presented both before and after deprojection to face-on. In  $\\S$~4.5, we present a first-order comparison of the bar fraction and properties at $z\\sim$~0 from OSUBSGS to those derived at  $z\\sim$~0.2--1.0 or lookback times of  3--8 Gyr from GEMS (Jogee \\etal  2004) and the Tadpole field (Elmegreen \\etal 2004). In $\\S$~4.6, we discuss the constraints set by our results for theoretical models addressing the  robustness of bars, and  the assembly of the Hubble sequence over cosmological times. $\\S$ 5 presents the summary and conclusions.   This paper is the first in a series of three based on the OSUBSGS. Paper II (Marinova \\etal in prep) will address the bulge properties and activity of barred and unbarred galaxies in the OSUBSGS sample. In paper III, we will present simulations that artificially redshift the  rest-frame optical and NIR images of the local OSUBSGS sample out to   $z\\sim$~1--2, in order to assess the impact of redshift-dependent systematic effects on the recovery rate of bars  in surveys conducted by current and future facilities in the optical and IR, such as the planned Wide Field Camera 3 (WFC3) and the James Webb Space Telescopes (JWST). ",
        "conclusions": "With the advent of high redshift $HST$ surveys, such as the Tadpole Field, GEMS, GOODS, and COSMOS, which trace bars in the rest-frame optical band out to $z\\sim$~1, it becomes increasingly important to provide a reference baseline for bars at $z\\sim$~0 in the optical band. Motivated by these considerations, we characterize the frequency and structural properties of bars  at $z\\sim$~0 in the  optical and NIR bands, by   ellipse-fitting the $B$ and $H$ images of 180 spirals in the OSUBSGS (Eskridge \\etal 2002), and applying quantitative criteria in order to identify and characterize bars. We determine the inclination of the outer disk and exclude highly inclined  ($i >$~60$^\\circ$) galaxies to derive a sample S4 of 136 moderately inclined  spirals. For this sample, we derive bar properties  both before and after deprojection to face-on. Our study complements existing work on OSUBSGS based on Fourier amplitudes (Block \\etal 2002; Buta \\etal 2005) and visual classification (Eskridge \\etal  2000), and it can be compared with studies (Jogee \\etal  2004; Elmegreen \\etal 2004; Zheng \\etal  2005) of intermediate redshift ($z \\sim$~0.2--1.0) bars employing the same ellipse-fitting methodology. Our results are summarized below.   {\\it \\rm (1)~The optical and NIR bar fraction at $z\\sim$~0: } For our sample, which  is dominated by galaxies with $M_{\\rm V} \\sim$~-20 to -22, we find a deprojected  bar fraction at $z\\sim$~0 of $f_{\\rm NIR1}$~$\\sim$~60\\%~$\\pm$~6\\% in the near-infrared $H$ band, and $f_{\\rm optical1}$$\\sim$~44\\%~$\\pm$~6\\%  in the optical $B$-band images. The latter  likely miss  bars obscured by dust and star formation. Deprojection does not make any significant changes to the global  $B$- and  $H$- band bar fractions, which are 45\\% and 58\\% before  deprojection, and change by only a factor of 0.97 and 1.03, respectively, after deprojection. This is encouraging for large studies of bars at intermediate redshift (e.g., Jogee \\etal  2004, Elmegreen \\etal 2004, Zheng \\etal  2005), where deprojection is not performed.  {\\it   \\rm (2) Comparison of optical properties of bars at $z \\sim$~0  and at  intermediate redshifts: } Studies of bars  at $z \\sim$~0.2--1.0  (lookback times of  3--8 Gyr) based on  $HST$ ACS observations in the  Tadpole field, the GEMS and GOODS fields, and COSMOS surveys trace bars in the {\\it rest-frame optical}. $R/I$ to $V/B$ bands (7586 \\AA~to 4550 \\AA). Therefore, in order to avoid the pernicious effects of bandpass shifting, it is essential that ACS studies of bars at $z \\sim$~0.2--1.0  compare their rest-frame optical results to the optical bar fraction at $z \\sim$~0,  rather than to the significantly higher  NIR bar fraction at  $z \\sim$~0.  \\noindent Furthermore, at $z \\sim$~0.2--1.0, it is essential to apply  cutoffs  in absolute magnitude, bar size, and  bar ellipticity in order to ensure a complete sample, adequate spatial resolution, and reliable bar identification. After applying  cutoffs in absolute magnitude ($M_{\\rm V}<$-19.3), bar size ($a_{\\rm bar} \\ge$ 1.5 kpc), and bar ellipticity ($e_{\\rm bar} \\ge$~0.4), Jogee \\etal (2004) found a  rest-frame  optical bar fraction of  $f_{\\rm optical2}$$\\sim$~30\\%~$\\pm$ 6\\% at  $z \\sim$~0.2--1.0. A constant and similar optical bar fraction  (23\\% to 40\\%) out to  $z \\sim$~1 is also reported by  Elmegreen \\etal (2004). In order to derive the equivalent optical bar fraction for comparison at $z\\sim$~0, we applied the exact same cutoffs to the OSUBSGS optical data. With a cut off of $M_{\\rm V}<$-19.3, the optical bar fraction  $z\\sim$~0 drops from 45\\% (61/136) to 43\\%. Applying a further cutoff of $a_{\\rm bar} \\ge$ 1.5 kpc makes it drop  to 36\\%. Finally, a third cutoff  of $e_{\\rm bar} \\ge$~0.4 reduces optical $B$-band bar fraction at $z\\sim$~0  to $f_{\\rm optical3}$$\\sim$~34\\%~$\\pm$~6\\%. The result that  $f_{\\rm optical2}$ is comparable to $f_{\\rm optical3}$ rules out scenarios where the {\\it optical}  bar fraction in {\\it bright} disks  declines strongly with redshift. It allows for models where the  optical bar fraction is either constant, or  rises with redshift.   {\\it  \\rm (3) Distribution of bar strengths  $z\\sim$~0  as characterized by ellipse-fitting:} In this study, we use the maximum  bar ellipticity   $e_{\\rm bar}$ from ellipse-fits as a partial measure of the bar strength. Only a  very small proportion (7\\% in $B$; 10\\% in $H$) of bars  are very weak  as characterized by  $e_{\\rm bar}$ from ellipse fits (0.25~$ \\le  e_{\\rm bar} \\le$~0.40), while the majority of bars  (70\\% in $B$; 71\\% in $H$) have moderate to high ellipticities  (0.50~$ \\le  e_{\\rm bar} \\le$~0.75). We find  no evidence for  bimodality in the distribution of bar strength, as characterized by  $e_{\\rm bar}$ in the $B$ or $H$ bands, in agreement with Buta \\etal (2005).    {\\it  \\rm (4) Bar fraction and strength, as characterized by  ellipse-fitting, as a function of RC3 Hubble type at $z\\sim$~0:} The  deprojected bar fraction is 60\\% in $H$ and  44\\% in $B$, confirming the ubiquity of local bars. In  the $B$ band, the bar fraction is lower with respect to the $H$ band by  ~$\\sim$~1.2--1.5  for Hubble types S0s to Scs, and by $\\sim$~2.5  for Sds/Sms. This is consistent with the higher obscuration in  dusty, gas-rich late types. The bar fraction and bar strength,  as characterized by  $e_{\\rm bar}$, in the  $H$ band  shows no systematic variation across Hubble types Sa to Scd.    {\\it  \\rm (5) Comparison with RC3 visual bar classes: } Of the 42, 47, and 46 galaxies in our sample that have an RC3 visual bar class of  `A' (unbarred), `B' (strongly barred), and `AB' (weakly barred), respectively, our quantitative characterization ($\\S$~3.3) shows that 5\\%, 85\\%, and 41\\% host bars in $B$-band images and 19\\%, 87\\%, and 65\\%  host bars in $H$-band images. Thus, quantitative characterization of bars differs significantly from RC3 bar classes for the RC3 bar class `AB'. Furthermore, the mean bar strength, as characterized by the maximum  bar ellipticity $e_{\\rm  bar}$,  is higher for RC3 visual class `B' than for class `AB', but the two classes have significant overlap in the range $e_{\\rm  bar} \\sim$~0.5--0.7. Thus, RC3 bar types should be used with caution and may be misleading.    {\\it  \\rm (6)~Sizes of bars and disks at $z\\sim$~0: } The sizes or semi-major axes  $a_{\\rm bar}$  of large-scale bars in the local Universe lie in the range $\\sim$~1 to 14 kpc, with  the majority  of bars (68\\% in $B$ and 76\\% in $H$) having  $a_{\\rm bar} \\le$~5 kpc. Bar and disk sizes are correlated   with an average slope of $\\sim$~0.9, albeit with a large scatter of several kpc in bar size at a given disk size.  The ratio ($a_{\\rm bar}$/$R_{\\rm 25}$) lies primarily in the range 0.1 to 0.5, with only a minority of galaxies having larger values out to 0.95. The correlation  between bar and disk sizes, and the narrow range in $a_{\\rm bar}$/$R_{\\rm 25}$ suggests that the growths of the bar and disk may be intimately tied. The fact that ($a_{\\rm bar}$/$R_{\\rm 25}$) is generally well below 1.0  suggests that the  CR of disk galaxies  lies well inside their  $R_{\\rm 25}$ radius, assuming that  bars end near the CR.     {\\it   \\rm (7) Constraints on the robustness of bars: } Our findings that the majority (60\\%) of spirals are barred in the infrared  and that most  ($\\sim$~71\\%--80\\%) of  these bars have primarily moderate to high ellipticities (0.50~$ \\le  e_{\\rm bar} \\le$~0.80) suggest that DM halos of present-day spirals have at most a mild triaxiality, with a maximum equatorial axis ratio $b$/$a$~$\\sim$~0.9 in  the potential. We also found that the bar fraction and mean bar strength  (as characterized by the maximum  bar ellipticity $e_{\\rm  bar}$) are relatively constant  across Hubble types Sa to Scd, and there is no  bimodality in $e_{\\rm  bar}$. Taken together, our results are easily reconciled with scenarios where bars  in  present-day  galaxies are  relatively robust against the range in  gas mass fractions, gas inflows, and CMC components  present across Hubble types Sa to Scd. Our results do not  necessarily rule out models where bars are easily destroyed by bar-driven gas inflows. They do, however, imply that  if such an easy destruction occurs, then  there  must be  a very efficient  mechanism that not only regenerates  bars on a short timescale,  but is also very well tuned to the bar destruction rate so that it can reproduce the observed  constant  optical bar fraction in bright galaxies over the last 8 Gyr.   S.J. and I.M. acknowledge support from NSF grant AST-0607748, NASA LTSA grant NAG5-13063, as well as $HST$ grants G0-1048  and G0-10395 from STScI, which is operated by AURA, Inc., for NASA, under NAS5-26555. The Ohio State University Bright Spiral Galaxy Survey, was funded by grants AST-9217716 and AST-9617006 from the United States National Science Foundation, with additional support from the Ohio State University. We thank Pat Shopbell, Peter Teuben, and Stuart Vogel for their assistance with the  Zodiac and MIRIAD packages; Seppo Laine for sharing his deprojection code from Laine et al. (2002); and James Davies for help with IRAF and IGI visualization  routines. We also thank Paul Eskridge, Ron Buta, Fabio Barazza, Debbie Elmegreen, Isaac Shlosman, Ingo  Berentzen, Seppo Laine, Juntai Shen, Victor Debattista, Lia Athanassoula, Francoise Combes, Frederic Bournaud, Jerry Sellwood, and Johan Knapen for useful discussions."
    },
    "0608/astro-ph0608513_arXiv.txt": {
        "abstract": "{We present a quantitative comparison between the Horizontal Branch morphology in the core of the Sagittarius dwarf spheroidal galaxy (Sgr) and in a wide field  sampling a portion of its tidal stream (Sgr Stream), located tens of kpc away from the center of the parent galaxy. We find that the Blue Horizontal Branch (BHB)  stars in that part of the Stream are five times more abundant than in the Sgr core, relative to Red Clump stars. The difference in the ratio of BHB to RC stars between the two fields is significant at the $\\simgt 4.8\\sigma$ level. This indicates that the old and metal-poor  population of Sgr was preferentially stripped from the galaxy in past peri-Galactic passages with respect to the intermediate-age metal rich population that presently dominates the bound core of Sgr, probably due to a strong radial gradient that was settled within the galaxy before its disruption. The technique adopted in the present study allows to trace population gradients along the whole extension of the Stream. ",
        "introduction": "The Sagittarius dwarf spheroidal galaxy (Sgr dSph; Ibata et al. \\cite{iba1}) provides an excellent case for the study of the process of tidal disruption and accretion of a dwarf satellite into a large galaxy. Its huge tidal tails form a coherent and  dynamically cold filamentary structure (hereafter Sgr Stream) extending for tens of kpc from the parent galaxy that has been probed with many different tracers (see, among the others, Ibata  et al. \\cite{iba2}, Newberg et al. \\cite{sdss} - hereafter N02, Majewski et al. \\cite{maj03} - hereafter M03, Mart\\'inez-Delgado et al. \\cite{david} - hereafter MD04, Vivas et al. \\cite{vivas}, Belokurov et al. \\cite{belok}, and references therein).    \\begin{figure} \\centering \\includegraphics[width=8.5cm]{HBSELN_lr.ps} \\caption{CMDs focused on the Horizontal Branch for fields in the core of Sgr (upper panel) and in the Sgr Stream (lower panel). The color strips adopted to select candidate RC and BHB stars are enclosed within vertical lines and labeled. The filled circles in the CMD of SGR34 are the ab type RR Lyrae counter-identified in the nearby SGRW field from the catalogue of Cseresnjes (\\cite{cser}). } \\label{HBSEL} \\end{figure}    It is quite clear that the progenitor of the presently observed Sgr remnant was a quite remarkable galaxy, possibly similar to the LMC (Monaco et al. 2005, hereafter M05).  While the stellar content of the bound core of the galaxy is dominated by an  intermediate-old metal-rich population ($[M/H]\\sim -0.5$, M05, Bellazzini et al. \\cite{conun}), the globular clusters possibly associated with the Stream are predominantly old and metal-poor (Bellazzini, Ferraro \\& Ibata \\cite{clus}; see also Belokurov et al. \\cite{belok}). M03 and MD04 provided the first indications of the possible presence of a metallicity gradient along the Stream, suggesting that the stars lost in previous peri-Galactic passages, were, on average, more metal poor than those recently lost or still bound to the main body. Chou et al. (\\cite{chou} - hereafter C06 ), studying the chemical abundances of relatively nearby Stream stars, have shown that the metallicity distribution within the Stream is significantly different from that of the main body of Sgr, the former being skewed toward lower metallicities. These results strongly suggest that the stellar content of the Sgr progenitor could have been quite different from the present-day remnant, and indicates that it may be possible to trace the change of stellar content along the Stream (see C06).  The Horizontal Branch (HB) morphology is a very powerful tool to study population gradients within galaxies (see Harbeck et al. 2001), since the color of HB stars depends very strongly on their age and metal content (see Fusi Pecci et al. \\cite{ffp1} for discussion and references). In particular, HB stars can be located to the {\\em blue} of the RR Lyrae instability strip (Blue HB stars) only  if they are very old (age $>$ 10 Gyr) and (typically) metal-poor, while HB stars lying to the {\\em red} of the instability strip (Red Clump stars) must be (comparatively) young  or  metal-rich, or both (see Monaco et al. \\cite{lbhb}, and B06 for the  interpretation of the HB  morphology of Sgr in this context). Monaco et al.(\\cite{lbhb})  noted that both the BHB and RC sequences observed at the center of Sgr have a clear counterpart in the Color-Magnitude Diagrams (CMDs) of Stream fields obtained by N02 using data from the Early Data Release of the Sloan Digital Sky Survey (SDSS). Here we perform a direct comparison of the relative abundance of BHB and RC stars in the Sgr core and in one of the Stream fields studied by N02, to search for a population gradient. While the results we obtain are qualitatively similar to those of C06, they refer to a different and more distant portion of the Stream, and provide an on-field demonstration that the whole Sgr Stream can be studied with this technique.  \\subsection{The data}  To obtain a representative sample of the stellar population in the core of Sgr we take the photometry of the $1\\degr \\times 1\\degr$ wide field located $\\sim 2\\degr$ eastward of the galaxy center, recently presented by B06, i.e. the Sgr34 field [(l,b)  $\\simeq(6.5\\degr,-16.5\\degr$), see B06 for details]. N02 provided two very clear detections of the Sgr Stream in their fields S341+57-22.5 and S167-54-21.5, where the first two numbers of the names are the mean Galactic longitude and latitude. While the BHB is clearly visible in the CMD of S167-54-21.5, the signal of the Stream RC is too weak in this field to derive fully reliable star counts, due to various intrinsic factors that will be discussed elsewhere (Bellazzini et al., in preparation). Hence exclude this field from the analysis and we focus on S341+57-22.5  (S341 hereafter, for brevity), a $\\sim 62.5$ deg$^2$ strip enclosed within $200\\degr \\le RA\\le 225\\degr$ and $-1.25\\degr\\le Dec\\le 1.25\\degr$, that samples a portion of the leading arm of the Stream, lying more than 40 kpc away from us in the Northern Galactic hemisphere (N02, Law et al. \\cite{law}). The g$^*$,r$^*$ photometry by N02 has been transformed into Jonhsons-Cousins B,V according to Smith et al. (\\cite{smith}). The photometry has been corrected for extinction following B06. Note that the average reddening in the S341 field is $E(B-V)=0.04$ mag. ",
        "conclusions": "We have detected a $\\simgt 4.8\\sigma$ difference in the ratio of BHB to RC stars between the core of Sgr and a distant portion of its tidal Stream. BHB stars are 5 times more abundant - relative to RC stars - in the considered Stream field than within the main body of the galaxy, a clear signature of the presence of an age/metallicity gradient along the Sgr remnant.  C06 clearly states that the observed metallicity gradient within the Stream cannot be caused by ``an intrinsic variation of the instantaneous mean metallicity of Sgr with time'', because the chemical enrichment of Sgr was essentially completed much before than the Stream stars that display the gradient effect were stripped from the main body of their parent galaxy (see also B06). This leaves the presence of a radial metallicity/age gradient in the progenitor as the only viable explanation for the observed core/Stream gradient (note, however, that the simultaneous presence of stars stripped from Sgr in different peri-Galactic passages in the same portion of the Stream may also play a role; see Fellhauer et al. 2006). C06 notes also that the implied radial gradient should have been significantly stronger than that typically encountered in dwarf spheroidal galaxies. A comparison with the results by Harbeck et al. (\\cite{harb}) shows that in the galaxies of their sample the $N_{BHB}/N_{RC}$ ratio typically varies by a factor of $<3$, over the considered radial range, that is less than the factor of 5 we observe between SGR34 and S341. However changes of a factor of $\\simgt 8$ are observed in Sculptor and Sextans, hence the value found here does not appear exceptional, although the exclusion of the hottest stars from our BHB selection may hide a larger difference.  According to the N-body models by Law et al. (\\cite{law}) the stars populating the branch of the Stream sampled by the S341 field have been stripped during different orbital revolutions, in the past. Independently of the adopted model (adopting a {\\em spherical}, {\\em oblate} or {\\em prolate} Milky Way halo), only $\\sim 30$\\% of the stars in this part of the Stream have been lost from the main body during the current peri-Galactic passage or during the immediately previous one: hence the observed population should be dominated by stars lost more than $\\simeq 1.5-2$ Gyr ago (see also Belokurov et al. \\cite{belok} and Fellhauer et al. \\cite{fel}). Thus the observed HB morphology should reflect the stellar mix in the outer regions of the Sgr progenitor as it was 2 to 4 peri-Galactic passages ago. We conclude that such a stellar mix was significantly more metal-poor, on average, than that observed in the present-day bound remnant."
    },
    "0608/astro-ph0608349_arXiv.txt": {
        "abstract": "A variety of on-orbit imaging and spectroscopic observations are used to characterize the Charge Transfer Efficiency (CTE) of the Charge-Coupled Device (CCD) of the Space Telescope Imaging Spectrograph (STIS) aboard the {\\it Hubble Space Telescope}. A set of formulae is presented to correct observations of point sources for CTE-related loss of signal. For data taken in imaging mode, the CTE loss is parametrized in terms of the location of the source on the CCD, the source signal level within the measurement aperture, the background level, and the time of observation. For spectroscopic data, it is found that one additional parameter is needed to provide an adequate calibration of the CTE loss, namely the signal in the point spread function located between the signal extraction box and the read-out amplifier. The effect of the latter parameter is significant for spectra taken using the G750L or G750M gratings of STIS.  The algorithms presented here correct flux calibration inaccuracies due to CTE losses as large as 30\\% to within $\\sim$\\,1.5\\% RMS throughout the wavelength range covered by the STIS CCD modes. This uncertainty is similar to the Poisson noise associated with a source detected at a signal level of about 2500 electrons per resolution element.  Using bi-directional CCD readouts, centroid shifts incurred due to CTE loss are also derived. A tight correlation is found between the CTE loss and the centroid shift (both for imaging and spectroscopic modes), thus enabling one to correct for both effects of imperfect charge transfer to STIS CCD observations. ",
        "introduction": "Astronomical observation was revolutionized about 25 years ago by charge-coupled device (CCD) technology, due to a combination of generally linear response over a very large dynamic range and high quantum efficiency. However, a shortcoming of CCDs is the fact that the transfer of charge from one pixel to the next is not perfect. Charge Transfer Efficiency (CTE) is the term commonly used to describe this effect, and it is quantified by the fraction of charge successfully moved (clocked) between adjacent pixels. In practice it is often more useful to use the term Charge Transfer {\\it Inefficiency\\/} (CTI $\\equiv$ 1$-$CTE). We will do so in this paper when considering quantitative measures. The main observational effect of CTE loss is that a star whose induced charge has to traverse many pixels before being read out appears to be fainter than the same star observed near the read-out amplifier.  Laboratory tests have shown that the CTE loss of CCDs increases significantly when being subjected to radiation damage \\citep[e.g.,][]{jane+91}. This is particularly relevant for space-borne CCDs such as those aboard {\\it Hubble Space Telescope (HST)}, where the flux of high-energy particles, particularly in the South Atlantic Anomaly, is significantly higher than on the ground. The effect is significant for all CCD detectors used on {\\it HST\\/} instruments (e.g., WFPC2: Whitmore, Heyer, \\& Casertano 1999; Dolphin 2000; STIS: Goudfrooij \\& Kimble 2003; ACS: Riess \\& Mack 2004), and charge losses increase with elapsed on-orbit time. The purpose of the current paper is to characterize the CTE loss of the CCD of the Space Telescope Imaging Spectrograph (STIS) for observations of point sources in terms of its dependencies on the X and Y positions, target intensity, background counts, measurement aperture size, observing mode, and elapsed on-orbit time. Several aspects of on-orbit characterizations of the CTE loss of the STIS CCD have been reported by \\citet*{gill+99}, \\citet*{kimb+00}, \\citet*{goukim03}, and \\citet*{bohgou03}. The current paper provides a comprehensive description of the final STIS CTE calibrations, using data taken through mid-2004 (i.e., just before STIS operations ended) which allows a more accurate temporal dependence of the CTE loss than the reports mentioned above. The current paper also includes analysis of data taken specifically to constrain the CTE loss at lower signal levels than before. Furthermore, we provide (for the first time) {\\it (i)\\/} separate characterizations of the CTE loss for spectroscopic and imaging modes, {\\it (ii)\\/} a comparison between data taken in gain = 1 and gain = 4 e$^-$ DN$^{-1}$ settings, and {\\it (iii)\\/} the relation of the CTE loss with its associated centroid shifts.  The STIS CCD is a 1024 $\\times$ 1024 pixel, backside-illuminated device with 21\\,$\\mu$m $\\times$ 21\\,$\\mu$m pixels. The CCD was built by Scientific Imaging Technology (SITe) with a coating process that allows coverage of the 200\\,--\\,1000 nm wavelength range for STIS in a wide variety of imaging and spectroscopic modes. Key features of the STIS CCD architecture are shown schematically in Figure~\\ref{f:ccdarch}. Two serial registers are available. A read-out amplifier is located at all four corners, each with an independent analog signal processing chain. The full image can be read out through any one of the four amplifiers, or through two-- and four-amplifier combinations. By default, science exposures employ full-frame readout through amplifier `D', which features the lowest read-out noise (4.0 e$^-$ RMS at launch).  Further technical details regarding the STIS CCD in particular are provided in \\citet{kimb+94}, while background information on the design of STIS in general can be found in \\citet{wood+98}.  This paper is organized as follows. Section~\\ref{s:moni} describes methods used to measure the CTE loss:\\ Two methods designed to quantify different aspects of CTE loss for observations of point sources in sparse fields for spectroscopic and imaging modes, and one method to improve the characterization of the CTE loss for point source spectroscopy, using on-orbit observations of spectrophotometric standard stars. Section~\\ref{s:results} presents the results of the CTE analyses, and Section~\\ref{s:concl} summarizes the results. ",
        "conclusions": "\\label{s:concl}  Using a variety of methods, a set of formulae to correct observations of point sources with the STIS CCD for CTE-related loss of signal has been derived and presented. For data taken in imaging mode, the CTE loss is parametrized in terms of the location of the source on the CCD, the source signal level within the measurement aperture, the background level, and the time of observation. For spectroscopic data, one extra parameter is needed to provide an adequate calibration of the CTE loss, namely the signal in the point spread function located between the signal extraction box and the read-out amplifier. The effect of the latter parameter is significant at wavelengths redward of 8000 \\AA\\ where the PSF of the STIS CCD exhibits a extended ``halo'', i.e., for spectra taken using the G750L or G750M gratings of STIS. The algorithms presented here correct flux calibration inaccuracies due to CTE losses as large as 30\\% to within $\\sim$\\,1.5\\% RMS throughout the wavelength range covered by the STIS CCD observing modes.  Using bi-directional CCD readouts, centroid shifts incurred due to CTE loss are also derived. Centroid shifts are up to $\\sim$\\,10$^{-1}$ pixel for imaging photometry of faint point sources, and a factor of $\\sim$\\,3 higher for spectroscopy mode. A strong correlation is found between the signal loss and the centroid shift (both for imaging and spectroscopic modes), thus enabling one to correct for both effects of imperfect CTE to STIS CCD observations. While centroid shifts due to CTE effects are rarely taken into account in the literature, they are relevant for science programs that aim to reach the highest possible positional accuracy such as proper motion measurements for which accuracy of order 0.01 pixel is often required.  Though we have demonstrated that the correction formulae work well, it is of course preferable (when practical) to minimize the magnitude of the CTE effects on an observing program at the data acquisition stage. To this end, STIS operations support acquisition of targets for CCD spectroscopic modes at the so-called `E1 pseudo-apertures', which are located $\\sim$\\,125 rows from the readout amplifier rather than at field center. The appropriate instrumental calibrations have been developed for these aperture positions. At the STIS plate scale of 0.05 arcseconds per pixel along the slit, this position gives 6 arcseconds of sky coverage to the near edge of the CCD; targets for which this angular coverage is sufficient can thus be observed with four times lower charge loss by selecting this observing position rather than the center of the CCD. Corrections for CTE loss and the use of the `E1 pseudo-apertures' will be of even greater importance in the future if STIS is successfully repaired as planned  in the hoped-for HST Servicing Mission 4, as the CCD usage will be extended for a number of additional years in the orbital radiation environment.  \\paragraph*"
    },
    "0608/astro-ph0608380_arXiv.txt": {
        "abstract": "We present evidence for a dual Active Galactic Nucleus (AGN) within an early-type galaxy at $z=0.709$ in the Extended Groth Strip. The galaxy lies on the red sequence, with absolute magnitude $M_B=-21.0$ (AB, with $h=0.7$) and rest-frame color $U-B=1.38$.  Its optical spectrum shows strong, double-peaked [\\ion{O}{3}]  emission lines and weak H$\\beta$ emission, with Seyfert-like line ratios.  The two narrow peaks are separated by $630 \\kms$ in velocity and arise from two distinct regions, spatially resolved in the DEIMOS spectrum, with a projected physical separation of $1.2$ kpc.  \\emph{HST}/ACS imaging shows an early-type (E/S0) galaxy with hints of disturbed structure, consistent with the remnant of a dissipationless merger. Multiwavelength photometric information from the AEGIS consortium confirms the identification of a dust-obscured AGN in an early-type galaxy, with detections in X-ray, optical, infrared and radio wavebands. These data are most readily explained as a single galaxy harboring two AGN---the first such system to be observed in an otherwise typical early-type galaxy. ",
        "introduction": "\\label{sec:introduction}  The standard picture of AGN emission---which identifies the engines as  accreting supermassive black holes (SMBHs)---is now well established.  Also well established is the hierarchical paradigm for structure formation, in which massive galaxies grow via a series of mergers of smaller objects.  An obvious consequence of this picture is that some galaxies should contain two SMBHs in their centers.  Such nuclear BH pairs will initially be widely separated ($\\ga 1$ kpc) ``dual'' SMBHs.  After $\\sim 100$ Myr they will become true \\emph{binary} SMBHs, gravitationally bound to one another and with parsec-scale separations; they may finally coalesce into a single central SMBH on much longer timescales~\\citep{BBR80}. The behavior of SMBH binaries has significant implications for the formation of galactic cores (e.g., \\citealt{MMRv02}), and their coalescence may be a significant source of gravitational radiation in the universe; this subject has received much theoretical study (for a review see~\\citealt{MM05}).  Observational study of these objects has been far more limited, however: only recently has direct detection of a true binary SMBH been reported, in the LINER-like elliptical radio galaxy 0402+379 \\citep{Maness04, Rodriguez06}. In addition, there is only one unambiguous detection of a dual AGN in a single galaxy, from \\emph{Chandra} studies of the merger remnant NGC 6240 \\citep{Komossa03}, although there are several examples of merging pairs of AGN hosts (\\emph{e.g.}, \\citealt{Ballo04, Guainazzi05, Hudson06}), and there is a population of close quasar pairs that are very unlikely to be lensed objects (\\emph{e.g.} \\citealt{Kochanek99, Junkkarinen01}).  Most of the merger examples were discovered in studies of merging late-type galaxies with significantly disturbed morphologies. Recent work on galaxy evolution (e.g., \\citealt{Faber05}) implies that significant stellar mass in elliptical galaxies is built up in dissipationless mergers of early-type progenitors. The remnants of such mergers exhibit only subtle visual signatures \\citep{vanD05}, suggesting that dual-SMBH systems should also occur in apparently normal early-types.  The Extended Groth Strip is a subregion of the DEEP2 Galaxy redshift survey that is also the site of intense, multiwavelength observational efforts by the All-wavelength Extended Groth strip International Survey (AEGIS) team.  Full details of the dataset can be found in \\citet{AEGIS}.  In this Letter, we present evidence from AEGIS for a dual AGN at $z=0.709$ in the $R_{AB} = 22.6$ early-type galaxy EGSD2 J142033.66+525917.5 The object was identified serendipitously in the process of visually inspecting DEEP2 spectra to confirm galaxy redshifts. Throughout this work we adopt a flat, $\\Lambda$CDM cosmology with $\\Omega_M=0.3$ and $h=0.7$; all distances are quoted in physical (not comoving) units. ",
        "conclusions": "\\label{sec:discussion}  These observations represent strong evidence for a dual AGN hosted by an early-type galaxy.  The [\\ion{O}{3}] $\\lambda 5007$/H$\\beta$ emission-line ratio of the galaxy is iron-clad evidence for an AGN; this is further supported by the X-ray detection.  The conclusion that there are in fact \\emph{two} accreting SMBHs in this galaxy is driven entirely by the DEEP2 spectroscopic data.  Double peaked emission lines like those seen in Fig.~\\ref{fig:1dspec} could be explained by three reasonable models: a rotating disk, an outflow (or inflow), or a pair of orbiting emission components. A reasonable disk model, however, cannot easily accommodate emission regions separated by $1.2$ kpc and $630 \\kms$. Typical accretion disks around SMBHs are smaller by several orders of magnitude and have relativistic velocities (\\emph{e.g.}, \\citealt{EH03}); it is difficult to understand how a larger disk could exhibit its brightest AGN-like emission at distances $\\sim 1$ kpc from the SMBH. On the other hand, outflows of gas and dust are commonly invoked to explain asymmetric [\\ion{O}{3}] emission from narrow-line AGN (\\emph{e.g.},  \\citealt{Heckman81}). These asymmetries, however, typically take the form of broad, blue wings, not strong second peaks, and in this picture we would expect one emission component to be at rest at the host galaxy's center, unlike what is observed here. There do exist, in a few nearby Seyferts, bipolar outflows that have spatial scales, bulk velocities, and emission-line ratios that could explain our observations (\\emph{e.g.}, \\citealt{Axon98, VSM01}), but these systems are quite kinematically disturbed, with emission components whose velocity dispersions are comparable to their velocity separations.  This contrasts strongly with the relatively narrow emission-line components in this object.   We thus regard a dual SMBH as the likeliest explanation for our observations.  The identification is quite consistent with theoretical expectations and with previous observations. As discussed in the Introduction, dual SMBHs are expected to remain separated at $\\ga 1$ kpc for $\\sim 100$ Myr after a galaxy merger and then rapidly form a parsec-scale binary Indeed, like this one, the previously observed dual AGN in NGC 6240 \\citep{Komossa03} has a separation $\\sim 1$ kpc.  Unlike this one, that galaxy also shows strong signatures of recent merger activity, but since such signatures are quite subtle in dissipationless mergers, they will not be as obvious in our imaging. However, the identification of a dual AGN in this galaxy is not as secure as in NGC 6240, where \\emph{Chandra} can resolve two X-ray sources in the galaxies' centers; the angular scale of interest here ($0^{\\prime\\prime}.17$) is simply too small to be resolved in our observations.    Nevertheless, assuming our interpretation is correct, it is possible to infer further information about the two SMBHs. First, they are clearly bound within the galaxy, since their peculiar velocities ($\\sim 300 \\kms$) are much smaller than the escape velocity ($\\sim 1000 \\kms$) implied by the host galaxy's measured stellar mass.  Second, our observations  permit an estimate of the SMBH masses. \\citet{GH05} show that, in nearby AGN, the width of the [\\ion{O}{3}] emission lines is a reasonable proxy for the stellar velocity dispersion of the galaxy bulge, $\\sigma_\\ast$, provided that any asymmetric wings are removed.  Our line-fitting procedure accounts for all obvious asymmetries, so we may use the measured linewidths to infer SMBH masses from the $M_{\\mathrm{BH}}$-$\\sigma_\\ast$ relation (\\emph{e.g.}, \\citealt{GH06}).   This relation implies $M_{\\mathrm{BH}}\\sim 5\\times 10^5 M_\\odot$ for the blueward component and $M_{\\mathrm{BH}}\\sim 5\\times 10^6 M_\\odot$ for the redward one.  By comparison, the relation between bulge mass and SMBH mass (\\emph{e.g.} \\citealt{HR04}) implies that this $M_\\ast= 10^{11} M_\\odot$ galaxy should harbor a BH of mass $\\sim 10^8 M_\\odot$, more than an order of magnitude larger than the total BH mass we infer.  The discrepancy is only at the $\\sim 2 \\sigma$ level, however, given the large scatter in the relations discussed here.   Simple considerations might lead one to expect that dual SMBHs reside in a relatively large fraction of galaxies: if a typical galaxy undergoes a major merger every few Gyr, and the resulting dual SMBH has a lifetime of $\\sim 100$ Myr, then a few percent of all galaxies---say $\\sim 1000$ of the $\\sim 40000$ galaxies in DEEP2---should host such systems.  The fact that we have found only one dual AGN in DEEP2 is not inconsistent with this estimate. First, it is not clear that \\emph{every} dual SMBH will exhibit AGN activity for its full lifetime, especially from both nuclei at once. Even if so, clean detection of double-peaked AGN emission lines is easiest if the bright, narrow [\\ion{O}{3}] emission lines fall in the DEEP2 wavelength range; this is true for only $\\sim 30\\%$ of the DEEP2 sample.  Also, dual AGN that occur in late-type galaxies may have the bulk of their optical emission obscured by dust (e.g., \\citealt{Zakamska05}), so we can expect objects with observational signatures like this one to occur mostly in early-type systems.  These constitute only $\\sim 15$\\% of the DEEP2 sample.    Thus we would expect DEEP2 to have only a few dozen objects like the one studied here if all such objects are active.  This dual AGN was discovered serendipitously; a comprehensive search of the DEEP2 catalog might find more.   We close by addressing the relatively weak emission bridging the two emission components in Figs.~\\ref{fig:1dspec} and~\\ref{fig:2dspec}.  As already mentioned, narrow [\\ion{O}{3}] lines in  AGN spectra frequently exhibit asymmetric blue wings, probably from outflows (\\emph{e.g.}, \\citealt{Heckman81}).  If such an outflow is present for the redward AGN here, it could explain the bridging emission. Given the unusual nature of this galaxy's spectrum, however, this explanation is by no means certain. Infrared imaging and integral-field unit observations on an 8-meter class telescope with adaptive optics, or narrow-band imaging of the [\\ion{O}{3}] line from space, could help to further characterize this emission and the system as a whole."
    },
    "0608/astro-ph0608455_arXiv.txt": {
        "abstract": "Evidence for non-thermal activity in clusters of galaxies is well established from radio observations of synchrotron emission by relativistic electrons. New windows in the Extreme Ultraviolet and Hard X-ray ranges have provided for more powerful tools for the investigation of this phenomenon. Detection of hard X-rays in the 20 to 100 keV range have been reported from several clusters of galaxies, notably from Coma and others. Based on these earlier observations we identified the relatively high redshift cluster 1E0657-56 (also known as RX J0658-5557) as a good candidate for hard X-ray observations. This cluster, also known as the bullet cluster, has many other interesting and unusual features, most notably that it is undergoing a merger, clearly visible in the X-ray images. Here we present results from a successful \\textit{RXTE} observations of this cluster. We summarize past observations and their theoretical interpretation which guided us in the selection process. We describe the new observations and present the constraints we can set on the flux and spectrum of the hard X-rays. Finally we discuss the constraints one can set on the characteristics of accelerated electrons which produce the hard X-rays and the radio radiation. ",
        "introduction": "\\label{intro}  The intra-cluster media (ICM) of several clusters of galaxies, in addition to the well studied thermal bremsstrahlung (TB) emission dominating in the $\\sim 2 - 10$ keV soft X-ray (SXR) region, show growing evidence for non-thermal activity. This activity was first observed in a form of {\\it diffuse radio} radiation (classified either as relic or halo sources). The first cluster with diffuse emission detected in the radio band was Coma, and recent systematic searches have identified more than 40 clusters with halo or relic sources.  In the case of Coma, the radio spectrum may be represented by a broken power law (Rephaeli 1979), or a power law with a rapid steepening (Thierbach \\ea 2003) or with an exponential cutoff (Schlickeiser \\ea 1987).  There is little doubt that this radiation is due to synchrotron emission by a population of relativistic electrons with similar spectra, however, from radio observations alone one cannot determine the energy of the electrons or the strength of the magnetic field.  Additional observations or assumptions are required.  Minimum total (particles plus field) energy or equipartition assumptions imply magnetic field strength of $B\\sim \\mu{\\rm G}$, in rough agreement with the Faraday rotation measurements (\\eg Kim \\ea 1990), and a population of relativistic electrons with Lorentz factor $\\g \\sim 10^4$. In the  papers cited above, it was also realized that because the  energy density of the Cosmic Microwave Background (CMB) radiation (temperature $T_0$) $u_{\\rm CMB}= 4\\times 10^{-13}(T_0/2.8 \\, {\\rm K})^{4}\\,{\\rm erg} \\,\\cc$ is larger than the  magnetic energy density $u_{\\rm B}= 3\\times 10^{-14}(B/\\mu{\\rm G})^2 \\, {\\rm erg} \\,\\cc$, most of the energy of the relativistic electrons will be radiated via inverse Compton (IC) scattering of the CMB photons, producing a broad photon spectrum (similar to that observed in the radio band) around 50 keV (for $\\g \\sim 10^4$). Thus, one expects a higher flux of non-thermal X-ray radiation than radio radiation. Detection of HXRs would then break the degeneracy and allow determination of the magnetic field and the energy of the radiating electrons. Moreover, since the redshift dependence of the CMB photons is known, then in principle, the cosmological evolution of these quantities can also be investigated.  While the detectability of such radiation would be easiest in the soft X-ray range, where very sensitive imaging instruments are available, in reality, in this range, it is masked by the prominent thermal bremsstrahlung emission, a general characteristic of clusters. With this, the most promising is either the hard X-ray (HXR) band, beyond the energies where TB flux dominates, or alternatively, in the very soft X-ray or extreme ultraviolet regime. Recently  HXR emission (in the 20 to 80 keV range) at significant levels above that expected from the thermal gas was detected by instruments on board {\\it Beppo}SAX and \\textit{ RXTE } satellites from Coma (Rephaeli \\ea 1999, Fusco-Femiano \\ea 1999, Rephaeli \\& Gruber 2002, Fusco-Femiano \\ea 2004% \\footnote{The results of this paper have been challenged by an analysis performed with different software by  Rossetti \\& Molendi (2004).}), Abell 2319 (Gruber \\& Rephaeli 2002), Abell 2256 (Fusco-Femiano \\ea 2000, Rephaeli \\& Gruber 2003, and Fusco-Femiano, Landi, \\& Orlandini 2005), and  a marginal ($\\sim3\\sigma$) detection from Abell 754 and an upper limit on Abell 119 (Fusco-Femiano \\ea 2003). We also note that a possible recent detection of non-thermal X-rays, albeit at lower energies, has been reported from a \\underline{poor} cluster IC 1262 by Hudson et al. (2003).  All those clusters are in the redshift range $0.023<z<0.056$. Notable recent exception at a higher redshift is Abell 2163 (Rephaeli, Gruber, \\& Arieli 2006) where the reported nonthermal flux is consistent with the upper limit set by {\\it Beppo}SAX (Feretti \\ea 2001). It should also be noted that excess radiation was detected in the 0.1 to 0.4 keV band by {\\it Rosat, Beppo}SAX and {\\it XMM-Newton} and in the EUV region (0.07 to 0.14 keV) similar excess radiation was detected by the {\\it Extreme Ultraviolet Explorer} from Coma (Lieu et al.  1996) and some other clusters. A cooler ($kT\\sim 2$ keV) component and IC scattering of CMB photons by lower energy ($\\g \\sim 10^3$) electrons are two possible ways of producing this excess radiation. However, some of the observations and the emission process are still controversial (see Bowyer 2003).  Here we present results from $\\sim 309$ ks \\textit{RXTE } observations of cluster 1E0657-56 (also known as RX J0658-5557) with a considerably higher redshift of $z=0.296$, which manifests many other interesting features (Markevitch 2005).  All those hard X-ray observations - especially with the addition of 1E0657-56 - encompass a wide range of temperature, redshift and luminosity, indicating that HXR emission may be a common property of all clusters with significant diffuse radio emission. In the next section we give brief descriptions of the emission processes and the considerations which led to the selection of 1E0657-56 as a target for HXR observations by \\textit{RXTE }. In \\S 3 we describe the observations and the results from our spectral fits. Finally in \\S 4 we discuss the significance of our results and their implication for the radiation and acceleration mechanisms. ",
        "conclusions": "\\label{disc}  Radio observations of diffuse emission from ICM show presence of non-thermal activity in many clusters, especially in those with high SXR emitting gas temperatures (and luminosity) and showing recent merger activity.  A list of such clusters is shown in Table 1. Detections of HXR emission exceeding the levels expected from the thermal gas have been reported by two groups using different instruments. This enables a quantitative investigation of the nature and origin of the non-thermal activity. The radio observations alone indicate presence of extreme relativistic electrons ($\\g\\sim 10^4$). Assuming equipartition between the relativistic electrons and the magnetic field we estimate a volume-averaged magnetic field value in the range of 0.5 to 2 $\\mug$ (see Table 1), in rough agreement with (line-of-sight averaged) field strengths deduced from Faraday rotation measurements. In the case of Coma the magnetic field deduced from Faraday rotation of $2-3 \\, \\mug$, implies conditions far from equipartition if we assume a homogeneous source;  the field energy density is about 25 times larger than that of the electrons.  This discrepancy can be resolved by an inhomogeneous model; \\eg the two phase model proposed by Brunetti \\ea (2001) that implies a magnetic field profile, or if  there was 25 times more energy in nonthermal protons than electrons.  If the HXR radiation is due to IC scattering of CMB photons by the the same electrons that produce the synchrotron radio emission, which we have argued to be the simplest scenario, such high magnetic files would imply IC fluxes in the $20-100$ keV range of about $10^{-2}F_0$ while both Coma and A2256 show fluxes higher than $F_0$. Production of HXRs at this level require magnetic fields of about $0.1 \\, \\mug$ (or partition parameter $\\zeta\\sim 10^3$) which raises the required energy of the non-thermal electrons (${\\cal E}_e\\propto \\zeta^{1/2}$). For example for Coma and A2256 the observed HXR fluxes require ${\\cal E}_e= 2\\times 10^{-13}$ and $1.5\\times 10^{-12}$ erg $\\cc$, respectively, which are comparable to the energy density of the CMB ($4.6\\times 10^{-13}$ erg $\\cc$). Assuming similar conditions for other clusters, several years ago we identified 1E0657-56 and A2219 as possible candidate for having detectable HXR fluxes. Here, we have described these observations of 1E0657-56 and the results of our spectral fitting to these data alone and to the \\textit{RXTE } data combined with the SXR data from \\textit{XMM-Newton}. In the future, attempts should be made to observe other clusters in Table 1, and any newly discovered cluster with a relatively high temperature and redshift, an ample sign of substructure and/or mergers, and, of course, with a strong diffuse radio emission. A low value of magnetic field indicated by Faraday rotation measurements or estimated via equipartition argument will also be good justification for HXR observations.  We have shown that a thermal component plus a non-thermal tail described as a power law with spectral index $\\Gamma=1.6$ containing $\\sim 0.5 F_0$ of the $20-100$ keV flux provides an acceptable fit. This flux is fortuitously close to the value of $0.52 \\, F_0$ we estimate by scaling the Coma and A2256 observations. (Note that the equipartition (\\ie $\\zeta=1$) magnetic field and flux for this cluster would be $\\sim 1 \\, \\mug$ and $0.02 \\, F_0$.) This agreement is somewhat encouraging so that we believe that the chances  are even better for  detection of HXR emission from Abell 2219, a good candidate  at a moderate redshift,  and a recently discovered cluster, MACJ7017, with even higher $z$ than 1E0657-56. As is the case for Coma and Abell 2256, these clusters also require high nonthermal electron energy densities; ${\\cal E}_e= [1.0, 1.8, 8.5]\\times 10^{-13}$ erg $\\cc$ for Abell 2219, 1E0657-56 and MACSJ0717, respectively. These are even closer to the CMB densities at their respectively higher redshifts. In general the required electron energies are about 1 to 10 \\%  of the energy density of the thermal gas of few times $10^{-11}$ erg $\\cc$ which is also comparable to the gravitational potential energy density of the cluster. This means a very efficient acceleration process.  Regarding the spectral fits to 1E0657-56, as shown in Table 2, we also obtain an acceptable fit with a double temperature thermal model with a second component with $kT=50$ keV and a normalization implying a volume emission measure ($EM$) for this component which is a significant fraction ($\\sim 10\\%$) of that of the lower temperature component. Indeed Markevitch (2005) finds evidence for such a high $T$ component behind the shock near the bullet based on the recent Chandra Observatory data.  However, this has a much smaller $EM$ than is required for fitting the data with a two temperature model. In fact a high $T$ component with such a large $EM$ would have been easily detected by Markevitch (2005). Clearly more and deeper observations of the clusters in Table 1 can clarify this situation considerably. In particular, deep observations with imaging HXR instruments -- such as those with the proposed {\\sl NuSTAR} satellite -- which can  provide information about the spatial distribution of the HXRs  would be very valuable.  In spite of the uncertainties about the exact character  of the observed radiation spectra and the emission mechanism, and in spite of the meagerness of the data, the acceleration mechanism of electrons can  be constrained significantly.  The lifetimes of electrons with energies in the range $200 \\, {\\rm keV} \\leq E \\leq 200 \\, {\\rm GeV}$ are longer than the crossing time, $T_\\tr\\sim 3\\times 10^6$ yr.  Therefore, these electrons will escape the cluster and radiate most of their energy outside of the cluster unless there exists some scattering agent with a mean free path $\\lambda_\\scat\\sim 1$ kpc to trap the electrons in the ICM for at least a timescale of $T_\\esc=(R/\\lambda_\\scat)T_\\tr\\sim 3\\times 10^9$ yr, for cluster size $R\\sim 1$Mpc. Turbulence can be this agent and can play a role in stochastic acceleration directly, or indirectly in acceleration by shocks, presumably arising from merger events. Several lines of argument point to an ICM which is highly turbulent, and there has been considerable discussion of these aspects of the problem in the recent literature (see, e.g., Cassano \\& Brunetti 2005). These aspect are explored in P01, but the upshot of which is that {\\it we require injection of high energy electrons, presumably from galaxies and AGNs, and that the acceleration process is episodic on time scale of $\\sim 10^8$ to $10^9$ yr.} Whether the electrons are injected directly by past and current AGN activity, or are produced by the interaction with the thermal gas of cosmic ray protons escaping the galaxies would be difficult to determine now. There are however constrains on both these scenarios (see \\eg Blasi 2003)."
    },
    "0608/astro-ph0608488_arXiv.txt": {
        "abstract": "We analyze thermonuclear and pycnonuclear fusion reactions in dense matter containing atomic nuclei of different types. We extend a phenomenological expression for the reaction rate, proposed recently by Gasques {\\it et al}.\\ \\cite{leandro05} for the one-component plasma of nuclei, to the multi-component plasma. The expression contains several fit parameters which we adjust to reproduce the best microscopic calculations available in the literature. Furthermore, we show that pycnonuclear burning is drastically affected by an (unknown) structure of the multi-component matter (a regular lattice, a uniform mix, etc.). We apply the results to study nuclear burning in a $^{12}{\\rm C}^{16}{\\rm O}$ mixture. In this context we present new calculations of the astrophysical $S$-factors for carbon-oxygen and oxygen-oxygen fusion reactions. We show that the presence of a CO lattice can strongly suppress carbon ignition in white dwarf cores and neutron star crusts at densities $\\rho \\gtrsim 3 \\times 10^9$ g~cm$^{-3}$ and temperatures $T \\lesssim 10^8$~K. ",
        "introduction": "\\label{introduction}  Nuclear reactions are most important for the physics of stars. They determine hydrogen burning in main-sequence stars, helium burning in red giants, and carbon, neon, and oxygen burning at later stages. They determine also nucleosynthesis in shock driven stellar explosions, such as type II supernovae, as well as ignition and burning in accreting stars.  Steady-state and explosive thermonuclear carbon burning during late stellar evolution \\cite{eleid-2004} and in shock fronts of type II supernovae \\cite{pardo-couch-arnett-74} is governed by the $^{12}$C+$^{12}$C and possibly by the $^{12}$C+$^{16}$O fusion processes. Similarly, thermonuclear oxygen burning is mainly determined by the $^{16}$O+$^{16}$O and possibly by the $^{16}$O+$^{20}$Ne  reactions \\cite{truran-arnett-1970}. The ignition and nucleosynthesis during these burning phases critically depend on the initial fuel abundance and on the thermonuclear reaction rates.  In high-density cores of white dwarfs and crusts of neutron stars, the thermonuclear reactions are strongly affected by plasma effects (especially important for carbon ignition in cores of accreting massive white dwarfs for triggering type Ia supernova explosions). The ignition conditions are sensitive to the $^{12}$C and $^{16}$O abundance and to the fusion reaction rates \\cite{niemeyer-woosley-1997}.  Carbon ignition has also been suggested as a trigger of superbursts in surface layers of accreting neutron stars \\cite{strohmeyer-bildsten}. However, the required ignition conditions seem to disagree with the observed superburst lightcurves \\cite{cumming-2005}. Alternative explanations are presently being discussed, such as carbon ignition in the crust of an accreting strange (quark) star, to accommodate the observed light curve characteristics \\cite{page-cumming-2005}. While in most of the scenarios pure carbon burning dominates the energy production, ignition conditions and associated nucleosynthesis are affected by the presence of other elements.  Pycnonuclear burning occurs in dense and cold cores of white dwarfs \\cite{svh69} and in crusts of accreting neutron stars \\cite{Schramm, hz}. Theoretical formalism has been mostly developed for pycnonuclear reactions between equal nuclei, but one often needs to consider a multi-component matter, for instance, carbon-oxygen cores of white dwarfs.  In a previous publication \\cite{leandro05} we have focused on fusion reactions between equal nuclei in a one-component plasma (OCP) of atomic nuclei (ions). In the present work we expand the study towards a multi-component plasma (MCP). The problem  has two aspects; the first one is associated with the underlying nuclear physics while the second one is concerned with the plasma physics. The nuclear part deals with the reliable determination of astrophysical $S$-factors at stellar energies. These energies are low (typically, lower than a few MeV), in particular if compared to the presently accessible range of low-energy fusion experiments. This prevents direct measurements of $S$-factors at laboratory conditions. Thus, one needs to calculate the $S$-factors theoretically and use these results to extrapolate measured $S$-factors towards lower stellar energy range. In Section \\ref{nuclear} we present calculations of the $S$-factors for the two reactions of astrophysical importance, $^{12}$C+$^{16}$O and $^{16}$O+$^{16}$O.  The plasma physics problem consists in calculating the Coulomb barrier penetration in nuclear reactions taken into account Coulomb fields of surrounding plasma particles. These fields modify the reaction rates and lead to five nuclear burning regimes \\cite{svh69} (two thermonuclear regimes, with weak and strong plasma screening; two pycnonuclear regimes, for zero-temperature and thermally enhanced burning; and the intermediate regime). These regimes are described in Section \\ref{fusion}; their validity conditions are specified in Section \\ref{phys}. In ordinary stars, nuclear burning proceeds in the weak screening thermonuclear regime \\cite{FCZ,clayton83}. The foundation of the theory of thermonuclear burning with strong plasma screening was laid by Salpeter \\cite{Salp}. The strict theory of pycnonuclear burning was developed by Salpeter and Van Horn \\cite{svh69}. References to other works can be found in Section \\ref{fusion} and in \\cite{leandro05}.  In Section~\\ref{fusion} we analyze  calculations of Coulomb barrier penetration in MCP and propose a phenomenological expression for a reaction rate valid in all five regimes for any non-resonant fusion reaction. This expression accurately reproduces well known results in thermonuclear regimes and gives a reasonable description of the Coulomb tunneling problem in other regimes. It is important for incorporating the plasma physics effects into computer codes which simulate nucleosynthesis, especially at high densities in compact stars, such as white dwarfs and neutron stars (see above). In Section \\ref{ignition} we illustrate the results of Sections \\ref{nuclear} and \\ref{fusion} by analyzing nuclear burning in $^{12}$C\\,$^{16}$O mixtures, with the emphasis on the carbon ignition curve. ",
        "conclusions": "\\label{conclusions}  We have studied the problem of Coulomb barrier penetration for non-resonant nuclear fusion reactions in a dense MCP of atomic nuclei. We have considered all five nuclear burning regimes (Sections \\ref{thermo}--\\ref{intermed}) and analyzed calculations of nuclear reaction rates in an MCP for these regimes, available in the literature. We have proposed (Section \\ref{together}) a unified phenomenological expression for the reaction rate valid for all regimes. It generalizes an analogous expression proposed recently \\cite{leandro05} for an OCP. The expression contains several parameters which can be varied to account for current theoretical uncertainties of the reaction rates.  Our main conclusions are: \\begin{enumerate}  \\item The reaction rates in the thermonuclear regimes (with weak and strong plasma screening) can be calculated sufficiently accurately. In the regime of strong screening and for reacting nuclei with non-equal charges, our expression is more accurate than those proposed in the literature (Section \\ref{thermoscreen}).  \\item The reaction rates in other regimes (zero-temperature and thermally enhanced pycnonuclear regimes; intermediate thermo-pycno nuclear regime) are much less certain. They are very sensitive to currently unknown microphysical correlation properties in an MCP (a uniform mix, a regular crystalline lattice, a phase separated matter, a matter with impurities and defects); they are much richer in physics than in the OCP case.  \\item At low temperatures, we have mainly considered reactions in a uniform mix. Other MCP microstructures can strongly decrease or increase the reaction rates. For instance, the reactions in a regular MCP lattice can be strongly suppressed by the absence of nearby reacting nuclei.  \\item Our phenomenological expression can be improved (Section \\ref{together}) after new calculations of the reaction rates or main properties of the MCP are performed. It would be important to know the actual microstructure of the MCP at low temperatures (first of all, the availability of closest neighbors, local separations and oscillation frequencies of neighboring nuclei, particularly in the presence of impurities and lattice defects).  \\item Although our main formula in Section \\ref{together} assumes a uniform mix at low $T$, the presented results are sufficient to understand qualitatively the reaction rates for other cases (following prescriptions of Section \\ref{pycnozero}).   \\end{enumerate}  For illustration, we have considered (Section \\ref{ignition}) C+C, C+O, and O+O nuclear reactions in a dense carbon-oxygen mixture, that is important for the structure and evolution of massive white dwarfs (supernova Ia explosions) and accreting neutron stars (as sources of superbursts). For this purpose we have calculated and parameterized the appropriate astrophysical factors (Section \\ref{nuclear}). The main results of our analysis are as follows:  \\begin{enumerate} \\item The ranges of densities and temperatures, where C+C, C+O, and O+O reactions are most important, look like narrow regions in the temperature-density diagram (Fig.~\\ref{fig:diag}); the regions do not strongly overlap which means strong heterogeneity of these reactions.  \\item  With increasing density and/or temperature, carbon starts burning first in the C+C reaction (because carbon nuclei have lower Coulomb barrier); this reaction is most important for the nuclear evolution of CO mixtures.  \\item  Carbon burning in the C+C reaction is affected by the presence of oxygen. The effect is simple in the thermonuclear regimes but more complicated in other regimes (at low temperatures).  \\item  Carbon ignition in a CO mixture occurs (Fig.\\ \\ref{fig:igni}) in thermonuclear regimes as long as $\\rho \\lesssim 10^9$ g~cm$^{-3}$ (and $T \\gtrsim 3 \\times 10^8$~K). It can be calculated quite accurately. With decreasing the fraction of carbon, the ignition curve shifts to higher $\\rho$ and $T$.  \\item At $\\rho \\gtrsim 10^9$ g~cm$^{-3}$ and $T \\lesssim 3 \\times 10^8$~K the ignition condition becomes uncertain (Section \\ref{ignition}). The formation of a regular CO lattice after the crystallization can block the C+C reaction and shift the carbon ignition to higher densities.  \\end{enumerate}   Our consideration in this paper was general. More quantitative nuclear network simulations involving thermonuclear and pycnonuclear burning in dense stellar matter are currently in progress."
    },
    "0608/astro-ph0608441_arXiv.txt": {
        "abstract": "Long duration Gamma-Ray Bursts (GRBs) have been strongly connected with core collapse supernovae, so it was surprising when the recent GRB060614 (with a reported redshift of 0.125) was found to have no visible supernova to deep limits.  Three separate groups have reached the same conclusion that this event forces the existence of a new previously-unsuspected class of GRBs.  The problem with this conclusion is that the redshift is not secure, as the measured galaxy emission lines simply give the redshift of the brightest galaxy near the line-of-sight to the GRB afterglow.  Fortunately, eight different luminosity indicators are known which can give the luminosity (and hence redshift) of the GRB itself.  In combination, these luminosity indicators have already proven to give redshifts with average one-sigma errors of 26\\% as based on 69 GRBs with spectroscopic redshifts.  For GRB060614, the luminosity indicators all uniformly give a high luminosity for the burst which implies a high redshift.  This can be seen by the spiky many-peaked light curve, the high $E_{peak}$ value, the near-zero spectral lag, and the timing of the jet break.  Quantitatively, we derive a redshift of $1.97_{-0.53}^{+0.84}$ from the eight luminosity indicators.  Additional information from the detection of the afterglow in all {\\it Swift UVOT} bands further constrains the redshift to be $<1.71$, so the final range of allowed redshift is $1.44<z<1.71$, which readily explains why no supernova was visible.  Thus, GRB060614 is at high redshift and is certainly not evidence for a new class of GRBs. ",
        "introduction": "Long-duration Gamma-Ray Bursts (GRBs) are now known to be associated with Type Ib/c supernovae.  The primary evidence for this is four identifications of spectroscopically-confirmed supernovae with core collapse at the same time and place as the GRBs (Galama et al. 1998; Hjorth et al. 2003; Stanek, K. Z. et al. 2003; Malesani, D. et al. 2004; Campana et al. 2006).  Along with secondary evidence (e.g., Schaefer et al. 2003) and strong theoretical support (e.g., MacFadyen, Woosley, \\& Heger 2001; Woosley \\& Bloom 2006), the GRB community is now confident that GRBs arise promptly from the core collapse of supernovae.  A prediction of this new paradigm is that all nearby GRBs should have visible supernovae.  After GRB060614 exploded, a redshift of 0.125 was determined (Price et al. 2006) and this event became a test of the paradigm prediction.  Three groups (Della Valle et al. 2006; Gal-Yam et al. 2006; Fynbo et al. 2006) tested this prediction and discovered that no supernova was visible to very strict limits.  This contradiction to the current paradigm forced all three groups to conclude that there must be a separate eruption mechanism different from the standard model.  Indeed, the titles of their three papers declare that \"GRB060614 requires a novel explosion process\", that GRB060614 is \"not due to a hypernova\", and that it represents \"a new type of massive stellar death\".  The weak point in the case for a new class of GRBs is the distance to GRB060614.  The reported wavelengths of the observed galaxy emission lines only gives the redshift of the brightest galaxy near the line of sight to the transient.  What is needed is a redshift for the GRB.  And exactly this can be provided with all needed accuracy by any of eight luminosity indicators.  The purpose of this paper is to determine the redshift to GRB060614 and test whether the lack of a visible supernova is simply due to the burst really being at high redshift. ",
        "conclusions": "Three papers have just appeared all claiming that GRB060614 is at z=0.125 and that no associated supernova is detected to severe limits, thus forcing a conclusion that \"GRB060614 requires a novel explosive process\".  The weak point in these claims is that the reported redshift is {\\it not} that of the GRB but merely of the brightest galaxy near the line of sight.  The simplest idea is that the galaxy is just a random foreground source unrelated to the burst and that the burst is at a high redshift so as to easily account for the lack of any visible supernova.  Indeed, a detailed calculation shows that roughly one-to-two such false alarms are {\\it expected} to date.  Occam's Razor argues strongly against the \"new type of massive stellar death\" proposal as being an unnecessary new phenomenon that is already well explained by ordinary chance alignment.  But we do have eight ways of getting a redshift to the GRB itself, and that is to use the eight luminosity indicators to derive luminosities and distances.  These tools are already known to work for 69 GRBs with spectroscopic redshifts (see Figure 1), for which the average error in redshift is 26\\%.  This confirmed accuracy is more than adequate to determine whether GRB060614 is at z=0.125 or at high redshift where any underlying supernova would be invisible.  We see that the burst light curve is very spiky with many peaks, the spectral lag is near zero, and the $E_{peak}$ value is high; all pointing to a high luminosity event and a high redshift.  Detailed analysis gives a consistent picture from all eight luminosity relations that GRB060614 is really at $z=1.97_{-0.53}^{+0.84}$ and certainly not at $z=0.125$.  This provides a simple and forced explanation as to why no supernova event was visible.  As such, the claim to use GRB060614 as evidence for a \"new type of massive stellar death\" is certainly wrong.  $~$  We thank David Palmer for information on GRB060505."
    },
    "0608/astro-ph0608394_arXiv.txt": {
        "abstract": "{ Optical spectroscopic identification of the nature of 21 unidentified southern hard X--ray objects is reported here in the framework of our campaign aimed at determining the nature of newly-discovered and/or unidentified sources detected by {\\it INTEGRAL}. Our results show that 5 of these objects are magnetic Cataclysmic Variables (CVs), 4 are High-Mass X--ray Binaries (HMXBs; one of which is in the Large Magellanic Cloud) and 12 are Active Galactic Nuclei (AGNs). When feasible, the main physical parameters for these hard X--ray sources are also computed using the multiwavelength information available in the literature. These identifications further underscore the importance of {\\it INTEGRAL} in the study of the hard X--ray spectrum of AGNs, HMXBs and CVs, and the usefulness of a strategy of catalogues cross-correlation plus optical spectroscopy to securely pinpoint the actual nature of the X--ray sources detected with {\\it INTEGRAL}. ",
        "introduction": "\\begin{figure*}% \\vspace{-2cm} \\hspace{-1.7cm} \\psfig{file=figure1.ps,width=21cm}%  \\vspace{-13.7cm} \\parbox{6.0cm}{ \\hspace{6.0cm} } \\parbox{11.7cm}{ \\caption{DSS-II-Red optical images of the fields of IGR J07597$-$3842, (upper left panel), IGR J10101$-$5654 (upper middle panel), XSS J12270$-$4859 (upper right panel), IGR J14175$-$4641 (central left panel), IGR J14471$-$6319 (central middle panel), IGR J14515$-$5542 (central right panel) and IGR J14536$-$5522 (lower left panel). When present, circles indicate the {\\it ROSAT} or {\\it Swift} error boxes of the soft X--ray counterparts of {\\it INTEGRAL} objects. The putative optical counterparts of the {\\it INTEGRAL} sources are always indicated with tick marks. Field sizes are 2$\\farcm$5$\\times$2$\\farcm$5. In all cases, North is up and East to the left.} \\vspace{1cm} } \\end{figure*}   \\begin{figure*}% \\vspace{-2cm} \\hspace{-1.7cm} \\psfig{file=figure2.ps,width=21cm}%  \\vspace{-13.5cm} \\parbox{6.0cm}{ \\hspace{6cm} } \\parbox{11.7cm}{ \\caption{As Fig. 1, but for the fields of IGR J14552$-$5133, (upper left panel), IGR J15094$-$6649 (upper middle panel), IGR J16185$-$5928 (upper right panel), IGR J16558$-$5203 (central left panel), IGR J17200$-$3116 (central middle panel), IGR J17488$-$3253 (central right panel) and IGR J17513$-$2011 (lower left panel). For the sake of clarity, the {\\it Chandra} error circles of the latter two sources were not reported in the corresponding panels.} \\vspace{1.5cm} } \\end{figure*}   Since its launch in October 2002, the {\\it INTEGRAL} satellite (Winkler et al. 2003) is boosting our knowledge of the hard X--ray sky above 20 keV in terms of both sensitivity and positional accuracy of the detected sources. Thanks to the capabilities of the IBIS instrument (Ubertini et al. 2003), {\\it INTEGRAL} is effectively detecting hard X--ray objects at the mCrab level with a typical localization accuracy of 2--3$'$ (Gros et al. 2003). This has made it possible, for the first time, to obtain all-sky maps in the 20--100 keV range with arcminute accuracy and down to mCrab sensitivities (e.g., Bird et al. 2006).  Most of the sources detected by {\\it INTEGRAL} are known Galactic X--ray binaries ($\\sim$50\\% of the total number of detected objects), plus a fraction of known Active Galactic Nuclei (AGNs; $\\sim$10\\%) and Cataclysmic Variables (CVs; $\\sim$5\\%). However, a large majority of the remaining objects (about 25\\% of all detections achieved with IBIS) has no obvious counterpart at other wavelengths and therefore cannot immediately be associated with any known class of high-energy emitting objects.  Recently, in order to fill this identification gap, we started a campaign aimed at identifying the nature of these still unknown sources through optical spectroscopy at northern and southern telescopes (Masetti et al. 2004, 2006a,b,c; hereafter Papers I-IV). Our results indicate that, despite {\\it INTEGRAL} doubled the number of Galactic High Mass X--ray Binaries (HMXBs; see Walter et al. 2006) and despite the expectation according to which most of these unidentified objects should be HMXBs (Dean et al. 2005), about half of them are actually identified in the optical as nearby ($z \\la$ 0.1) AGNs (Papers I-IV).  In the framework of our continuing effort to identify unknown {\\it INTEGRAL} sources, we present here the optical spectroscopic observations obtained on 21 southern objects at the 1.5-metre telescope of the Cerro Tololo Interamerican Observatory (CTIO) located in Cerro Tololo (Chile). In Sect. 2 we introduce the criteria with which the sample of {\\it INTEGRAL} and optical objects was chosen for the present observational campaign, whereas in Sect. 3 a description of the observations is given; Sect. 4 reports and discusses the results, divided into three broad classes of sources (CVs, HMXBs and AGNs), together with a statistical outline of the identifications of {\\it INTEGRAL} sources obtained up to now. Conclusions are drawn in Sect. 5. ",
        "conclusions": "In our continuing work of identification of {\\it INTEGRAL} sources by means of optical spectroscopy (Papers I-IV), we have identified and studied 21 southern hard X--ray objects of unknown nature by means of the 1.5m CTIO telescope of Cerro Tololo (Chile).  We found that the selected sample is made of 12 AGNs (6 of which are of Seyfert 1 type and 6 are Seyfert 2 AGNs), 5 magnetic CVs and 4 HMXBs (one of which in the LMC). In terms of relative sizes of the three groups, we notice the absolute majority of AGNs in the sample, and a comparatively large fraction ($\\sim$25\\%) of CVs.  We recall that in three cases (IGR J14175$-$4641, IGR J14452$-$5133 and IGR J18244$-$5622), all identified as AGNs, only a tentative albeit likely optical counterpart was given because of the the lack of an univocal arcsecond-sized soft X--ray position. Thus, for them an observation with soft X--ray satellites affording arcsecond localizations (such as {\\it Chandra}, {\\it XMM-Newton} or {\\it Swift}) is needed to confirm the proposed association.  The results presented in this work further indicate the capabilities of {\\it INTEGRAL} to reveal not only high- and low-mass X-ray binaries, but also (if not mostly) extragalactic objects in the nearby Universe ($z <$ 0.1) and magnetic dwarf novae."
    },
    "0608/astro-ph0608507_arXiv.txt": {
        "abstract": "We observed the narrow-line quasar SDSS~J094857.3+002225, which has the highest known radio loudness for a narrow-line Seyfert~1 galaxy~(NLS1), at 1.7--15.4~GHz with the Very Long Baseline Array~(VLBA).  This is the first very-long-baseline interferometry~(VLBI) investigation for a radio-loud NLS1.  We independently found very high brightness temperatures from (1)~its compactness in a VLBA image and (2)~flux variation among the VLBA observation, our other observations with the VLBA, and the Very Large Array~(VLA).  A Doppler factor larger than 2.7--5.5 was required to meet an intrinsic limit of brightness temperature in the rest frame.  This is evidence for highly relativistic nonthermal jets in an NLS1.  We suggest that the Doppler factor is one of the most crucial parameters determining the radio loudness of NLS1s.  The accretion disk of SDSS~J094857.3+002225 is probably in the very high state, rather than the high/soft state, by analogy with X-ray binaries with strong radio outbursts and superluminal jets such as GRS~1915+105. ",
        "introduction": "Narrow-line Seyfert~1 galaxies (NLS1s; \\cite{Osterbrock&Pogge1985,Pogge2000}) are usually found as radio-quiet objects \\citep{Zhou_Wang2002}.  Various properties observed through optical and X-ray bands suggest that NLS1s are active galactic nuclei~(AGNs) of highly accreting~($L_\\mathrm{bol}/L_\\mathrm{Edd}\\sim1$, where $L_\\mathrm{bol}$ and $L_\\mathrm{Edd}$ are bolometric and Eddington luminosities, respectively) systems~(e.g.,~\\cite{Pounds_etal.1995,Boroson2002}).  The radio-quietness of NLS1s may be related to the suppression of radio jets emanated from accretion disks with high accretion rates \\citep{Greene_etal.2006}, as well as the disks of X-ray binaries in the {\\it high/soft} state (see, e.g.,~\\cite{McClintock&Remillard2003} for a review).  On the other hand, radio-loud NLS1s do exist, although they are very rare \\citep{Zhou_Wang2002}.  There is no significant difference between radio-quiet and radio-loud NLS1s in optical and X-ray properties (SDSS~J094857.3+002225,~\\cite{Zhou_etal.2003}; RX~J0134.2$-$4258,~\\cite{Grupe_etal.2000}; RGB~J0044+193,~\\cite{Siebert_etal.1999}; PKS~2004$-$447,~\\cite{Oshlack_etal.2001}; PKS~0558$-$504,~\\cite{Wang_etal.2001}).  The origin of the radio excess of the radio-loud NLS1s has been unknown.  Hardening of X-ray spectra during rapid X-ray flares of the radio-loud NLS1 PKS~0558$-$504 in a few minutes could arise from transient spectral dominance of synchrotron emission from relativistically boosted jets \\citep{Wang_etal.2001} similar to radio-loud quasars (e.g.,~\\cite{Reeves_etal.1997}).  However, properties of radio jets in NLS1s are slightly known.  Their linear sizes are $\\lesssim 300$~pc \\citep{Ulvestad_etal.1995}, which correspond to $\\sim1$~arcsec resolutions of connected interferometries; there are only a few images that have been obtained by very-long-baseline interferometry~(VLBI) observations~\\citep{Lal_etal.2004,Middelberg_etal.2004}.     \\begin{table*} \\caption{Parameters of observations.\\label{table1}} \\begin{center} \\begin{tabular}{lcrccc} \\hline\\hline \\multicolumn{1}{c}{Date} & Array & \\multicolumn{1}{c}{$\\nu_\\mathrm{c}$} & $\\Delta\\nu$ & pol. & $t_\\mathrm{scan} \\times N_\\mathrm{scan}$ \\\\ \\multicolumn{1}{c}{} &  & \\multicolumn{1}{c}{(GHz)} & (MHz) &  & (s) \\\\ \\multicolumn{1}{c}{(1)} & (2) & \\multicolumn{1}{c}{(3)} & (4) & (5) & (6) \\\\\\hline 2003Oct30 & VLBA & 15.365 & 32 & L & 46$\\times$52 \\\\ &  & 8.421 & 32 & R & 35$\\times$18 \\\\ &  & 4.987 & 32 & L & 56$\\times$7 \\\\ &  & 2.271 & 32 & R & 85$\\times$6 \\\\ &  & 1.667 & 32 & L & 94$\\times$6 \\\\ 2005May09 & VLBA & 1.667 & 32 & L & 46$\\times$12 \\\\ 2003Dec15 & VLA-B & 14.94 & 100 & dual & 43$\\times$2 \\\\ &  & 8.460 & 100 & dual & 33$\\times$1 \\\\ &  & 4.860 & 100 & dual & 33$\\times$1 \\\\ &  & 1.425 & 100 & dual & 43$\\times$1 \\\\\\hline \\end{tabular} \\end{center} {\\footnotesize Col.~(1) observation date; Col.~(2) array; Col.~(3) center frequency; Col.~(4) bandwidth; Col.~(5) polarization.  L and R denote left and right circular polarization, respectively; Col.~(6) scan lengths $\\times$ number of scans. } \\end{table*}     We started VLBI-imaging several NLS1s.  SDSS~J094857.3+002225 was discovered at $z=0.584$ in the Sloan Digital Sky Survey (SDSS) early data release \\citep{Williams_etal.2002}.  With FWHM(H$\\beta$)$\\approx$1500~km~s$^{-1}$ and undetected [OIII] lines, this source was consistent with the conventional definition of NLS1s \\citep{Zhou_etal.2003}.  Its radio counterpart was found in the Faint Images of the Radio Sky at Twenty-centimeter~(FIRST) survey \\citep{Becker_etal.1995}.  This has been the most radio-loud~($R\\approx 2000$) object in NLS1s \\citep{Zhou_etal.2003}, where $R$ is ``radio loudness'' conventionally defined as the ratio of 6~cm radio to {\\it B}-band flux densities with a threshold of $R=10$ separating radio-loud and radio-quiet objects (e.g.,~\\cite{Visnovsky_etal.1992,Stocke_etal.1992,Kellermann_etal.1994}).  \\citet{Zhou_etal.2003} suggested that its rapid ($\\sim$years) radio-flux variations were caused by relativistic boosting of a jet.  This is appropriate for the first target of VLBI observations, because of the brightest ($\\sim300$~mJy) NLS1 in radio bands.  In the present paper, we report VLBI observations in milliarcsecond~(mas) resolutions for SDSS~J094857.3+002225; this is the first VLBI study for a radio-loud NLS1.  Throughout this paper, a flat cosmology is assumed, with $H_0=71$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_\\mathrm{M}=0.27$, and $\\Omega_\\mathrm{\\Lambda}=0.73$ \\citep{Spergel_etal.2003}.  Therefore, 1~mas corresponds to 6.6~pc at the distance to SDSS~J094857.3+002225.  Its optical magnitude $m_B \\mathrm{(SDSS)} \\approx 18.89$~mag \\citep{Zhou_etal.2003} corresponding to $M_B=-23.8$~mag implies that this source is a quasar rather than a Seyfert galaxy according to the standard definition of a quasar, $M_B<-23$~mag. ",
        "conclusions": "\\label{section:discussion}  \\subsection{Very high brightness temperatures}\\label{section:brightnesstemperature}  We can derive brightness temperatures from both the VLBA image and the flux variability.  Using imaging analysis, brightness temperature in the rest frame is calculated using (cf. \\cite{Lahteenmaki_etal.1999}) \\begin{equation} T_\\mathrm{Brest}(\\mathrm{img})=\\frac{1+z}{\\delta} \\frac{c^2 S_{\\nu\\mathrm{obs}}^\\mathrm{obs}}{2k_\\mathrm{B}\\nu_\\mathrm{obs}^2 (\\theta_\\mathrm{obs} / 2)^2 \\pi}, \\label{equation:TB} \\end{equation} where $c$ is the speed of light, $k_\\mathrm{B}$ is the Boltzmann constant, $S_\\nu$ is the flux density at frequency $\\nu$, $\\theta$ is the angular diameter of a source, and suffixes of ``obs'' and ``rest'' denote quantities in the observer frame and the rest frame, respectively.  The Doppler factor is defined as $\\delta \\equiv \\sqrt{1-\\beta^2}/(1-\\beta\\cos{\\phi})$, where $\\beta \\equiv v/c$ ($v$ is the source speed), and $\\phi$ is the ``viewing angle'' between the direction of the source velocity and our line of sight.  Given an apparent diameter of $<$0.10~mas for C0 in the 15-GHz VLBA data (Section~\\ref{section:resultVLBA}), $T_\\mathrm{Brest}(\\mathrm{img})>5.5\\times10^{11}/\\delta$~K.  Alternatively, with flux variation, a brightness temperature in the rest frame can be calculated independently using (cf. \\cite{Lahteenmaki_etal.1999}) \\begin{equation} T_\\mathrm{Brest}(\\mathrm{var}) > \\frac{1+z}{\\delta^3} \\frac{D_\\mathrm{co}^2 \\Delta S_{\\nu_\\mathrm{obs}}^\\mathrm{obs}}{2 \\pi k_\\mathrm{B} \\nu_\\mathrm{obs}^2 (\\Delta t_\\mathrm{obs})^2}, \\label{equation:TBrest_variable} \\end{equation} where $D_\\mathrm{co}$ is the comoving distance, $\\Delta S_{\\nu_\\mathrm{obs}}^\\mathrm{obs}$ is a change of observed flux density in a period, of $\\Delta t_\\mathrm{obs}$.  In this equation, we assume exponential decay as seen in radio outbursts in blazars \\citep{Valtaoja_etal.1999}.  We found that the flux density decreases by 31\\% in 46~days at 15~GHz, and by 50\\% in 19 months at 1.7~GHz (Section~\\ref{section:resultflux}), which implies that $T_\\mathrm{Brest}(\\mathrm{var})>3.3\\times10^{13}/\\delta^3$~K and $T_\\mathrm{Brest}(\\mathrm{var})>1.1\\times10^{13}/\\delta^3$~K, respectively.  A very high brightness temperature in this source has also been suggested by \\citet{Zhou_etal.2003} from 1.4-GHz flux variation in $\\sim$years time scale.    These very high brightness temperatures form clear evidence for a nonthermal process in the central engine of this NLS1.           \\subsection{Highly relativistic jets}\\label{section:relativisticjet}  Based on near equipartition of energies between radiating particles and a magnetic field, the brightness temperature of nonthermal radio emission should be within a physically realistic upper limit of $\\sim2\\times10^{11}$~K in the rest frame \\citep{Readhead1994}.  Therefore, the Doppler factor should be $\\delta>2.7$--5.5 for the observed brightness temperatures (Section~\\ref{section:brightnesstemperature}).  The observed radio emissions are presumably from highly relativistic jets and are enhanced by Doppler boosting.  The jets should approach us with a viewing angle of $\\theta<$\\timeform{22D} and a speed of $\\beta>0.76$ for $\\delta>2.7$ (Section~\\ref{section:brightnesstemperature}).  C1, spatially resolved in the VLBA data (Section~\\ref{section:resultVLBA}), may be one of the expanding jet components.  Because of the very high brightness temperature, synchrotron self-absorption could be a possible explanation for the inverted spectrum as observed.  We show an example of spectral model fitting to the observed VLBA total fluxes using power-law spectra with a low-frequency cut-off due to synchrotron self-absorption: $S_\\nu = S_{\\nu_0} \\nu^{2.5}[1-\\exp{(-\\tau\\nu^{\\alpha_0-2.5})}]$, where $S_{\\nu_0}$ is a scaling constant, $\\tau$ is the optical depth, and $\\alpha_0$ is the intrinsic (not absorbed) spectral index.  Spectral model fitting was performed on the assumptions that the observed spectrum of C0 is caused by a superposition of self-absorbed synchrotron spectra of two components, and that the observed spectrum of C1 can be represented by a single power-law spectrum.  As a result, C0 consists of two convex spectra with peak frequencies at $\\sim$16~GHz~(``C0A'') and $\\sim$3~GHz~(``C0B''), as shown in Fig.~\\ref{figure2}.  The Doppler boosting can be responsible for such high frequencies of the spectral peaks, since $\\nu_\\mathrm{obs} = [\\delta/(1+z)] \\nu_\\mathrm{rest}$.  The intrinsic peak frequencies could be at several hundreds of MHz.  Such high-frequency peaked spectra are reminiscent of those of the Seyfert galaxy III~Zw~2 with two compact radio components and an inverted high-frequency spectrum in total flux \\citep{Falcke_etal.1999}.  These components in III~Zw~2 are also highly relativistic jets \\citep{Brunthaler_etal.2000}.      \\begin{figure} \\includegraphics[width=1.0\\linewidth]{figure2.eps} \\caption{Radio continuum spectra of SDSS~J094857.3+002225.  Symbols representing total flux densities are identified in top corner.  Filled circles particularly describe the flux densities for component C1.  The solid curve represents the spectral model that was determined by spectral model-fitting to the VLBA total fluxes using two self-absorbed synchrotron spectra (``C0A'' and ``C0B'' as dashed curves) with an intrinsic spectral index of $\\alpha_0=-0.6$ and one simple power-law spectrum (C1 as a dashed line).  The spatially single component C0 was divided into two spectrally separate components, C0A and C0B with peak frequencies of $\\sim$16 and 3~GHz, respectively.  See Section~\\ref{section:relativisticjet} for details.} \\label{figure2} \\end{figure}         \\subsection{Implications of SDSS~J094857.3+002225} Our observations showed a slightly inverted spectrum at $\\sim$5~GHz and $R\\approx2530$, which is derived from $m_B \\mathrm{(SDSS)} \\approx 18.89$~mag \\citep{Zhou_etal.2003} and the VLBA total flux density at 5~GHz.  Given $\\delta>5.5$ (Section~\\ref{section:brightnesstemperature}), the peak frequency of C0A, $\\sim$16~GHz, should be $<$2.9~GHz before Doppler beaming.  In the case where $\\alpha=-0.6$, a flux density should be amplified by a factor of $>$462 from an intrinsic spectrum at the same frequency.  Therefore, there is a possibility that SDSS~J094857.3+002225 is a radio-quiet object with a steep spectrum in the rest frame, like typical NLS1s \\citep{Moran2000,Zhou_Wang2002,Ulvestad_etal.1995}.  Flat or inverted radio spectra have also been found in the other radio-loud NLS1s, e.g., RGB~0044+193 \\citep{Siebert_etal.1999}, 2E~1640+5345, RXS~J16290+4007, and RX~J16446+2619 \\citep{Zhou_Wang2002}.  Doppler boosting may be responsible for their spectra and radio loudness as well as those of SDSS~J094857.3+002225.  The observed optical and X-ray properties and black-hole masses of radio-loud NLS1s are not significantly different from those of radio-quiet ones (e.g., \\cite{Siebert_etal.1999,Wang_etal.2001}).  From our observations, we suggest that the Doppler factor is a crucial parameter determining the radio loudness of NLS1s: at least some of NLS1s could possess the ability to generate highly relativistic jets.   Analogies have been drawn between X-ray binaries and AGNs (e.g., \\cite{Maccarone_etal.2003}).  NLS1s are thought to be AGNs with high accretion rates.  Although X-ray binaries in the {\\it high/soft} state with accretion rates of $L/L_\\mathrm{Edd}\\sim0.03$--0.1 have no relevance to jets, strong radio flares have been seen during the {\\it very high} state, which is thought to be a disk in $L/L_\\mathrm{Edd}>0.1$, such as GRS~1915+105, with highly relativistic jets (e.g., \\cite{Fender2003}).  SDSS~J094857.3+002225 has the potential to be compatible with the {\\it very high} state, rather than the {\\it high/soft} state.  Radio-loud NLS1s will be appropriate targets to investigate for the {\\it very high} state on super-massive black hole systems."
    },
    "0608/astro-ph0608057_arXiv.txt": {
        "abstract": "{ We present models for the mid- and far- infrared emission from the Narrow Line Region (NLR) of Active Galactic Nuclei (AGN). Using the \\mapiii\\ code we explore the effect of typical NLR parameters on the spectral characteristics of the IR emission. These include useful IR emission line ratio diagnostic diagrams for the determination of these parameters, as well as Star formation--AGN mixing diagnostics. We also examine emission line to continuum correlations which would assist in separating the IR emission arising from the NLR from that coming from the inner torus. We find for AGN like NGC 1068 and NGC 4151 that the NLR only contributes $\\sim$10\\% to the total IRAS 25 \\mum\\ flux, and that other components such as a dusty torus are necessary to explain the total AGN  IR emission.} {} ",
        "introduction": "Active galactic nuclei (AGN) are well known to emit strongly at all wavelengths, from gamma-rays and X-rays to the infrared and radio. Their emission characteristics at each wavelength gives us insight into the central engine of these galaxies, whether observed directly, such as at X-ray wavelengths, or indirectly through reprocessed emission, such as by study of the emission lines from the narrow line region (NLR).  One area that has seen recent rapid advances is studies of the mid-- and far--infrared characteristics of AGN, with space telescopes such as IRAS, ISO and Spitzer making this wavelength regime readily accessible. As the continuum in this wavelength region is dominated by  warm and hot  dust re-emission, it is particularly useful in testing the paradigm of the unified  model, which argues for a dense equatorial band of dust surrounding the black hole central engine \\citep{Antonucci85}.  Various attempts have been made to model the mid-- and far--IR emission.  These range from simple empirical fits of dust grey-body emission to the observations \\citep{Edelson86,Blain03}, to more complicated distributions of dust which include the stochastic heating of dust by an AGN power-law continuum \\citep{Dullemond05}. With the prevalence of the unified model, most analyses of the IR from AGN have concentrated on the emission from dust lying in a dense torus surrounding a nuclear heating source \\citep{Pier92,Pier93,Granato94,Efstathiou95, Granato97,vanBemmel04,Nenkova02}. However, there have been some authors who have considered different geometries \\citep[i.e.][]{Siebenmorgen04} or have considered in more detail the collisional heating and destruction of dust \\citep{Contini04}.  Often ignored in the considerations of AGN IR emission is the contribution from the narrow line region (NLR). This region is known to contain dust \\citep[e.g.][]{Radomski03} and, like the dusty torus, will be heated by the central engine and re-emit at IR wavelengths \\citep{Netzer93}. With lower densities than found in the central torus, the NLR  also emits strong forbidden lines  such as the  [\\oiii] 500.7 nm line. Such lines are often one of the main diagnostics for the presence of an AGN, and also provide important constraints on the form and intensity of the radiation field which excites them. They also provide insight into the physical conditions such as density or pressure, and constraints on the metallicity of the circum-nuclear ISM.  As AGN are often heavily dust-embedded, much of the line emission from the near nuclear parts of the NLR may be totally obscured at optical wavelengths. To penetrate this circum-nuclear dust while retaining adequate spatial resolution to investigate the structure of the dense parts of the NLR, we need to observe at mid-IR wavelengths. Instruments such as the \\emph{IRS} instrument of the \\emph{Spitzer Space Telescope} are finally providing the data needed to study this region. In order to interpret this data we need both line and continuum diagnostics in the mid-IR.  Here,  we will explore the expected characteristics of the infrared emission from the NLR, examine the dependence of this emission on the parameters that define the NLR, and investigate possible diagnostics that can be used to define these parameters and so give us insight into the NLR and the AGN as a whole. ",
        "conclusions": "We have explored within this work a series of Dusty Narrow Line Region (NLR) models, covering a range of realistic pressures and ionization parameters, with the aim of investigating the Infrared emission from dust in these regions.  The presented spectral energy distributions (SEDs) for these models reveal the possible range of the IR emission from the NLR, and the correlation of this emission with the ionization parameter and the intensity of the ionizing radiation. An increase in either of these parameters will increase both the total IR flux from the dust and the peak temperature of the dust.  Comparisons with observed NLR IR spectra show that simple 1-D photoionization models are not enough to represent the complexity of the NLR. Combinations of different NLR clouds and NLR dynamics may be needed to represent a more accurate model of the NLR IR emission.  While examining the spectral energy distributions of these models gives an idea of both the overall emission from the NLR of active galaxies, the greatest benefit of these models is that they give a direct connection between the emission by dust and the line emission from the photoionized gas.  A selection of correlations between line and dust emission are presented here to give possible diagnostics for the IR emission from AGN and galaxies. In addition we have explored further possible IR emission line diagnostics to assist in the determination of the NLR physical conditions, and to help distinguish the contribution that starformation makes to the total spectrum.  Comparisons of the models with the nearby AGN, NGC 1068 and NGC 4151, appear to show that only $\\sim$10\\% of the 25\\mum\\ emission arises from the [\\oiii] line emitting region. This means a significant fraction must arise from either the putative torus, nuclear star formation, coronal--shock excited region, or some combination thereof.  While the fraction the NLR contributes to the overall infrared continuum of an active galaxy may be small, it is important to consider this component when trying to understand the full spectral energy distribution of AGN. The final aim will be to connect the emission from the NLR with what is occurring in the central engine and thus connect the emission lines with the total IR output of the AGN."
    },
    "0608/astro-ph0608261_arXiv.txt": {
        "abstract": "We study the dynamical evolution of globular clusters containing populations of primordial binaries, using our newly updated Monte Carlo cluster evolution code with the inclusion of direct integration of binary scattering interactions. We describe the modifications we have made to the code, as well as improvements we have made to the core Monte Carlo method.  We present several test calculations to verify the validity of the new code, and perform many comparisons with previous analytical and numerical work in the literature.  We simulate the evolution of a large grid of models, with a wide range of initial cluster profiles, and with binary fractions ranging from 0 to 1, and compare with observations of Galactic globular clusters.  We find that our code yields very good agreement with direct $N$-body simulations of clusters with primordial binaries, but yields some results that differ significantly from other approximate methods.  Notably, the direct integration of binary interactions reduces their energy generation rate relative to the simple recipes used in Paper III, and yields smaller core radii.  Our results for the structural parameters of clusters during the binary-burning phase are now in the tail of the range of parameters for observed clusters, implying that either clusters are born significantly more or less centrally concentrated than has been previously considered, or that there are additional physical processes beyond two-body relaxation and binary interactions that affect the structural characteristics of clusters. ",
        "introduction": "\\label{sec:intro}  Observations \\citep[e.g.,][]{1996AJ....112..574C,2002AJ....123.2541B,1997ApJ...474..701R,2002ASPC..263..163C,2002AJ....123.1509B}, in combination with recent theoretical work \\citep{2005MNRAS.358..572I}, suggest that although the currently observed binary fractions in the cores of globular clusters may be small ($\\lesssim 10\\%$), the initial cluster binary fraction may have been significantly larger ($\\gtrsim 50\\%$).  As has been understood theoretically for some time, primordial binaries in star clusters generally act as an energy source (through super-elastic scattering encounters), producing energy in the core and postponing deep core collapse in a quasi-steady state ``binary-burning'' phase. This is analogous to the long-lived main-sequence in stars, in which hydrogen is burned to prevent collapse.  An initial binary fraction of a few percent is enough to postpone deep core collapse for many initial relaxation times \\citep[see][for discussion and references]{2003ApJ...593..772F}. In addition to playing a large part in the global evolution of a star cluster, dynamical interactions of binaries also strongly affect the formation and evolution of stellar and binary exotica, which include low-mass X-ray binaries, recycled pulsars, cataclysmic variables, and blue stragglers \\citep{1991A&A...241..137H,1995ApJS...99..609S,1998MNRAS.301...15D,2000ApJ...532L..47R,2005MNRAS.358..572I}.  Similarly to dynamical binary interactions, physical stellar collisions (both direct star--star collisions, and those mediated by resonant binary interactions) also play an important role in the evolution of globular cluster populations.  Stellar collisions are thought to be one of the two primary mechanisms by which blue stragglers are created in dense star clusters \\citep[e.g.,][]{2004ApJ...605L..29M}.  A runaway sequence of stellar collisions of massive main sequence stars early in the lifetime of a dense cluster may yield a very massive star ($\\gtrsim 10^3\\,M_\\sun$), which may then become an intermediate mass black hole \\citep[see, e.g.,][for discussion and references]{2006ApJ...640L..39G}.  The complete picture of star cluster evolution can be rather complicated, as it includes, in addition to the two physical processes just mentioned, single star evolution, binary star evolution, and tidal stripping due to the field of the Galaxy. Even if one ignores these additional processes, modeling a dense stellar cluster with primordial binaries still presents a formidable computational challenge.  There are at least two reasons for this: 1) dynamical interactions of binaries---which are typically resonant, lasting for many orbits---must be resolved on their natural timescale, which is orders of magnitude shorter than the cluster relaxation time, and 2) primordial binaries extend the life of a cluster.  Although the GRAPE series of special-purpose computers is steadily increasing in performance, direct $N$-body simulation of the evolution of clusters with more than a few percent binaries and a moderate number of stars ($\\sim 10^5$) is still quite computationally expensive, with computational timescales on the order of months.  Faster, more approximate methods, such as the anisotropic gas model or direct solution of the Fokker-Planck equation, suffer from the difficulty inherent in incorporating physics beyond two-body relaxation, such as stellar evolution or binary interactions, in these methods. The Monte Carlo method bridges the gap between these two computational extremes, since it allows for the relatively facile inclusion of additional layers of physics, provides a star-by-star description of a cluster, and is computationally inexpensive.  In previous studies using approximate methods like Monte Carlo or Fokker-Planck, binary interactions had generally been treated using recipes culled from the results of large numbers of numerical scattering experiments.  (The work of \\citet{2003MNRAS.343..781G}, which incorporates direct integration of binary interactions, is one notable exception.)  The recipes are typically known only for equal-mass binary interactions, thus prohibiting the use of a cluster mass function. Thus in order to model realistic clusters, which contain a wide range of masses, one must numerically integrate each binary interaction in order to resolve it properly. We have now incorporated into our Monte Carlo code a dynamical integrator to exactly integrate dynamical interactions of binaries, allowing us to evolve clusters with mass spectra, and perform more realistic comparisons with direct $N$-body calculations.  As we demonstrate below, the direct integration of binary interactions reduces their energy generation rate relative to the simple recipes used in Paper III, and yields smaller core radii.  This is the fourth paper in a series studying the evolution of globular clusters using the Monte Carlo method.  Paper I describes the core method and presents several test calculations exhibiting the validity of the code \\citep{2000ApJ...540..969J}. Paper II treats the evolution of tidally-truncated clusters with mass spectra \\citep{2001ApJ...550..691J}.  Paper III adds a recipes-based treatment of binary scattering interactions and considers the evolution of an ensemble of clusters of varying initial binary fraction and central concentration, finding that even a small fraction of binaries in a cluster is sufficient to support the core against collapse significantly beyond the normal core-collapse time \\citep[][hereafter Paper III]{2003ApJ...593..772F}.  In this, the fourth paper in the series, we describe our new code, perform several tests to ensure its validity, and perform a large set of simulations of clusters with primordial binaries, which we compare with previous results in the literature and with observations of Galactic globular clusters.  Section \\ref{sec:method} describes our new code in detail, including the additional physical processes we have added (numerical integration of binary scattering interactions, and star--star physical collisions), as well as the improvements we have made to the core Monte Carlo method.  Section \\ref{sec:exampleresults} presents a few example results, and compares them with semi-analytical theory and previous numerical calculations. Section \\ref{sec:results} describes the trends evident in the grid of cluster models we simulated, and compares our results with observations.  Finally, in section \\ref{sec:summary} we summarize and conclude. ",
        "conclusions": "\\label{sec:summary}  In this paper we have described our new Monte Carlo evolution code and used it to perform a large set of cluster evolution simulations, which we compared with previous results in the literature, as well as observations of Galactic globular clusters.  In section \\ref{sec:method} we described our new code in detail, including the implementation of direct integration of binary scattering interactions and star--star physical collisions, and the fundamental modifications we have made to the core Monte Carlo method.  We performed several test calculations with the code, finding that it reproduces well several standard results.  It yields a core collapse time of $\\approx 18\\,t_{\\rm rh}$ for an isolated Plummer model, in good agreement with results in the literature.  It also produces an $r^{-2.3}$ density profile during the late stages of core collapse, in good agreement with the theoretical expectation.  We also performed comparisons of clusters with increasingly wide mass spectra with $N$-body, finding that for moderately wide mass spectra ($1$ to $10\\,M_\\sun$) the agreement with $N$-body is satisfactory, but for very wide mass spectra ($0.2$ to $120\\,M_\\sun$) the agreement is not as good. In particular, for such wide mass spectra, our Monte Carlo code tends to overestimate the mass segregation timescale at early times, and underestimate it at later times. The sense of the disagreement is the same as found by \\citet{2006MNRAS.368..121F} with their Monte Carlo code.  In section \\ref{sec:exampleresults} we displayed a few example results and compared with theory and previous numerical calculations in the literature.  We found that the code conserves energy well over the long timescales of our runs. We compared our predicted values of $r_c/r_h$ during the quasi-steady state binary burning phase with the semi-analytical work of \\citet{1994ApJ...431..231V}, finding good agreement.  We compared our results with previous numerical results in the literature for isolated and tidally-truncated cluster models, finding excellent agreement with $N$-body calculations.  There are much larger discrepancies with the other approximate methods (Fokker-Planck and other Monte Carlo codes), which is to be expected, since most used recipes for binary interactions, which are known to overestimate the energy generation rate in binaries.  In section \\ref{sec:results} we surveyed the results from all our simulations, and compared with observations.  Our models cover a large range in parameter space, using Plummer and King models with $W_0=3$ to 11 for the initial cluster profile, and with initial binary fractions from 0 to 1.  The resulting structural parameters in the binary burning phase span a remarkably small range, with $r_c/r_h$ varying from 0.03 to 0.12, and $c$ varying from 1.7 to 2.4. Our results for these structural parameters are distinctly different from the results found with the previous version of our code (which used recipes for binary interactions), with $r_c/r_h$ now smaller than what we found in Paper III by a factor of up to $\\sim 10$, and with $c$ larger by $\\sim 0.5$.  Although our new results agree much better with $N$-body calculations, they unfortunately agree much less well than in Paper III with the observations.  The disagreement implies one of at least two possibilities.  It may be that the initial conditions for Galactic globular clusters are outside the range of initial conditions we have sampled in this work.   Extrapolation of our results suggests that clusters with $W_0 \\gtrsim 11$ and $f_b \\lesssim 1$ may match the observations.  This is intriguing since a cluster with $W_0 \\gtrsim 11$ will likely form an IMBH early in its lifetime via a collisional runaway, and clusters with central IMBHs are known to have larger values of $r_c/r_h$ \\citep{2006MNRAS.368..141F,astro-ph/0610342}. Alternatively, stellar evolution and collisions, which are not included in the simulations in this paper, could possibly explain the disagreement. However, it appears that the effect of these processes should act in the opposite sense of ameliorating the disagreement with observations.  More detailed simulations including the effects of single- and binary-star evolution and physical collisions should be performed to test this."
    },
    "0608/astro-ph0608111_arXiv.txt": {
        "abstract": "RBS (Rosat Bright Source) 0490 is a cataclysmic variable star (CV) with unusually strong emission lines.  The strength of the emission lines led to a suggestion that the object is intrinsically faint and correspondingly nearby ($\\sim 33$ pc), which, if true, would strongly affect estimates of the CV space density. Here we report astrometry, filter photometry, and time-series spectroscopy of this object.  The astrometry gives an absolute parallax $\\pi_{\\rm abs} = 4.5 \\pm 1.5$ mas and a relative proper motion of 102 mas yr$^{-1}$.  A Bayesian procedure gives a very uncertain distance estimate of $d \\sim 300$ pc, and the small parallax alone implies $d > 133$ pc (at two standard deviations). The mean $V$ magnitude is 17.4, which implies $M_V = 10.9 - 5 \\log(d/200\\ {\\rm pc})$, neglecting extinction.  At 200 pc, the space velocity would be over 90 km s$^{-1}$ with respect to the LSR.  The time-series spectroscopy shows a possible emission-line radial-velocity period near 46 min. This would be unusually short for an orbital period and it may represent some other clock in the system. ",
        "introduction": "Cataclysmic variable stars (CVs) are close binaries in which a white dwarf accretes matter from a less-evolved companion (the secondary), which usually resembles a lower main sequence star.  CVs have a rich phenomenology, and the theory of CVs ranges across wide areas of astrophysics; \\citet{warn} presents a comprehensive review.  CVs are common enough that there are several with distance $d < 100$ pc, but their space density remains somewhat controversial, in part because of distance uncertainties.  Also, the completeness of the known sample is not perfectly understood. For example, all-sky camera surveys (e.g., \\citealt{asasdescr}) frequently turn up previously undiscovered, bright dwarf novae in outburst. Many new CVs have also been found in the Sloan Digital Sky Survey (SDSS; \\citealt{szkodysdss1,szkodysdss2, szkodysdss3,szkodysdss4,szkodysdss5}), but as \\citet{gaensickegoettingen} points out, the colors of the CVs turned up by SDSS range down to the cutoff used to trigger spectroscopic followup, so more presumably remain undiscovered.  \\citet{pattlate} gives a judicious estimate of the space density as $10^{-5}$ pc$^{-3}$. The space density of CVs is predicted to be an order of magnitude higher by the evolution scenario of cataclysmic binaries \\citep{shara86}. Unravelling the long-term evolution of CVs is largely dependent on an accurate measure of their space density.  Essentially all CVs are X-ray sources at some level, so a great many CVs have been discovered as the optical counterparts of X-ray sources. Because the X-ray surveys have well-understood completeness limits, these surveys have the potential to put CV space density estimates on a firmer footing (though very weak X-ray emitters will still be missed).  \\citet{schwope} presented new optical identifications of CVs in the Rosat Bright Source (RBS) list.  For many of their sources they estimated the absolute magnitude $M_V$ using an empirical relationship between $M_V$ and the emission equivalent width of H$\\beta$ (\\citealt{patt84}; his eqn. 16).  One source -- RBS 0490 -- had a very large EW(H$\\beta$) = 230 \\AA , for which the Patterson relation predicts a quite faint $M_V = 13.4$.  This, combined with an estimated $m_V = 16$, gave a distance of only 33 pc. RBS 0490 was one of two objects for which \\citet{schwope} estimated $d \\sim 30$ pc, the other being RBS 1955.  Applying the $V/V_{\\rm max}$ method to their data set, \\citet{schwope} estimated a relatively high CV space density, of order $3 \\times 10^{-5}$ pc$^{-3}$.  However, the inference of this high space density rested almost entirely on these two objects.  Largely because of the influence these two objects exert on the the space density estimate, we undertook to characterize them and estimate their distances.  A study of RBS 1955 will be published elsewhere; here we concentrate on RBS 0490, which lies in Eridanus at $\\alpha = 3^{\\rm h} 54^{\\rm m} 10^{\\rm s}.3, \\delta = -16^\\circ 52' 50''$ (ICRS). ",
        "conclusions": "\\subsection{Orbital Period}  Ordinarily in CV studies, an emission-line radial velocity period can be confidently identified with the orbital period.  Our apparent $\\sim 46$ min period challenges that assumption.  CVs with `normal' hydrogen-burning secondary stars are thought to evolve to shorter periods as they lose angular momentum, pass through a period mimimum near $P_{\\rm orb} \\sim 70$ min, and then evolve toward longer periods \\citep{pattlate}.  This is called the `period bounce' (although it does not take place as abruptly as that term suggests). Most of the relatively small number of known CVs with periods below the `bounce' show no hydrogen lines and appear to have helium secondaries; these are the AM CVn stars.  Of the small handful of CVs with periods below the period minimum that do not obviously have helium secondaries, at least one -- EI Psc --- has an anomalously warm secondary that is evidently somewhat evolved \\citep{eipsc}, and the helium lines in both EI Psc and the 59-minute binary V485 Cen \\citep{augusteijnv485} appear at least a little strong when compared with the ordinary hydrogen-accretors at longer periods.  In RBS 0490 the emission lines are remarkably strong, but the He lines do not appear to be enhanced, so an orbital period as short as 46 minutes seems implausible.  In recent years, several objects have been discovered that show radial velocity periods which are clearly not the same as $P_{\\rm orb}$.  The emission velocities in HS 2331+3905 show a remarkably convincing quasi-period near 3.5 hr \\citep{hasitall}, but the periodicity does not retain coherence over longer timescales, and weak eclipses indicate a stable a $P_{\\rm orb}$ near 81 min.   In DW Cnc, the emission-line radial velocities reveal stable periodicities at 86.1 and 38.6 min, with comparable amplitude; it appears that these periodicities represent respectively the orbit and the white-dwarf spin \\citep{dwcnc,dwcncrodriguez}.  These examples add weight to the argument that the $\\sim 46$ min period in RBS 0490 is not orbital.  But, if the 46-min period is {\\it not} $P_{\\rm orb}$, then the absence of any radial-velocity modulation at a more convincing orbital frequency is somewhat surprising; while CV emission lines often do not trace the motion of the white dwarf accurately, they usually do show some modulation at $P_{\\rm orb}$.   The full-width at half maximum of the H$\\alpha$ emission line is 1000 km s$^{-1}$, so it is unlikely that the system is nearly face-on.  It is possible that the secondary's mass is very small.  The spectrum shows no trace of the secondary, but given the still-unknown $P_{\\rm orb}$ and rather uncertain distance, the non-detection is not readily interpreted physically.  Fortunately, rather than crying vaguely for `more studies', we can suggest a very specific approach to unscramble this object: extensive time-series photometry.  There is a good chance that this would reveal other periodicities that would point the way toward understanding this object.  \\subsection{Distance and Demographics}  Our mean $V = 17.4$ implies an absolute magnitude $M_V = 10.9 - 5 \\log_{10}(d / 200 \\hbox{ pc})$, or $M_V < 11.8$ using our lower-limit distance of 133 pc.  As noted earlier, the $M_V$-EW(H$\\beta$) relation from \\citet{patt84} predicts $M_V = 13.4$ from the very strong emission. However, the relationship shows considerable scatter, and the calibration points used by \\citet{patt84} do not extend to objects with lines this strong.  One must clearly be cautious in using this broad-brush correlation to estimate the luminosities of individual objects; also, the lines in this object really do appear extraordinarily strong for the luminosity.  The very short distance estimated by \\citet{schwope} was also based on a value of $V \\sim 16$, estimated from spectrophotometry; this is over 1 magnitude brighter than we have seen it in our four photometric visits.   In any case, the short distance is excluded by the small parallax.   Also, an MDM Observatory parallax study of RBS 1955 (in preparation) gives a distance of $\\sim 160$ pc for that source, over five times the \\citet{schwope} estimate.  Thus it appears that the high space density derived in that paper should be disregarded.  The appropriate CV subclassification for RBS 0490 is unclear. The relatively short period in the emission lines suggests a white dwarf spin period in a magnetic system, but HeII $\\lambda$4686 is rather weak, while in most magnetic CVs it has considerable strength.  On the other hand, HeII $\\lambda$4686 is not particularly strong in DW Cnc or HS 2331+3905, either, so this is not conclusive. The absolute magnitude is comparable to those of dwarf novae in quiescence, and much less bright than novalike variables.  Using our upper-limit distance of 405 pc, we find $M_V > 9.4$, which is still well short of novalike brightness.  {\\it Acknowledgments.} We gratefully acknowledge support through NSF grants AST 9987334 and AST 0307413, and thank Bill Fenton for taking some of the spectroscopic data.  \\clearpage"
    },
    "0608/astro-ph0608327_arXiv.txt": {
        "abstract": "We present a high-resolution Keck/ESI spectrum of \\grb, which exhibits four absorption systems at $z=1.04329$, $1.95260$, $1.96337$, and $1.98691$.  The two highest redshift systems, separated by about 2400 km s$^{-1}$, have been previously suspected as kinematic features arising in the circumstellar wind around the progenitor star. However, the high column densities of low-ionization species (including possibly neutral hydrogen) in the blue-shifted system, are inconsistent with the expected highly ionized state of the circumstellar wind from the massive progenitor star, even prior to the GRB explosion.  This conclusion is also supported by the lack of detectable absorption from fine-structure transitions of \\ion{Si}{2} and \\ion{Fe}{2}.  Instead we conclude that the two redshift systems are similar to multiple DLAs found in QSO sight lines with a similar velocity separation and chemical abundance of [Cr/Fe] and [Zn/Fe]. The absorption system at $z=1.96337$ is likely an intervening low-mass galaxy, possibly related to the GRB host as part of a forming large-scale structure. ",
        "introduction": "\\label{sec:intro}  The idea that the progenitors of long-duration $\\gamma$-ray bursts (GRBs) are massive stars is now well-established, based in particular on the detection of associated type Ic supernovae (e.g., \\citealt{gvv+98,smg+03}).  The nature of these supernovae, and theoretical considerations, have lead to the suggestion that the progenitors are massive Wolf-Rayet (WR) stars \\citep{mw99}.  Such stars shed a considerable fraction of their mass in a fast ($\\sim 10^3$ km s$^{-1}$) wind (e.g., \\citealt{abb78}).  As a result, absorption spectroscopy of GRB afterglows can in principle probe the composition and kinematics of the wind, providing in turn direct information on the nature and identity of the progenitor star. Previous observations have in fact uncovered such blue-shifted absorption systems, which have been used to argue for carbon-rich (WC) Wolf-Rayet progenitors (e.g., \\citealt{mhc+03,fdl+05,bpc+06}), potentially in binary systems \\citep{swh+05,bpc+06}.  One central question to answer in applying the potentially powerful diagnostics from absorption lines is whether blue-shifted absorbers arise from outflowing circumstellar gas, or separate intervening absorbers arising from different galaxies.  In addition, the intense radiation field produced by the GRB can modify the circumstellar environment through ionization, excitation, and destruction of dust grains (e.g., \\citealt{pl98,dh02}).  However, these processes, as well as the relative abundance of low- and high-ionization ions, may in fact serve to distinguish the location of the absorbers \\citep{pcb06}.  Here we analyze a high resolution spectrum of \\grb\\ in a continued effort to characterize the circumstellar and interstellar environments around GRBs.  This burst is of particular interest because it is one of only a few that exhibit a blue-shifted absorption system with a velocity of a ${\\rm few}\\times 10^3$ km s$^{-1}$ \\citep{kgr+04}, possibly representing the circumstellar wind from the progenitor star. However, based on the relative strengths of the low- and high-ionization lines, and the lack of detectable fine-structure transitions we conclude that the blue-shifted absorption system is not circumstellar.  Instead we conclude that the blue-shifted system is an intervening galaxy similar to the multiple DLAs found in some quasar spectra \\citep{el01,le03}. ",
        "conclusions": "\\label{sec:disc}  We find four absorption systems in the spectrum of \\grb, of which the two highest redshift systems are separated by only 2400 km s$^{-1}$. Absorber $z_2$ is separated by about 3400 km s$^{-1}$ from $z_4$ and 1100 km s$^{-1}$ from $z_3$, but is detected only in \\ion{C}{4}.  As such it most likely arises in the halo of an intervening galaxy. Although $z_3$ was previously suspected to arise in the outflowing circumstellar gas surrounding the progenitor star \\citep{kgr+04}, we conclude that this scenario is unlikely based on the presence of high column density low-ionization species, the lack of excited fine-structure lines, the large column density of iron, and the possible presence of large column density of neutral hydrogen.  A more likely scenario is that system $z_3$ is an intervening absorber.  This conclusion leaves only GRBs 021004 and 050505 as showing potential evidence for wind outflows from the progenitor star.  In the former case, the presence of high velocity neutral hydrogen, as well as \\ion{Mg}{2} and \\ion{Fe}{2}, presents a challenge in the context of WR wind models and the ionizing radiation of the burst \\citep{swh+05}. In the latter case, only high-ionization \\ion{C}{4} is detected in outflow \\citep{bpc+06}, removing the problem of explaining the presence of low-ionization lines, but the low resolution of the spectrum prevents a detailed study.  Given the influence of the kinematic and ionization structure of the wind in the pre-explosion environment, as well as the effects of the GRB radiation field itself, it is of prime importance to make model predictions for the type and column density of various ionic species that can be observed in the afterglow spectrum.  Such models can be used to assess whether any of the outflowing systems observed to date are in fact signatures of the stellar wind.  Assessing $z_3$ as an intervening absorber, we note that both $z_3$ and the host galaxy absorber $z_4$ exhibit similar abundances and kinematics compared to previous GRB-DLAs and multiple DLAs found in QSO sight lines.  Assuming that iron is weakly depleted (as suggested by several lines of evidence, see \\S\\ref{sec:abund}) the gas metallicity is somewhat lower than in most previous GRB-DLAs, ${\\rm [Fe/H]}\\sim -1.5$ and $\\sim -1.2$, respectively.  Since $z_3$ most likely represents an intervening galaxy, we may be able to identify it in deep imaging around the GRB position, and spectroscopy of possible intervening galaxies in the field.  Previous narrow-band imaging observations aimed at finding Ly$\\alpha$ emission in the range $z\\approx 1.94-2.0$, have uncovered three galaxies that may have similar redshifts to the $z_2$ and $z_3$ absorbers \\citep{jbf+05}.  No source was detected at the GRB position.  These putative counterparts are faint and it is reasonable to interpret both the undetected GRB host galaxy and the intervening systems as low mass galaxies with low metallicity and dust content.  The remaining open question is whether there is a link between the GRB host absorber and the intervening systems.  It has recently been shown that the probability of finding multiple intervening absorbers toward GRBs is larger compared to QSO sight lines \\citep{ppc+06}.  However, the majority of these systems are separated by $>10^4$ km s$^{-1}$ and are unlikely to be physically connected.  With a velocity separation of 2400 km s$^{-1}$ it is conceivable that $z_3$ and the GRB host are part of the same large-scale structure, as has been suggested for the multiple DLAs in QSO sight lines \\citep{le03}.  Future observations of the environment of \\grb, as well as similar absorption systems may shed light on the dependence of the number and velocity difference of intervening systems on the properties of GRB host galaxies and their large-scale environments."
    },
    "0608/nucl-th0608080_arXiv.txt": {
        "abstract": "We have studied the effects of the nucleon weak magnetism and the neutrino electromagnetic properties on the neutrino and antineutrino mean free paths for all types of neutrinos. In the calculation, we consider matters with and without neutrino trapping as the target. We have found that the difference between the two mean free paths depends not only on the neutrino energy, but also on the matter density as well as matter constituents. Compared to the nucleon weak magnetism, the neutrino electromagnetic properties are found to have negligible effects on this difference. ",
        "introduction": "In the standard model, neutrinos have zero mass, magnetic moment, as well as electronic charge. However, there have been experimental upper limits on the electron-, muon- and tau-neutrino dipole moments given by ${\\mu}_{\\nu_{e}} < 1.5\\times 10^{-10} {\\mu}_{B}$ \\cite{Beac99} , ${\\mu}_{\\nu_{\\mu}} < 7.4\\times 10^{-10} {\\mu}_{B}$ \\cite{Krak90} and ${\\mu}_{\\nu_{\\tau}} < 5.4\\times 10^{-7}{\\mu}_{B}$ \\cite{Coop92}, where ${\\mu}_{B}$ is the Bohr magneton. While neutrino charge radii of $R = (0.54 \\pm 0.82)\\times 10^{-6}~{\\rm MeV}^{-1}$  (for electron-neutrino) and $R <0.40 \\times 10^{-6}~{\\rm MeV}^{-1}$ \\cite{Vila95} (for muon-neutrino) have been experimentally evident, no experimental bound has been found so far on the charge radius of tau-neutrino \\cite{Josh01}. On the other hand, the corresponding bounds from Super-K and SNO observations \\cite{Josh01} are substantially larger, i.e., ${\\mu}_{\\nu_{\\tau}}, {\\mu}_{\\nu_{\\mu}} < 6.73 ~(5.77)\\times 10^{-10} {\\mu}_{B}$ and ${\\mu}_{\\nu_{e}} < 6.45 ~(5.65) \\times 10^{-10} {\\mu}_{B}$, whereas for the neutrino charge radii they obtained that $R_{\\nu_{\\mu}}, R_{\\nu_{\\tau}} < 2.31~(1.98) \\times 10^{-6}~{\\rm MeV}^{-1}$ and $R_{\\nu_{e}} <1.33~(1.16) \\times 10^{-6}~{\\rm MeV}^{-1}$. Moreover, various astrophysical observations provide a limit on the neutrino magnetic moment in the range of $(1-4) \\times 10^{-10} {\\mu}_{B}$ \\cite{Hagi02}, while from the plasmon decay in globular-cluster stars it is found that ${\\mu}_{\\nu} < 3\\times 10^{-12} {\\mu}_{B}$  and $e_{\\nu} < 2\\times 10^{-14}e$ \\cite{Raff99}, where $e$ is the proton charge.  Using the standard neutrino properties, Horowitz and Garc\\'ia \\cite{Horo05} found that the muon-antineutrino mean free path in high density matter was considerably larger than the muon-neutrino one, provided that the weak magnetism of the nucleons is included. The sensitivity of differential cross section of the interaction between neutrinos and dense matter to the possibly nonzero neutrino electromagnetic properties has been also investigated \\cite{Will05,Sul06}. It is found that the effects of the neutrino electromagnetic properties on the differential cross section become more significant for the neutrino magnetic moment ${\\mu}_{\\nu} > 10^{-10} {\\mu}_{B}$ and the neutrino charge radius $R> 10^{-5}~{\\rm MeV}^{-1}$. Motivated by this fact, in this paper we present the effect of the neutrino electromagnetic properties on the difference between neutrino and its antineutrino mean free paths. To this end, we follow the same procedure as in Ref.~\\cite{Sul06}, i.e., the differential cross section is calculated in a linear response with zero temperature approximation, while the leptons are assumed to be Fermi gas. In modeling the interacting nucleons, the relativistic mean field model inspired by effective field theory \\cite{Furu96} has been used to describe the non strange dense matter with and without neutrino trapping. ",
        "conclusions": "We have calculated the neutrino and antineutrino mean free paths for all types of neutrinos and study their differences with various mechanisms, i.e., by using different types of matters (with and without neutrino trapping) as well as by including the weak magnetism of the nucleon and the electromagnetic form factors of the neutrinos. We have found that the difference in the neutrino and antineutrino mean free paths depends not only on the matter density, but also on the neutrino energy and the matter constituents. However, variation of the latter affects the mean free path difference in a distinct way. The effect of the neutrino electromagnetic form factors has been also studied. It is found that the corresponding contribution can be neglected as compared to the contribution from the nucleon weak magnetism."
    },
    "0608/astro-ph0608518_arXiv.txt": {
        "abstract": "{}{We present an analysis of the accreting X-ray pulsar system \\fu\\ based on \\inte\\ data. The main focus of this analysis is a study of the timing behavior of this source. In addition we also show an analysis of the 5--90\\,\\kev\\ spectrum.}  {The data were extracted using the official \\inte\\ software OSA~5.1. Timing analysis was performed using epoch folding and pulsar pulse phasing.}  {We have measured 12 individual pulse periods for the years 2003 to 2005. We confirm earlier \\xte\\ results that during 2003 the spin down became slower and show furthermore that after this phase \\fu\\ started to spin up with $\\dot{P}_\\text{pulse}=-0.158\\,\\text{s}\\,\\text{yr}^{-1}$ in 2004. The similarity of the pulse period histories of \\fu\\  and 4U\\,1626$-$26 suggests that accretion onto an oblique rotator, as recently proposed by Perna et al., is a possible explanation for this change.}{} ",
        "introduction": "\\label{sect:intro} The wind-accreting High Mass X-ray Binary system \\object{\\fu} \\citep{giacconi71} consists of a neutron star in an eccentric ($e=0.28$) 8.3753\\,d orbit \\citep{intzand98} around its companion, which has been identified optically with a highly reddened $m_\\text{V}=16.37$\\,mag star \\citep{schwartz80}.  Using interstellar atomic lines of \\ion{Na}{i} and \\ion{K}{i}, \\citet{cox05} set a lower limit of 5\\,kpc for the distance. According to this value a lower limit of the X-ray luminosity above 1\\,\\kev\\ is given by $2\\,10^{36}\\,\\text{erg}\\,\\text{s}^{-1}$ \\citep{intzand97}. \\citet{cox05} also confirm earlier suggestions \\citep{schwartz80,marshall80,kerkwijk89} that the stellar companion is a O8--O9 Ia supergiant with an effective temperature of 30500\\,K, a radius of 26\\,\\rsun, a luminosity of $5\\,10^{5}\\,\\lsun$, and a mass loss rate of $7\\,10^{-6}\\,\\msun\\,\\text{yr}^{-1}$. Note that the presence of X-ray flaring seen twice per neutron star orbit \\citep[][ see also Fig.~\\ref{fig:asm}]{marshall80} had led some authors \\citep[e.g.,][]{makishima84,iye86,cook87} to the suggestion of a Be star companion, however, this classification would require a distance of $<$1.5\\,kpc, which is in contradiction to the significant interstellar extinction measured in optical observations \\citep{kerkwijk89}.  Similar to other accreting neutron stars, the X-ray continuum of \\fu\\  can be described by a power-law spectrum with an exponential turnover at $\\sim$13\\,keV \\citep{mihara95,intzand97,cusumano98}.  The spectrum is modified by strong photoelectric absorption with a column $N_\\text{H}=1.5$--$5.7\\, 10^{22}\\,\\text{cm}^{-2}$ \\citep{schwartz80,marshall80,makishima84,cook87,chitnis93,cusumano98}.  As the absorbing material is situated in the dense stellar wind, $N_\\text{H}$ is strongly variable over the orbit.  The column is maximal between the end of the primary X-ray flare and the start of the secondary flare (Fig.~\\ref{fig:asm}) of the orbital light curve \\citep{roberts01}. Fluorescence in the absorbing material gives rise to an Fe K$\\alpha$ line at 6.4\\,keV with an equivalent width of $\\sim$60\\,eV \\citep{cusumano98}.  At higher energies, the spectrum exhibits cyclotron resonant scattering features (CRSF) at $\\sim$19\\,\\kev\\ and $\\sim$40\\,keV \\citep{makishima92,mihara95,cusumano98}.  With a pulse period of $\\sim440$\\,s, \\fu\\  is a slowly rotating neutron star.  Since the discovery of the pulsations by \\citet{makishima84} the neutron star has exhibited a steady linear spin down with an average of $\\dot{P}_\\text{pulse}=+0.225\\,\\text{s}\\,\\text{yr}^{-1}$ from $P_{\\text{pulse}}=437.5$\\,s in 1983 to 440.76\\,s in 1998 \\citep{intzand98}. Recently, \\citet{baykal05} reported a decrease in $\\dot{P}_\\text{pulse}$, which in 2002 was $\\sim$$0.115\\,\\text{s}\\,\\text{yr}^{-1}$ and therefore $\\sim$0.5 times the long term value.  In this paper we report on an analysis of \\inte\\ observations of \\fu\\ performed since 2003. In Sect.~\\ref{sec:datareduction} we describe the \\inte\\ data reduction, followed by a discussion of the broad-band X-ray spectrum of the source (Sect.~\\ref{sec:specevol}).  We then present the results of the long term X-ray timing analysis (Sect.~\\ref{sec:lc}), showing that the source underwent a torque reversal towards a spin up (Sect.~\\ref{sec:pulsevol}). We summarize and discuss our results in Sect.~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary} Since 2003, \\fu\\ was observed frequently with \\inte. In this paper, we presented the results of the timing analysis of $\\sim$2280\\,ksec of IBIS data and of the spectral analysis of 1258\\,ksec of IBIS and 200\\,ksec of JEM-X2 data.  Two cyclotron lines at $\\sim$19\\,\\kev\\ and $\\sim$40\\,\\kev\\ are detected, consistent with earlier \\ginga\\ and \\sax\\ results \\citep{mihara95,cusumano98}.  Assuming a gravitational redshift of $z=(1-(2GM/rc^{2}))^{-1/2}-1=0.31$ (assuming $M=1.4\\msun$ and $r=10^{6}\\,\\text{cm}$), from these CRSF observations a $B$-field strength of $2.15\\times10^{12}\\,\\text{G}$ can be inferred. Four flares in the X-ray lightcurves with durations between $\\sim$2500\\,s and $\\sim$7000\\,s were observed, three of which are associated with the primary flare in the orbital lightcurve of \\fu.  In addition to these confirmations of earlier results, our \\inte\\ data show, for the first time, a clear spin up phase in \\fu\\ (Table~\\ref{tab:periods} and in Figs.~\\ref{fig:perievo} and~\\ref{fig:inteperievo}).  After almost 20 years of constant spin down at a rate of $\\dot{P}_{\\text{pulse}}=+0.225\\,\\text{s}\\,\\text{yr}^{-1}$, with the available data being consistent with $\\ddot{P}=0\\,\\text{s}^{-1}$, in 2002 and 2003 first indications for a decrease in the magnitude of spin down were detected \\citep{baykal05}.  As shown in Sect.~\\ref{sec:pulsevol}, the new \\inte\\ data show a torque reversal from $\\sim$MJD\\,53131 onwards, with the source now exhibiting a spin up with a rate of $\\dot{P}_\\text{pulse}=-0.158\\,\\text{s}\\,\\text{yr}^{-1}$.  So far, the \\inte\\ results appear to be consistent with $\\ddot{P}=0\\,\\text{s}^{-1}$ (the $1\\sigma$ upper limit for $\\ddot{P}$ during the spin up is $-4\\,10^{-5}\\,\\text{s}^{-1}\\,\\text{yr}^{-2}$).  The traditional interpretation of spin up in X-ray binaries with a strongly magnetized, disk-accreting neutron star is that the accretion disk is truncated at the inner edge of the disk by the magnetic field of the neutron star \\citep[see, e.g.,][and references therein]{ghosh77,ghosh79a,ghosh79b}. For prograde disks, the magnetic coupling between the accretion disk and the neutron star and the transfer of angular momentum from the accreted matter onto the neutron star will lead to a torque onto the neutron star resulting in spin up. The general expectation of these models is that the accretion disk is close to equilibrium, i.e., the Alfv\\'en or magnetospheric radius, $r_\\text{m}$, where the accretion flow couples to the magnetic field, is assumed to be close to the corotation radius, $r_\\text{co}$, where the disk's Kepler frequency equals the rotational frequency of the neutron star. For accretion to occur, $r_\\text{m}\\lesssim r_\\text{co}$, as otherwise the onset of the centrifugal barrier (the ``propeller effect'') prevents the material from falling onto the neutron star \\citep{illarionov75}.  Most accreting neutron star systems, such as \\object{Her~X-1} or \\object{Vela X-1}, show a secular spin up or spin down trend onto which short, large magnitude episodes of spin down or spin up are superimposed \\citep{bildsten97}. These shorter episodes are often interpreted to be due to short term variations in $\\dot{M}$, which give rise to short-term torque fluctuations. In contrast to these systems, GX~1+4, 4U\\,1626$-$67, and \\fu\\ showed decade-long phases where $\\dot{P}$ did not change its sign.  \\object{GX~1+4} was seen to spin up with a very short characteristic timescale, $P/\\dot{P}\\sim 40$\\,years, from the early 1970s until the source flux dropped below detectability in the \\textsl{EXOSAT} era of the early 1980s. From 1985 onwards the source reemerged with a strong spin down \\citep{makishima88,mony91,chakrabarty97}. Recent \\inte\\ observations show that this spin down has continued until at least 2004 \\citep{ferrigno06}.  A second source with a long phase of unchanged sign of $\\dot{P}$ is the ultracompact low mass X-ray binary \\object{4U\\,1626$-$67}, which was spinning up between 1977 and 1989, followed by a torque reversal and an extended spin down phase starting in 1991 \\citep{chakrabarty97b}.  In contrast to GX~1+4, no luminosity drop was seen during the torque reversal and with $|P/\\dot{P}|\\sim 5000$\\,years the characteristic spin up and spin down timescales are significantly larger.  Provided that the torque reversal of \\fu\\ described in this paper does not change back in the next years, in its overall period behavior, \\fu\\ resembles 4U\\,1626$-$67: both sources have $|P/\\dot{P}| \\sim$ several thousands of years and no change in luminosity was observed during the torque reversal. On the other hand, however, \\fu\\ is a high mass X-ray binary with a much longer orbital period, and thus shares similarities with the low mass X-ray binary GX~1+4, and furthermore \\fu\\ switched from spin down to spin up, while GX~1+4 and 4U~1626$-$67 switched from spin up to spin down.  Common for the three sources is that the magnitude of $|P/\\dot{P}|$ is similar for the spin up and the spin down.  With the simple magnetic torquing model outlined above, the long distinct episodes of roughly constant $\\dot{P}$ and the magnitude and sign of $\\dot{P}$ are difficult to explain: As discussed above, for systems close to equilibrium it is expected that $r_\\text{m}\\sim r_\\text{co}$. In reality, however, this is not the case.  In the case of \\fu, it was first pointed out by \\citet{intzand98} that the magnetospheric radius inferred from the cyclotron line measurements is $r_\\text{m}\\sim 2400$\\,km, while $r_\\text{co}\\sim12000$\\,km. Moving $r_\\text{m}$ out to $r_\\text{co}$ would require a magnetic field of $\\sim$$10^{14}\\,\\text{G}$, which is two orders of magnitude larger than the magnetic field deduced from the observed cyclotron lines. Alternatively, $\\dot{M}$ could be significantly lower than the value inferred from the X-ray luminosity of the system. Assuming the efficiency of accretion is comparable to other systems, however, a low $\\dot{M}$ would imply a significantly smaller distance to \\fu\\ than allowed by its optical reddening, which seems equally unlikely. \\citet{intzand98} argued that the presence of quasi-periodic oscillations (QPOs) and the long term systematic trend in the spin down are strong indicators for the presence of a accretion disk in the system with a small inner disk radius.  In analogy to GX~1+4 \\citep[][ and references therein]{chakrabarty97}, these authors show that a transient retrograde disk with a duty cycle of 1--5\\% could provide a sufficiently strong torque to explain the observed spin down.  This torquing would coincide with a short term increase of source luminosity. \\citet{intzand98} speculate that the X-ray flares indicate the presence of the retrograde disk, since a switch to a retrograde disk would imply an increase in X-ray luminosity. They point out, however, that the $\\sim$1000\\,s duration of the flares seen in their \\textsl{RXTE} data is too small with respect to the required duty cycle for a retrograde disk.  We note that the \\inte\\ results show flare durations of several ksec duration (Tab.~\\ref{tab:flares}), which is more in line with the expected duty cycle. Furthermore, the folded lightcurve of \\fu\\ indicates a brightening of the source around phase 0.6 (Fig.~\\ref{fig:asm}). With $0.1P_\\text{orb}$ duration, this flare is long enough to easily accommodate the required duty cycle of the retrograde disk and to explain the observed $\\dot{P}$. We note, however, that not all of the brightening seen in Fig.~\\ref{fig:asm} is due to the flares identified by us as removing these flares from the data still results in a clear peak at phase 0.6 in the folded lightcurve. Part of the brightening could therefore be also due to an enhanced $\\dot{M}$ during this phase of the elliptical orbit of the neutron star, although we cannot firmly exclude the presence of smaller flares which would not be picked up due to our conservative definition of a flare.  While the model of a retrograde disk can explain the magnitude of $P/\\dot{P}$, it is more difficult to reconcile the torque reversal found with \\inte\\ with this model. First of all, as discussed above, the four intervals of different $\\dot{P}$ represent distinct episodes of different torques on the neutron star.  In the retrograde disk model, it is difficult to understand why the characteristic spin up and spin down timescales appear so similar. Furthermore, these three phases would then correspond to different duty cycles of the retrograde disk, resulting in a change in observed flux. Such a behavior has been observed, e.g., in GX~1+4, where torque and luminosity are correlated during the spin down phase \\citep{chakrabarty97}. No such correlations are observable in \\fu, where the flux has not appreciably changed during the \\inte\\ and the pointed \\xte\\ observations \\citep[see also][]{baykal05}. In addition, an analysis of the orbit-averaged 2--12\\,keV \\xte\\ \\asm\\ lightcurves also does not reveal a change in the soft X-ray behavior of the source, and neither is a change in the X-ray light curve with orbital phase observed.  The \\inte\\ observations also rule out the application of the torque reversals put forward by \\citet{murray99} to explain GX~1+4. In this model, the accretion disk is assumed to consist of several rings with opposite rotation.  Between these rings a gap is created such that a torque reversal should be accompanied by a minimum in the accretion luminosity.  The model predicts a change in pulse profile from ``leading-edge bright'' during the spin down to ``trailing-edge bright'' during the spin up. Both, a minimum in luminosity and a change in the shape of the pulse profile are observed in GX 1+4 \\citep{greenhill99}, but neither effect is present in \\fu\\ (see discussion above and Fig.~\\ref{fig:pp_rev}).  Recently, \\citet{perna06} presented a new \\textit{Ansatz} to explain torque changes in accreting neutron stars, which does not require the presence of retrograde disks. This model makes use of the fact that X-ray pulsars are oblique rotators, i.e., the neutron star's magnetic field is tilted by an angle $\\chi$ with respect to the axis of rotation of the neutron star. Assuming the accretion disk is situated in the neutron star's equatorial plane, for a ring of matter in the accretion disk, the magnetic field strength then depends on the azimuthal angle and thus the boundary of the magnetosphere is asymmetric.  As shown by \\citet{perna06}, such a configuration can lead to regions in the disk where the propeller effect is locally at work, while accretion from other regions is not inhibited (i.e., on a ring on the disk, there are some regions with $r_\\text{m}<r_\\text{co}$ and other regions with $r_\\text{m}>r_\\text{co}$). This results in a nonlinear dependence of the accretion luminosity from the $\\dot{M}$ through the outer parts of the disk.  A nontrivial consequence of this more realistic accretion geometry is that for values of $\\chi$ greater than a critical value, $\\chi_\\text{crit}$, limit cycles are present, where cyclic torque reversal episodes are possible without a change in $\\dot{M}$. For typical parameters. $\\chi_\\text{crit}$ is between $\\sim$$25^\\circ$ and $\\sim$$45^\\circ$.  We stress that the existence of the limit cycles depends only on $\\chi$ and on the pulsar's polar magnetic field, $B$, and no variation of external parameters is required to trigger torque reversals.  This fact is a big advantage of this model over the models discussed above.  \\citet{perna06} show that their model can explain the principal observed properties of GX~1+4, including the luminosity drop during torque reversal, except for the observed correlation between torque and luminosity during the spin down phase. For 4U\\,1626$-$67, assuming $B=2.5\\times 10^{12}$\\,G and $\\chi=68^\\circ$, all observed properties of the spin history of 4U\\,1626$-$67 including the large and values of $|P/\\dot{P}|$ before and after the torque reversal and the virtually unchanged luminosity of the source can be explained (in the model, the required change in luminosity is only $\\sim$5\\%). For these model parameters, the limit cycle of 4U\\,1626$-$67 is predicted to have a long period.  The similarity between the pulse histories of \\fu\\ and 4U\\,1626$-$67 suggests that a similar model would also work for the former source. We stress, however, that further quantitative work, which is outside the scope of this paper, is clearly required. Specifically, from a theoretical point of view we note that the model of \\citet{perna06} does not yet take changes in $\\dot{M}$ due to the orbital eccentricity into account and that the current version of the model assumes the disk to be flat and in the equatorial plane of the neutron star, while more realistic disks onto magnetized objects can be expected to be warped and precessing \\citep[e.g.,][and therein]{pfeiffer04}. From an observational point of view, the donor stars in both systems are very different, with 4U\\,1626$-$67 being a low mass and \\fu\\ being a high mass system.  Finally, we also emphasize that the torque reversal in 4U\\,1626$-$67 was from a spin up to a spin down, while for \\fu\\ it was from a spin down to a spin up. Since the torque reversal episodes are cyclic, however, we do not expect this latter difference to be of major importance.  In conclusion, the torque reversal of \\fu\\ with no associated drop in luminosity and no change in the shape of the pulse profile seem to be difficult to reconcile with models explaining the large $P/\\dot{P}$ through the presence of a retrograde disk. On the other hand, the proposal of \\citet{perna06} to explain torque reversals through accretion onto an oblique rotator seems to be able to explain the long term constancy of $\\dot{P}$ trends in \\fu\\ and 4U\\,1626$-$67, the large magnitude of $|P/\\dot{P}|$, and the constancy of X-ray flux over the torque reversal, although further and more detailed theoretical studies are required. As a prediction of the model for systems such as \\fu\\ and 4U\\,1626$-$67 is that torque reversal episodes are rare events, further monitoring of the pulse period of \\fu\\ is required to determine whether this change is a sign of a long term torque reversal or whether we are just observing a short deviation from a long term spin down trend."
    },
    "0608/astro-ph0608204_arXiv.txt": {
        "abstract": "We apply time-distance helioseismology, local correlation tracking and Fourier spatial-temporal filtering methods to realistic supergranule scale simulations of solar convection and compare the results with high-resolution observations from the SOHO Michelson Doppler Imager (MDI).  Our objective is to investigate the surface and sub-surface convective structures and test helioseismic measurements.  The size and grid of the computational domain are sufficient to resolve various convective scales from granulation to supergranulation. The spatial velocity spectrum is approximately a power law for scales larger than granules, with a continuous decrease in velocity amplitude with increasing size. Aside from granulation no special scales exist, although a small enhancement in power at supergranulation scales can be seen. We calculate the time-distance diagram for $f$- and $p$-modes and show that it is consistent with the SOHO/MDI observations. From the simulation data we calculate travel time maps for surface gravity waves ($f$-mode).  We also apply correlation tracking to the simulated vertical velocity in the photosphere to calculate the corresponding horizontal flows.  We compare both of these to the actual large-scale (filtered) simulation velocities. All three methods reveal similar large scale convective patterns and provide an initial test of time-distance methods. ",
        "introduction": "\\label{intro} Local helioseismology uses the observed oscillation signals to probe the near surface structure of the Sun to determine sound speed, flow velocities and magnetic field structures. One of several methods used in this field is called the time-distance helioseismology. It measures the time taken by a wave packet to travel from one point on the solar surface to another \\citep{1993Natur.362..430D, 1997scor.proc..241K}. The travel time for a wave depends on the wave speed and the flow velocities along the ray paths. These effects can be separated by measuring the wave travel time for waves propagating in opposite directions along the same ray paths. Inversion methods applied to infer the properties of subsurface convection and structures are based on various approximations, such as a geometric acoustic ray approximation \\citep{2000SoPh..192..159K, 2001ApJ...557..384Z}, Fresnel-zone approximation \\citep{2001ApJ...553L.193J, 2004ApJ...607..554C} and Born-approximation \\citep{2002ApJ...571..966G}. Unfortunately, until now, there has been no direct verification of these methods \\citep{2003soho...12..319J,2004soho...14..172W}.  Realistic 3D simulations of solar convection \\citep{{2000SoPh..192...91S},{Benson2006}} provide a way to evaluate the accuracy and consistency of time-distance and other local helioseismology techniques. These simulations provide a model of convective motions in the upper convection zone and possesses a rich spectrum of $f$- and $p$-mode oscillations excited by the turbulent convection.  They can be used for testing helioseismology methods. Therefore, it is very important to apply some of the existing local helioseismology methods to the simulated solar convection, and check whether the results are consistent with the subsurface properties present in the simulated flows. In addition, the comparison of the oscillation properties obtained from the numerical model, such as the  velocity spectrum, $k-\\omega$ diagram and time-distance (cross-covariance) diagram, with observations provides validation for the simulation model.  Previously, the realistic simulations of solar convection have been used to study the excitation mechanism of solar and stellar oscillations \\citep{2004SoPh..220..229S}, and some properties of oscillation modes, such as the line asymmetry \\citep{2000ApJ...530L.139G}. However, these simulations were carried out for small, granulation-size, computational domains. Recently, with the progress of parallel supercomputing it became possible to substantially expand the computational domain, both horizontally and in depth, and to simulate the multi-scale solar convection, from granulation to supergranular scales \\citep{Benson2006}. These simulations allow us for the first time to investigate the properties of large-scale convection by applying helioseismology and other methods used to analyze solar observations, and thus test these methods and also evaluate how close the realistic simulations are to the real Sun. Such tests are particularly important for the local helioseismology methods, which are based on simplified models of wave propagation.  The goal of this paper is to investigate the basic helioseismic properties of the large-scale 3D simulations (the oscillation power spectrum and wave travel times) and compare these with the high-resolution observations from SOHO/MDI \\citep{1995SoPh..162..129S}.  An initial time-distance analysis is carried out for surface gravity waves ($f$-mode), which are well suited for studying the surface and sub-surface structure of solar convection on supergranulation scales \\citep{2000SoPh..192..177D}. Our analysis shows that travel times from the simulation agree very well with travel times from the SOHO/MDI observations. This demonstrates that the simulated data are sufficiently close to the observed solar data, providing ground for further detailed helioseismic measurements and inversions. Large-scale structures can be detected in the simulated data set using $f$-mode time-distance and local correlation tracking techniques and are observed directly in filtered flow fields. ",
        "conclusions": "\\label{conc}  The new large-scale realistic 3D radiative-hydrodynamic simulations of the upper layers of solar convection provide an excellent opportunity to validate various techniques, widely used in solar physics and helioseismology for directly obtaining otherwise inaccessible properties (subsurface flows and structures, etc). On the other hand, these analysis techniques also help to examine how realistic the simulations are. We have performed an initial analysis of the simulated data and compared our results with the outcome of the SOHO/MDI observations. The similarity between the simulated and observed $\\ell-\\nu$ and time-distance diagrams demonstrates that the simulations can be efficiently used to perform and validate local helioseismology techniques. This agreement also reveals a potential for high-resolution time-distance measurements. We carried out $f$-mode time-distance calculations and local correlation tracking to obtain horizontal velocities. We compare them to the simulated horizontal flows and find a good qualitative agreement. We see large-scale structures in both actual horizontal velocities and their proxies. The results of this investigation provide the basis for further detailed helioseismic analyses of the simulated data, including various scheme of measuring p-mode travel times, approximations for the sensitivity kernels and inversion procedures.  {\\bf Acknowledgements.}  DG is grateful for the support of the SOHO/MDI group at Stanford University.  DB and RFS were supported by NASA grants NNG04GB92G and NAG 512450 and NSF grants AST-0205500 and AST-0605738.  The work of {\\AA}N was supported by a grant from the Danish Natural Science Research Council.  Computing time was provided by the Danish Center for Scientific Computing, and by grants from the NASA Advanced Supercomputing Division."
    },
    "0608/astro-ph0608174_arXiv.txt": {
        "abstract": "\\noindent We present a method for accelerating the calculation of CMB power spectra, matter power spectra and likelihood functions for use in cosmological parameter estimation. The algorithm, called {\\sc CosmoNet}, is based on training a multilayer perceptron neural network and shares all the advantages of the recently released {\\sc Pico} algorithm of Fendt \\& Wandelt, but has several additional benefits in terms of simplicity, computational speed, memory requirements and ease of training.  We demonstrate the capabilities of {\\sc CosmoNet} by computing CMB power spectra over a box in the parameter space of flat $\\Lambda$CDM models containing the $3\\sigma$ WMAP1 confidence region. We also use {\\sc CosmoNet} to compute the WMAP3 likelihood for flat $\\Lambda$CDM models and show that marginalised posteriors on parameters derived are very similar to those obtained using {\\sc CAMB} and the WMAP3 code. We find that the average error in the power spectra is typically $2-3\\%$ of cosmic variance, and that {\\sc CosmoNet} is $\\sim 7 \\times 10^4$ faster than {\\sc CAMB} (for flat models) and $\\sim 6 \\times 10^6$ times faster than the official WMAP3 likelihood code. {\\sc CosmoNet} and an interface to {\\sc CosmoMC} are publically available at {\\tt www.mrao.cam.ac.uk/software/cosmonet}. ",
        "introduction": "\\label{sec:intro}   In the analysis of increasingly high-precision data sets, it is now common practice in cosmology to constrain cosmological parameters using sampling based methods, most notably Markov chain Monte Carlo (MCMC) techniques (Christensen et al. 2001; Knox, Christensen \\& Skordis 2001; Lewis \\& Bridle 2002). This approach typically requires one to calculate theoretical CMB power spectra (i.e. some subset of the TT, TE, EE and BB $C_\\ell$ spectra) and/or the matter power spectrum $P(k)$ at a large number of points (typically $\\sim 10^5$ or more) in the cosmological parameter space. In addition, one must also evaluate at each point the corresponding (combined) likelihood function for the data set(s) under consideration. As a result, the process can be computational very demanding.  The purist would calculate the required power spectra at each point using codes such as {\\sc CMBfast} (Seljak \\& Zaldarriaga 1996) or {\\sc CAMB} (Lewis, Challinor \\& Lasenby 2000), which typically require around 10 secs for spatially-flat models and 50 secs for non-flat models. This approach is therefore computationally demanding, but does have the advantage that it is simple to generalise if one wishes to include new physics or change the form of the initial power spectra. MCMC parameter estimation codes such as {\\sc CosmoMC} (Lewis \\& Bridle 2002) attempt to decrease the overall computational burden by dividing the cosmological parameter space into `fast' parameters (governing the initial primoridal power spectra of scalar and tensor perturbations) and `slow' parameters (governing the perturbation evolution) and making judicious proposals for how the chain is propagated in parameter space.  Even with this technique, however, the total computational cost is still usually very high.  If one is willing to forego the full calculation of the required power spectra at each point in parameter space, there are a number of ways in which suitably accurate spectra can be generated somewhat more rapidly. If the cosmological parameter space of interest is sufficiently small, then it is possible simply to create spectra for a regular grid of models in parameter space and interpolate between them in some way. As the number of parameters increases, however, the computational cost of constructing the grid grows exponentially. Fast grid generation schemes have been proposed, such as the $\\ell$-splitting scheme of Tegmark \\& Zaldarriaga (2000) that exploits analytic approximations at high-$\\ell$ and insensitivity to certain parameters at low-$\\ell$. Nevertheless, the pre-compute of the grid of models remains extremely time-consuming and such approaches become difficult to implement accurately when second-order effects such as gravitational lensing are important.  More extensive use of analytic and semi-analytic approximations can reduce the required number of pre-computed models, but only at the cost of a loss of accuracy and/or placing restrictions on the parameters that are available as input.  Such approaches are usually based on a relatively sparse grid of base models in the parameter space from which the spectra of more general models are computed rapidly on-the-fly using various (semi-)analytic approximations.  The {\\sc DASh} code of Kaplinghat, Knox \\& Skordis (2002), instead stores a sparse grid of transfer functions (rather than $C_\\ell$), uses efficient choices for grid parameters and makes considerable use of analytic approximations. Following $\\sim 40$ hrs of computation on a typical desktop to calculate the grid, DASh provides a speed-up factor of $\\sim 30$ relative to CMBfast in calculating a $C^{\\rm TT}_\\ell$, $C^{\\rm TE}_\\ell$ or $C^{\\rm EE}_\\ell$ spectrum.  More recently, the need to pre-compute a grid of models has been removed in the {\\sc CMBwarp} package (Jimenez et al. 2004), which builds on the method introduced by Kosowsky, Milosavljevic \\& Jimenez (2002).  In this approach, a new set of nearly uncorrelated `physical parameters' are introduced upon which the CMB power spectra have a simple dependence. {\\sc CMBwarp} uses a modified polynomial fit in these parameters in which the coefficients are based on the spectra $C^{\\rm TT}_\\ell$, $C^{\\rm TE}_\\ell$ or $C^{\\rm EE}_\\ell$ for just a single fiducial model in the parameter space. Spectra for other models can then be calculated around $\\sim 3000$ times faster than {\\sc CMBfast}.  By taking the fiducial model to be the best-fit model to the WMAP1 data, {\\sc CMBwarp} gives better than 0.5 per cent accuracy for the $C^{\\rm TT}_\\ell$ spectrum throughout the entire region of parameter space lying within the WMAP1 3$\\sigma$ confidence region, although the accuracy quickly reduces as one moves further away from the fiducial model.   Although the above methods have proved extremely useful in performing cosmological parameter estimation, they do exhibit a number of drawbacks, as we have outlined. Most recently, this has led Fendt \\& Wandelt (2006) to propose a more flexible and robust machine-learning approach (called {\\sc Pico}) to accelerating both power spectra and likelihood evaluations.  In this method, one first calculates the required spectra (usually $C^{\\rm TT}_\\ell$, $C^{\\rm TE}_\\ell$ or $C^{\\rm EE}_\\ell$) using {\\sc CAMB} and the corresponding likelihoods for the experiments of interest (in particular WMAP3) at $\\sim 10^4$ points chosen uniformly within a box in parameter space that encompasses (say) the $3\\sigma$ confidence region of the WMAP3 likelihood. This constitutes the training set for the {\\sc Pico} code -- note that only power spectra values at the limited number of $\\ell$-values output by {\\sc CAMB} are used (typically $50$ values for $\\ell_{\\rm max}=1500$). In short, the basic algorithm used by {\\sc Pico} consists of three major parts. First, the training set is compressed using Karhunen--Lo\\`eve eigenmodes (essentially a principal component analysis) which typically results in a reduction in the dimensionality of the training set by a factor of two. Second, the training set is used to divide the parameter space into ($\\sim 100$) smaller regions using a $k$-means clustering algorithm (see e.g. MacKay 1997) with the goal that all clusters encompass volume of parameter space over which the power spectra vary roughly equally. Finally, a (4th order) polynomial is fitted within each cluster (by minimising the squared error) to provide a local interpolation of the power spectra within the cluster as a function of cosmological parameters.  The reason for dividing up the parameter space in the second step is that the interpolation method used fails to model accurately the power spectra over the entire parameter space.  The {\\sc Pico} approach provides about the same speed-up in spectrum calculation as {\\sc CMBWarp} (which is an order of magnitude faster than {\\sc DASh}), but is an order of magnitude more accurate. It also has several other important advantages. First, it is very flexible and can easily be applied to the fast calculation of any observables relevant to a particular data set, such as scalar, tensor and lensed power spectra, transfer functions or even higher-order correlation functions. Second, it allows the calculation of such observables from an arbitrary number of cosmological models and in any range of $\\ell$ (or $k$) values. Lastly, the algorithm is sufficiently generic to allow the direct fitting of likelihood functions, thereby incorporating the functionality of the {\\sc CMBfit} code of Sandvik et al.  This last capability allows an additional order of magnitude speed-up in cosmological parameter estimation beyond that resulting from faster power spectrum calculations, and is particularly important for experiments such as WMAP3 for which the likelihood calculation is very expensive.  In this letter, we present an independent approach to using machine-learning techniques for accelerating both power spectra and likelihood evaluations. Our approach is based on training a neural network in the form of a 3-layer perceptron. The resulting {\\sc CosmoNet} code shares all the advantages of {\\sc Pico}, but we believe also has some additional benefits in terms of simplicity, computational speed, accuracy, memory requirements and ease of training. The letter is organised as follows. In Section~\\ref{sec:nn} we give a brief introduction to neural networks and our training algorithm. The resulting network output is discussed in Section~\\ref{sec:interp}, where we investigate the accuracy of our approach.  The {\\sc CosmoNet} code is then used to perform a cosmological parameter estimation from WMAP3 data in Section~\\ref{sec:cosmo}. Our conclusions are presented in Section~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have presented a method for accelerating power spectrum and likelihood evaluations based on the training of multilayer perceptron neural networks, which we have shown to be fast, robust and accurate. Our {\\sc CosmoNet} method shares all the advantages of the Pico algorithm of Fendt \\& Wandelt, achieving similar accuracies on both spectra interpolations and cosmological parameter constraints, but there are several differences between the two methods that we believe give {\\sc CosmoNet} a number of additional benefits, which we now discuss.  {\\em Simplicity}.  Despite requiring the optimisation of a highly non-linear multi-dimensional function using MemSys, we consider the principal advantage of our method to be the relative simplicity of the trained interpolation for the user to implement. {\\sc CosmoNet} provides a {\\em single} simple, closed-form function for each interpolation over the {\\em whole} of the parameter space under consideration.    {\\em Memory usage}. A neural network with $N_{\\rm in}$ inputs, $N_{\\rm hid}$ nodes in the hidden layer, and $N_{\\rm out}$ outputs has $(N_{\\rm in} + 1)N_{\\rm hid} + (N_{\\rm hid} + 1)N_{\\rm out} \\approx N_{\\rm hid}N_{\\rm out}$ parameters. This is far less than in the {\\sc Pico} approach, where the use of clustering and individual interpolations for each $C_\\ell$ requires far more parameters. In the case of the flat $\\Lambda$CDM example demonstrated in section \\ref{sec:interp}, we require about $100$ kB of parameter memory for all three power spectra and the likelihood, whereas {\\sc Pico} would require $15$ MB. While the memory requirements of {\\sc Pico} will increase with the number of cosmological parameters, this should make little difference to the memory requirements of our method, as we have found the required number of nodes in the hidden layer does not increase beyond a factor of 2 for the 11 parameter non-flat model parameterised by $\\Omega_b h^2$, $\\Omega_c h^2$, $\\Omega_k$, $\\theta$, $\\tau$, massive neutrino fraction $f_{\\nu}$, varying equation of state of dark energy $w$, scalar perturbation amplitude and spectral index $A_s$, $n_s$ and tensor modes with amplitude ratio and spectral index $R$, $n_t$ (results to appear in a forthcoming paper), thus representing a linear rise.    {\\em Speed}. The number of calculations required to perform the feed-forward network mapping is $2 N_{\\rm in} \\ N_{\\rm hid} \\ + \\ 2 N_{\\rm hid} \\ N_{\\rm out} \\approx 2 N_{\\rm hid} N_{\\rm out} \\ $. In the example presented in section \\ref{sec:interp} the calculation of the 50 uninterpolated $C_\\ell$ values for each spectrum required $\\sim$ 120 microseconds, whereas each WMAP3 likelihood took $\\sim$ 10 microseconds; this is $\\sim 25$ times faster than {\\sc Pico} in performing the interpolation. Moreover, the CPU requirements of the {\\sc Pico} interpolation scales as $( \\frac{N_{\\rm in}}{p})^p$, for {\\sc Pico} $p$ is 4.  {\\em Ease of training}. Training using the {\\sc MemSys} package is almost totally automated and relatively quick. In fact the bottleneck in providing appropriate network weights for more complex cosmological models is the calculation of the training data using the {\\sc CAMB} package. The networks used in section \\ref{sec:interp} took around 100 hours of training (on a standard PC workstation). Additionally, MemSys training time scales linearly with the number of network nodes and again linearly with the number of entries in the training set. However, networks with an accuracy of roughly $2-3$ times worse, that are sufficient to provide good parameter constraints can be trained in under an hour using just 1000 training samples. These simpler neural networks did not require the l-splitting performed in section \\ref{sec:interp} and had only 50 nodes in the hidden layer.  \\medskip      \\subsection*"
    },
    "0608/astro-ph0608497_arXiv.txt": {
        "abstract": "NGC 5128, a giant elliptical galaxy only $\\sim 4$ Mpc away, is the dominant member of a galaxy group of over 80 probable members.  The Centaurus group provides an excellent sample for a kinematic comparison between the halo of NGC 5128 and its surrounding satellite galaxies.   A new study, presented here, shows no kinematic difference in rotation amplitude, rotation axis, and velocity dispersion between the halo of NGC 5128, determined from over $\\sim340$ of its globular clusters, and those of the Centaurus group as a whole. These results suggest NGC 5128 could be behaving in part as the inner component to the galaxy group, and could have begun as a large initial seed galaxy, gradually built up by minor mergers and satellite accretions, consistent with simple cold dark matter models.   The mass and mass-to-light ratios in the B-band, corrected for projection effects, are determined to be $(1.3\\pm0.5) \\times 10^{12}$ M$_{\\sun}$ and $52\\pm22$ M$_{\\sun}$/L$_{\\sun}$ for NGC 5128 out to a galactocentric radius of 45 kpc, and $(9.2\\pm3.0) \\times 10^{12}$ M$_{\\sun}$ and $153\\pm50$ M$_{\\sun}$/L$_{\\sun}$ for the Centaurus group, consistent with previous studies. ",
        "introduction": "The study of galaxy groups can provide essential insight into galaxy formation, particularly because many groups are much less dynamically evolved than rich clusters and thus are closer to their earliest conditions.  The density of formation environment affects the dynamical evolution of galaxies; for example, the luminosity and color distributions of galaxies differ in field versus group environments \\citep{girardi03}.  Furthermore, for groups with large central members, the properties of the central galaxy correlate with properties of their satellite galaxies, such as the satellite-type fraction found in the group \\citep{weinmann06}, the alignment of satellite galaxies preferentially along the dominant galaxy's major axis \\citep{yang06,brainerd05,zentner05}, and the rotation of satellite galaxies in the same direction as the host halo \\citep{warnick05}.  The connections between the centrally dominant galaxy in a group environment and its surrounding satellite galaxies raise the question whether there might be a more direct kinematic connection between them as well.  Kinematic similarities between the group and its brightest group member would be consistent with both cold dark matter models, suggesting the central giant and its satellite group are part of the same large-scale structure, and the hierarchical formation scenario, whereby the central giant would build up by accretion of satellites in the group.  NGC 5128 is an E/S0 giant galaxy in the positional and dynamical center of the Centaurus group of galaxies.  Its remarkably close proximity (just 4 Mpc distant) also provides a genuinely rare opportunity to study and compare the dynamics of the halo of NGC 5128 with the dynamics of the surrounding group at a level of detail that is difficult elsewhere.  The Centaurus group consists of 25 confirmed galaxy members \\citep{kar06}, 13 of which have radial velocity measurements used in this kinematic study; see Figure~\\ref{fig:galpos}. Within a projected distance of $\\sim1$ Mpc from NGC 5128 lies NGC 5236 (M83).  NGC 5236 is a dominant spiral galaxy surrounded by 9 other confirmed members within the M83 group.  Surrounding these two systems are 53 smaller galaxies that do not have confirmed membership in either the Centaurus group or the M83 group.  Therefore the bounds separating these two groups are still quite uncertain, and it has been suggested that these two systems create a dumbbell system, similar to our own Milky Way and M31 \\citep{kar06}.  Only one other kinematic comparison between the halo of a dominant galaxy and its surrounding satellite galaxies (M87 in the Virgo cluster) has been presented in the literature \\citep{cote01}.  Virgo is a much more dynamically evolved environment than the Centaurus group.  The Centaurus group, therefore, provides the chance to carry the comparison a step further than just alignments of satellites and simple number statistics by looking at an environment closer to its initial structure. ",
        "conclusions": "\\label{discussion}  Investigating the kinematics of the GCs in the giant elliptical galaxy, NGC 5128, and the kinematics of over 60 satellite galaxies in and surrounding the Centaurus group, has shown interesting new results.  The rotation amplitude, rotation axis, and velocity dispersion of both the halo of NGC 5128 and surrounding galaxies both have means of $\\Omega R = 67\\pm27$ km s$^{-1}$, $\\Theta_o = 162\\pm22^o$ E of N, and $\\sigma_v = 103\\pm14$ km s$^{-1}$. The masses determined with the Tracer Mass Estimator and the spherical component of the Jeans Equation, lead to projection corrected masses of $(1.3\\pm0.5) \\times 10^{12}$ M$_{\\sun}$ for NGC 5128 and  $(9.2\\pm3.0) \\times 10^{12}$ M$_{\\sun}$ for the entire Centaurus group.  Mass-to-light ratios in the B-magnitude were subsequently determined to be $52\\pm22$ M$_{\\sun}$/L$_{\\sun}$ for NGC 5128 and $153\\pm50$ M$_{\\sun}$/L$_{\\sun}$ for the Centaurus group.  The similar kinematics between the halo of the brightest group member with its surrounding satellites suggests that the halo of NGC 5128 could therefore be thought of as the innermost component of the entire group that surrounds it.  The Centaurus group is likely bound, although its outer regions are unlikely to be virialized beyond $R\\simeq1.5$ Mpc.  The global dynamics of the group that we are seeing, could therefore represent the initial internal shear of the pregalactic filamentary structure from which it condensed.   NGC 5128 could have formed early on as a galaxy that was already centrally dominant, and now continues to build up by a series of minor mergers and accretions, leading to an averaging of the kinematics from the surrounding satellite galaxies that we now measure in the GCs in the NGC 5128 halo.  A single or small number of major merger events would not necessarily produce these kinematic similarities.  This view is consistent with the recent work on the halo stars in NGC 5128 which have a mean age of $8^{+3}_{-3.5}$ Gyr \\citep{rejkuba05}, and the age distribution of the GCs which shows that a high fraction of both metal-rich and metal-poor GCs are old (M. Beasley et al. 2006, in preparation).  It would be interesting to extend this type of study to other nearby groups with centrally dominant giants, although large samples of GCs and satellite galaxies with measured radial velocities are required in each case."
    },
    "0608/astro-ph0608342_arXiv.txt": {
        "abstract": "{ We investigate the evolution and the statistical properties of the {\\Lya} absorbers of the intergalactic medium (IGM) in the largely unexplored redshift range $z=0.5-1.9$. We use high-resolution ($R \\geq 30\\,000$) UV (STIS) and optical (VLT/UVES and Keck/HIRES) spectra of nine bright quasars with $z_{\\rm em} < 1.94$. The {\\Lya} lines detected in the lines of sight (LOS) towards these quasars are evaluated with a software package which determines simultaneously the quasar continuum and the line profiles. The main results for the combined {\\Lya} line sample are summarized as follows:  \\begin{enumerate}  \\item{The evolution of the number density of the absorbers can be described by the power law $\\frac{{\\rm d}n}{{\\rm d}z} \\propto (1+z)^{\\gamma}$. The number density of the low column density lines ($N_{\\rm \\ion{H}{i}}=(10^{12.90}-10^{14.00})$ cm$^{-2}$) decreases with decreasing $z$ with $\\gamma=0.74\\pm0.31$ in the interval $z=0.7-1.9$. \\,\\, A comparison with results at higher redshifts shows that it is decelerated in the explored redshift range and turns into a flat evolution for $z \\rightarrow 0$. The stronger absorbers (\\mbox{$N_{\\rm \\ion{H}{i}} > 10^{13.64}$ cm$^{-2}$}) thin out faster ($\\gamma=1.50\\pm0.45$). The break in their evolution predicted for $z=1.5-1.7$ cannot be seen down to $z=0.7$. On the other hand, a comparison with values from the literature for the local number density gives a hint that this break occurs at lower redshift.}  \\item{The distribution of the column densities of the absorbers is complete down to \\mbox{$N_{\\rm HI}=10^{12.90}$ cm$^{-2}$}. It can be approximated by a single power law with the exponent $\\beta=1.60\\pm0.03$ over almost three orders of magnitude. $\\beta$ is redshift independent.}  \\item{The {\\Lya} lines with lower column densities as well as the higher column density lines show marginal clustering with a $2\\sigma$ significance over short distances ($\\Delta v < 200$ {\\kms} and $\\Delta v < 100$ {\\kms}, respectively). We do not see any difference in the clustering with either column density or redshift.}  \\item{The distribution of the Doppler parameters has a mean value of $\\overline{b}=(34\\pm22)$ {\\kms}. This value is typical for the analyzed region. It does not change significantly with $z$.}   \\end{enumerate} ",
        "introduction": "It is traditional to decompose spectra of the {\\Lya} forest into individual absorbers that can be fit with Voigt profiles, each with a redshift $z$, a column density of neutral hydrogen $N_{\\rm \\ion{H}{i}}$ and a Doppler parameter $b$. Examining a sufficiently large number of LOS to background quasi-stellar objects (QSOs), one can derive the distributions of these parameters as well as dependences between them and in this way deduce characteristics of the {\\Lya} forest. Observers have often focused on the evolution of the number density of {\\Lya} lines, usually approximated by the power law  \\begin{equation} \\label{dn/dz} \\frac{{\\rm d}n}{{\\rm d}z} = \\left(\\frac{{\\rm d}n}{{\\rm d}z}\\right)_{0} (1+z)^{\\gamma} \\end{equation} \\citep{Sargent}. At high redshift ($z>1.5$), the evolution of the strong lines (log~$N_{\\rm \\ion{H}{i}}>14.00$) is steep ($\\gamma=2-3$), transiting into a flat evolution ($\\gamma=0.1-0.3$) at lower redshift. This abrupt slow-down in the evolution, first detected by observers \\citep{Bahcall, Impey1996, Weymann} and later also seen in simulations \\citep{Theuns1998, Dave1999}, is suspected to happen at \\mbox{$z=1.5-1.7$}.  However, since the apparent break in the evolution rate seems to occur just at the transition from high-resolution optical spectra to low-resolution spectra taken with the HST, the location of this break needs to be confirmed using high-resolution UV spectra. In a previous study, using UV observations of the extremely bright $z_{\\rm em}=1.73$ QSO HE~0515-4414 with a resolution of 30\\,000, it appeared that the slow-down in the evolution rate did not occur earlier than $z \\sim 1$ \\citep{Janknecht2002}. However, at low redshift there are few absorbers, and more LOS are necessary.  In this paper we use UV and optical high-resolution data of nine bright quasars to explore the {\\Lya} forest in the redshift range $z=0.5-1.9$. We believe that this is the first study of the lower column density absorbers in the redshift region $z<1.5$ using spectra with $R \\geq 30\\,000$. The paper is organized as follows: in Section 2, we explain our selection criteria for the quasars, describe the observations and define the redshift ranges that we investigate. Section 3 illustrates how the {\\Lya} lines in the spectra were detected and modelled. After mentioning some interesting individual features in the individual LOS in Section 4, we discuss our results in Section 5. The main results are summarized in Section 6. ",
        "conclusions": "We have analyzed the combined sample of all {\\Lya} absorption lines which were detected in the spectra of nine quasars over a total redshift path length $\\Delta\\, z=5.176$ within the range \\mbox{$z=0.5-1.9$}. We classified 1325 absorption features as {\\Lya} lines. Examining the distributions of their fit parameters {\\logn}, $b$ and $z$, we find the following:  \\begin{enumerate}  \\item{The {\\N} distribution might be complete down to log~$N_{\\rm \\ion{H}{i}}=12.90$. It can be approximated over roughly three orders of magnitude (log~$N_{\\rm \\ion{H}{i}}=12.90-15.70$) by a simple power law. The slope $\\beta=1.60\\pm0.03$ is consistent with other studies in comparable redshift intervals. We do not see any change in $\\beta$ with $z$.}  \\item{The distribution of the Doppler widths has the expected form of nearly Gaussian with an additional expanded tail to high $b$ and an average value $\\overline{b}=(34\\pm22)\\, {\\rm km\\, s}^{-1}$. We cannot detect any evolution of the Doppler parameter distribution with $z$ in the examined {\\z} phase.}  \\item{The weaker (log~$N_{\\rm \\ion{H}{i}}=12.90-14.00$) as well as the stronger {\\Lya} lines (log~$N_{\\rm \\ion{H}{i}} > 13.64$) show marginal clustering with a $2\\sigma $ significance on short velocity intervals ($\\Delta v < 200$~{\\kms} and \\mbox{$\\Delta v < 100$~{\\kms}}, respectively). However, the clustering signal is too weak to be seen in individual LOS. We do not see any change with column density or redshift.}  \\item{Using the customary power law $\\frac{{\\rm d}n}{{\\rm d}z} \\propto (1+z)^{\\gamma}$, the evolution of the number density of the weaker {\\Lya} lines declines with $\\gamma=0.74\\pm0.31$ with decreasing $z$ within the range $z=0.7-1.9$. From comparisons with investigations of other authors at higher and lower redshifts we conclude that the decrease of the weaker lines is decelerated in the phase \\mbox{$z=1-2$}, turning into a nearly flat evolution for $z \\rightarrow 0$, without showing any hint for a sharp break in the evolution.}  \\item{The number density of the stronger absorbers also decreases with decreasing $z$ ($\\gamma=1.50\\pm0.45$ for $z=0.7-1.9$), faster than for the weaker lines, though the trends are consistent with each other within $1\\sigma$.}  \\item{The decline of the line density of the stronger absorbers is also decelerated compared with higher {\\z}s \\mbox{($z=1.5-4.0$)}. The break in the evolution predicted for $z=1.5-1.7$ cannot be detected, however. The deduced $z$ dependence of the number density as well as comparisons with other studies for the local universe rather suggest a later slow-down ($z < 0.7$), followed by an approximately flat evolution. The precise $z$ value of the break in the evolution probably has to be corrected somewhat upwards due to the uncertain conversion between equivalent width (used by \\cite{Weymann} as observable) and column density.}  \\item{Though having analyzed nine high-resolution ($R \\geq 30\\,000$) quasar spectra altogether, the scatter of the data points in the diagrams of the number density evolution is large. As Tytler et al. (2004) showed, the strong variations between the LOS are consistent with the effects of normal structure formation.}  \\item{An analysis of the flux distribution of the three quasars with LOS in the upper {\\z} range ($z=1.81-1.91$) gives $\\left<F_{\\rm norm}\\right > = 0.895\\pm0.020$ for the average normalized flux and $\\tau_{\\rm eff} = 0.111\\pm0.020$ for the effective optical depth, in good agreement with the values from Kirkman et al. (2004).}  \\end{enumerate}"
    },
    "0608/math-ph0608012_arXiv.txt": {
        "abstract": "Symplectic integrators evolve dynamical systems according to modified Hamiltonians whose error terms are also well-defined Hamiltonians. The error of the algorithm is the sum of each error Hamiltonian's perturbation on the exact solution. When symplectic integrators are applied to the Kepler problem, these error terms cause the orbit to precess. In this work, by developing a general method of computing the perihelion advance via the Laplace-Runge-Lenz vector even for non-separable Hamiltonians, I show that the precession error in symplectic integrators can be computed analytically. It is found that at each order, each paired error Hamiltonians cause the orbit to precess oppositely by exactly the same amount after each period. Hence, symplectic corrector, or process integrators, which have equal coefficients for these paired error terms, will have their precession errors exactly cancel after each period. Thus the physics of symplectic integrators determines the optimal algorithm for integrating long time periodic motions. ",
        "introduction": " ",
        "conclusions": "When solving physical problems, symplectic integrators approximate the original Hamiltonian by a modified Hamiltonian with a well-defined error structure. For time-reversible integrators, the error Hamiltonians come in pairs in the form of $\\{T,Q_i\\}$ and $\\{V,Q_i\\}$. There is a clear separation between the mathematics of the algorithm, which fixes the error coefficients $e_{TQ_i}$ and $\\e_{VQ_i}$, and the physics of the problem, which determines the error Hamiltonians $\\{T,Q_i\\}$ and $\\{V,Q_i\\}$. In the past, when symplectic integrators are studied as numerical methods, only the error coefficients are analyzed so that they can be set to zero. Here, by a well-chosen example, we have shown that the physical effects of the error Hamiltonians determine how the error coefficients should be chosen. That is, the underlying physics of the problem determines the best algorithm for its own solution.  For solving celestial mechanics problems dominated by Keplerian orbits, this work shows that the optimal integrators at each order are symplectic corrector algorithms. Unfortunately, for forward algorithms without any unphysical backward intermediate time steps, this cannot be implemented without including extra error Hamiltonians. In second order, it is easy to include $H_{VTV}$, which is just a local potential. In fourth order, $H_{VTTTV}$ is a non-separable Hamiltonian too cumbersome to be solved in general. One must find ways of including $H_{VTTTV}$ without solving it directly.  The analytical results for the precession error are useful for verifying numerical calculations, however, it is a tedious way of proving the equality $\\dt_{TQ_i}=-\\dt_{VQ_i}$. It should be possible to prove this equality without explicitly evaluating individual precession angles.  We have shown in Ref.\\cite{chinsante} that the phase error in the harmonic oscillator vanishes when $e_{TQ_i}=e_{VQ_i}$. It was simply not realized in that context that $H_{TQ_i}$ and $H_{VQ_i}$ are also generating exactly opposite phase angles. From these two examples, maybe one can prove that for a general Hamiltonian with periodic orbits, only symplectic corrector algorithms can yield periodic errors for both the action and the angle variable.  Finally, this work demonstrated that one must rethink the usual practice of minimizing the sum-of-square of the error coefficients as a mean of optimizing algorithms. The error Hamiltonians are not random; they come in pairs with opposite signs. The error coefficients should therefore be chosen to be pair-wise equal, {\\it i.e.}, one should seek optimal algorithms within the class of corrector algorithms."
    },
    "0608/astro-ph0608338_arXiv.txt": {
        "abstract": "We analyse high-resolution archival UVES data of turnoff and subgiant stars in the nearby globular cluster NGC\\,6397 ([Fe/H]\\,$\\approx$\\,$-$2). Balmer-profile analyses are performed to derive reddening-free effective temperatures. Due to the limited S/N and uncertainties related to blaze removal, we find the data quality insufficient to exclude the existence of gravitational settling. If the newly derived effective temperatures are taken as a basis for an abundance analysis, the photospheric iron (Fe\\,{\\sc ii}) abundance in the turnoff stars is 0.11\\,dex lower than in the (well-mixed) subgiants. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608562_arXiv.txt": {
        "abstract": " ",
        "introduction": "The abundances of the light nuclei produced by primordial Big-Bang nucleosynthesis (BBN) provide some of the most stringent constraints on the decays of unstable massive particles during the early Universe~\\cite{holtmann, kkm, kohri, cefo, grant, kkm2, kmy, Jedamzik:2006xz,cefos2}. This is because the astrophysical determinations of the abundances of deuterium (D) and \\he4  agree well with those predicted by homogeneous BBN calculations, and also the baryon-to-photon ratio $\\eta \\equiv n_b/n_\\gamma \\propto \\Omega_b h^2$ needed for the success of these calculations~\\cite{cfo1,bbn2} agrees very well with that inferred~\\cite{cfo2} from observations of the power spectrum of fluctuations in the cosmic microwave background (CMB). The value of $\\eta = (6.11 \\pm 0.25) \\times 10^{-10}$ that they indicate~\\cite{wmap} is now quite precise, reducing one of the principal uncertainties in the previous BBN calculations.  However, it is still difficult to reconcile the BBN predictions for the lithium isotope abundances with observational indications on the primordial abundances.  The discovery of the ``Spite'' plateau~\\cite{spite}, which demonstrates a near-independence of the \\li7 abundance from the metallicity in Population-II stars, suggests a primordial abundance in the range ${\\rm \\li7/H} \\sim (1 - 2) \\times 10^{-10} $ ~\\cite{rbofn}, whereas standard BBN with the CMB value of $\\eta$ would predict ${\\rm \\li7/H} \\sim 4 \\times 10^{-10} $~\\cite{cfo1,bbn2}. In the case of $\\li6$, the data \\cite{li6obs} lie a factor $\\sim 1000$ above the BBN predictions~\\cite{bbnli6}, and fail to exhibit the dependence on metallicity expected in models based on nucleosynthesis by Galactic cosmic rays~\\cite{li6cr}. On the other hand, the $\\li6$ abundance may be explained by pre-Galactic Population-III stars, without additional over-production of $\\li7$ \\cite{rvo}.  The concordance between BBN predictions and the observed abundances of D and \\he4 is relatively fragile and could have been upset by decays of massive unstable particles.  Electromagnetic and hadronic showers produced in decays occurring during or shortly after BBN induce new reactions which may either create or destroy light nuclides. Concrete examples of unstable but long-lived particles are found in supersymmetric theories. In particular, models based on the constrained version of the minimal $R$-parity conserving supersymmetric standard model (CMSSM) with a gravitino as the lightest supersymmetric particle (LSP) have been considered in this context~\\cite{SFT,eoss5,vcmssm,EOV,Jed1,Jed2}. In these models with a gravitino LSP, the next-to-lightest supersymmetric particle (NSP) may be either the lightest neutralino $\\chi$ or the lighter stau ${\\tilde \\tau_1}$, and the decay lifetime of the NSP can range from seconds to years depending on the specific model and gravitino mass.  Our previous results concentrated on relatively long lifetimes ($\\tau > 10^4$~s), and the effects of the electromagnetic showers on the light-element abundances~\\cite{cefo,eoss5,vcmssm,EOV}.  The effects of hadronic injections due to late decays of the NSP during BBN have also been studied extensively by other authors~\\cite{kohri,kkm2,kmy,karsten,jed,stef}. In particular, it has been shown for relatively short lifetimes of order $10^3$~s that decays may simultaneously increase the \\li6 abundance and decrease the \\li7 abundance~\\cite{karsten,jed}.  On the other hand, it has been shown~\\cite{EOV} that purely electromagnetic showers cannot reduce the \\li7 abundance sufficiently without also overproducing D relative to \\he3~\\cite{sigl}. In~\\cite{SFT} the authors have calculated the principal three-body decays of a stau NSP and, using the results of the analysis~\\cite{kkm2} have explored regions where the \\li7 puzzle can be solved in an unconstrained supersymmetric model.  The authors of~\\cite{kkm2} used the {\\small PYTHIA}~\\cite{pythia} model for $e^+e^-$ annihilation to hadrons in order to simulate the hadronic decays of the unstable particle. An analogous simulation was used in~\\cite{Jed1}, both to locate regions in the parameter space of the CMSSM which are compatible with the BBN constraints, and to solve the lithium problem~\\cite{Jed2}.  It has recently been pointed out that, if it has electric charge, the NSP forms bound states with several nuclei \\cite{maxim}. Due to the large NSP mass ($m_{\\rm nsp} \\gg m_{\\rm nucleon}$), the Bohr radii of these bound states $\\sim \\alpha^{-1} m_{\\rm nucleon}^{-1} \\sim 1 \\ {\\rm fm}$ are of order the nuclear size. Consequently, nuclear reactions with nuclei in bound states are catalyzed, due to partial screening of the Coulomb barrier, and due to the opening of virtual photon channels in radiative capture reactions.  Ref.~\\cite{maxim} used analytic approximations to argue that the $d(\\alpha,\\gamma)\\li6$ reaction is enhanced by an enormous factor $\\sim 10^6$, leading to \\li6 production far beyond tolerable levels over large regions of parameter space. Other effects of bound states were considered in~\\cite{kota,kapling}.  Here, we present results from a new BBN code that includes the nuclear reactions induced by hadronic and electromagnetic showers generated by late gravitational decays of the NSP, together with the familiar network of nuclear reactions used to calculate the primordial abundances of the light elements Deuterium (D), \\he3, \\he4 and \\li7. In addition, we include the effects of the bound states when the decaying particle is charged.  As claimed in~\\cite{maxim}, these bound states lead to large enhancements in otherwise heavily suppressed rates. We calculate the abundances of the various bound states and include both the Coulomb enhancement as well as virtual photon effects in radiative capture reactions.  We use as frameworks for this study both the CMSSM and mSUGRA models~\\cite{vcmssm}, where the NSP could be either the lighter stau or the lightest neutralino.  We have calculated the dominant two- and three-body gravitational decays of these sparticles and use the {\\small PYTHIA} Monte Carlo event generator and a cascade equation to model the resulting electromagnetic (EM) and hadronic (HD)  spectra for each point of the parameter space of the supersymmetric model.  The resulting accurate determinations of the abundances of the light elements, as altered by these late injections and the bound-state effects, enable us to delineate regions of the parameter space of the supersymmetric models which are compatible with the BBN constraints. In addition, we look for regions where the \\li6 and \\li7 puzzles can be solved in the context of these supersymmetric models. We find that for lifetimes $\\tau < 10^3 - 10^4$~s, the enhanced rates of \\li6 and \\li7 production, exclude gravitino dark matter (GDM) with a stau NSP.  At smaller lifetimes, we see that it is the \\li7 destruction rates which are enhanced, facilitating a solution to the Li problems. ",
        "conclusions": "We have calculated in this paper the cosmological light-element abundances in the presence of the electromagnetic and hadronic showers due to late decays of the NSP in the context of the CMSSM and mSUGRA models, incorporating the effects of the bound states that would form between a metastable stau NSP and the light nuclei.  Late decays of the neutralino NSP constrain significantly the neutralino region, since in general they yield large light-element abundances.  The bound-state effects are significant in the stau NSP region, where excessive \\li6 and \\li7 abundances exclude regions where the stau lifetime is longer than $10^3-10^4$~s. For lifetimes shorter than $1000$~s, there is a possibility that the stau decays can reduce the \\li7 abundance from the standard BBN value, while at the same time enhancing the \\li6 abundance. A more complete account of our calculations will be given in~\\cite{cefos2}, where more examples of CMSSM and mSUGRA parameter planes will be presented, and the possibility of matching the favoured lithium abundances will be discussed in more detail."
    },
    "0608/astro-ph0608048_arXiv.txt": {
        "abstract": "Intermittency is an essential property of astrophysical fluids, which demonstrate an extended inertial range. As intermittency violates self-similarity of motions, it gets impossible to naively extrapolate the properties of fluid obtained computationally with relatively low resolution to the actual astrophysical situations. In terms of Astrophysics, intermittency affects  turbulent heating, momentum transfer, interaction with cosmic rays, radio waves and many more essential processes. Because of its significance, studies of intermittency call for coordinated efforts from both theorists and observers. In terms of theoretical understanding we are still just scratching a surface of a very rich subject. We have some theoretically well justified models that are poorly supported by experiments, we also have She-Leveque model,  which could be vulnerable on theoretical grounds, but, nevertheless, is well supported by experimental and laboratory data. I briefly discuss a rather mysterious property of turbulence called ``extended self-similarity'' and the possibilities that it opens for the intermittency research. Then I analyze simulations of MHD intermittency performed by different groups and show that their results do not contradict to each other. Finally, I discuss the intermittency of density, intermittency of turbulence in the viscosity-dominated regime as well as the intermittency of polarization of Alfvenic modes. The latter provides an attractive solution to account for a slower cascading rate that is observed in some of the numerical experiments. I conclude by claiming that a substantial progress in the field may be achieved by studies of the turbulence intermittency via observations. ",
        "introduction": "Astrophysical fluids are, as a rule, turbulent and magnetized. These two properties are closely interrelated. Turbulence itself can amplify magnetic fields stretching and winding of magnetic field lines (see Batchelor 1950) and it is an important ingredient of a more regular astrophysical dynamo (see Mofatt 1978) that is generally believed to be the origin of magnetic fields on the scale from stars and accretion disks to galaxies (see a discussion of the problems of the traditional approach in a review by Vishniac, Lazarian \\& Cho 2003 and references therein). Magnetic fields, in their turn, constrain motions of ions, which decreases the diffusivity and increases the Reynolds number $Re$ of the flow. The latter is the ratio of the eddy turnover time of a parcel of gas to the time required for viscous forces to show it appreciably. Mathematically $Re=LV_L/\\nu$, where $L$ and $V_L$ are the scale of the flow and its velocity, respectively, while $\\nu$ is the viscous diffusivity. Numerically, for most astrophysical astrophysical flows the Reynolds number is huge, e.g. $Re>10^8$. Its magnetic counterpart that characterizes to what extend the magnetic field is frozen in the fluid may be even larger, e.g. $Rm>10^{16}$. Such tremendous $Re$ and $Rm$ are not feasible to reproduce in numerical simulations neither now nor in at any foreseeable future\\footnote{Note, that the currently available 3D simulations for $1024$ cubes have $Re$ and $Rm$ limited by numerical diffusion of the order of only $10^4$ and those numbers scale linearly with the cube size.}. Is it possible to understand MHD turbulence in these circumstances?  A lot of understanding of the turbulence has been achieved by studies of turbulence self-similarity. This property, which is also called scale-invariance, implies that fluid turbulence can be reproduced by the magnification of some part of it. Take the famous model of incompressible Kolmogorov turbulence as an example. In this model the energy is injected at a large scale $L$ and forms eddies that transfer the energy to smaller and smaller scales. At the scales where the corresponding Reynolds number $Re_l\\sim lv_l/\\nu$ is much larger than unity, the dissipation over eddy turnover time $t_l\\sim l/v_l$ is negligible. As the result, the energy cascades to smaller and smaller scales without much dissipation, i.e. $v_l^2/t_l\\sim const$, which gives the well-known Kolmogorov power-law scaling for the eddies of scale $l$, namely, $v_l\\sim l^{1/3}$. The cascade terminates at the dissipation scale, which provides an {\\it inertial range} from $L$ to $L Re^{-3/4}$, where $Re$ is the Reynolds number corresponding to the flow at the injection scale $L$.  At the dissipation scales the self-similarity is known to fail with turbulence forming non-Gaussian dissipation structures as exemplified, e.g. in Biskamp (2003). Interestingly enough, present-day research shows that self-similarity is not exactly true even along the inertial range. Instead the fluctuations tend to get increasingly sparse in time and space at smaller scales. This property is called {\\it intermittency}. Note, that the power-law scaling does not guarantee the scale-invariance or absence of intermittency.  From the practical point of view, self-similarity of turbulence simplifies our description of turbulent phenomena, while intermittency spoils the nice picture. If intermittency is $Re$ or $Rm$ number dependent phenomenon one cannot believe to the results of numerical simulations because of the aforementioned disparity in the numbers. In other words, it invalidates naive extrapolations of the dynamics of the turbulent fluid obtained to with computers to real astrophysical systems. In terms of astrophysical implications, intermittency can substantially change heating of interstellar gas. According to Falgarone et al.  (2005, Falgarone 1999 and references therein) many chemical endothermal reactions (e.g. formation of CH$^+$) take place in intensive vortices within interstellar medium. Cho \\& Lazarian (2003) speculated that the mysterious structures in ISM on the AU spatial scale (see Marscher et al. 1993, Heiles 1997) can be explained by intermittency of MHD turbulence in the viscosity-damped regime. Particle diffusion and acceleration may be very different\\footnote{Recent changes in understanding of MHD turbulence have induced substantial shifts in understanding of cosmic ray diffusion and acceleration (see Yan \\& Lazarian (2004), Cho \\& Lazarian (2005) and references therein). It may happen that intermittency is another vital missing ingredient of the modern cosmic ray propagation models.} in intermittent systems (see Honda \\& Honda 2005). These and many other issues can be settled when we know more about intermittency of MHD turbulence.     One way to do such studies is to investigate the scaling powers of longitudinal velocity fluctuations, i.e. $(\\delta V)^p$, where $\\delta V\\equiv ({\\bf V}({\\bf x}+{\\bf r})-{\\bf V}({\\bf x})) {\\bf r}/r$. The infinite set of various powers of $S^p\\equiv \\langle (\\delta V)^p \\rangle$, where $\\langle ..\\rangle$ denote ensemble\\footnote{In astrophysics spatial or temporal averaging is used.} averaging, is equivalent to the p.d.f. of the velocity increments. For those powers one can write $S^p(r)= a_p r^\\xi_p$ to fully characterize the isotropic turbulent field in the inertial range. While the scaling coefficients $a_p$ are given by the values of the function $S^p$ e.g. at  the injection scale,  the scaling exponents $\\xi_p$ are very non-trivial. It is possible to show that for a self-similar flow the scaling exponents are linear function of $n$, i.e. $\\xi_p\\sim p$, which for Kolmogorov model $S^1\\sim v_l\\sim l^{1/3}$ gives $\\xi_p=p/3$. Experimental studies, however, give different results which shows that the Kolmogorov model is an oversimplified one.  MHD turbulence, unlike hydro turbulence, deals not only with velocity fluctuations, but also with the magnetic ones. The intermittencies of the two fields can be different. In addition, MHD turbulence is anisotropic as magnetic field affects motions parallel to the local direction of ${\\bf B}$ very different. This all makes it more challenging to understand the properties of MHD intermittency more interesting.   In spite of the aforementioned limitations of the computational approach to studies of astrophysical turbulence, computers provide the most cost-effective way of {\\it testing} turbulence scalings. The problem with both computational and laboratory studies of $S^n$ is a relatively small inertial range.  An interesting and yet not understood property of structure functions, however, helps to extend the range over which $S^p$ can be studied. Benzi et al. (1993) reported that for hydrodynamic turbulence the functions $S^p(S^3)$ exhibit much broader power-law range compared to $S^p(r)$. While for the inertial range a similarity in scaling of the two functions stem from the Kolmogorov scaling $S^3\\sim r$, the power-law scaling of $S^p(S^3)$ protrudes well beyond the inertial range into the dissipation range\\footnote{In practical terms this means that instead of obtaining $S^p$ as a function of $r$, one gets $S^p$ as a function of $S^3$, which is nonlinear in a way to correct for the distortions of $S^p$.} . This observation shows that the dissipation ``spoils'' different orders of $S$ in the same manner. Therefore there is no particular need to use the third moment, but one can use any other moment $S^m\\sim r^m$ and obtain a good power law of the function $S^p\\sim (S^m)^{\\xi_p/\\xi_m}$ (see Biskamp 2003).  Even not understood, the above property of $S^n$ allows to get insight into the scaling of realistic high $Re$ flows with limited $Re$ number simulations or laboratory experiments. This way the work of Benzi et al. (1993) allowed much of the progress that we describe below. ",
        "conclusions": ""
    },
    "0608/astro-ph0608095_arXiv.txt": {
        "abstract": "Satellite galaxies currently undergoing tidal disruption offer a unique opportunity to constrain an effective violation of the equivalence principle in the dark sector.  While dark matter in the standard scenario interacts solely through gravity on large scales, a new long-range force between dark-matter particles may naturally arise in theories in which the dark matter couples to a light scalar field.  An inverse-square-law force of this kind would manifest itself as a violation of the equivalence principle in the dynamics of dark matter compared to baryons in the form of gas or stars.  In a previous paper, we showed that an attractive force would displace stars outwards from the bottom of the satellite's gravitational potential well, leading to a higher fraction of stars being disrupted from the tidal bulge further from the Galactic center.  Since stars disrupted from the far (near) side of the satellite go on to form the trailing (leading) tidal stream, an attractive dark-matter force will produce a relative enhancement of the trailing stream compared to the leading stream.  This distinctive signature of a dark-matter force might be detected through detailed observations of the tidal tails of a disrupting satellite, such as those recently performed by the Two-Micron All-Sky Survey (2MASS) and Sloan Digital Sky Survey (SDSS) on the Sagittarius (Sgr) dwarf galaxy.  Here we show that this signature is robust to changes in our models for both the satellite and Milky Way, suggesting that we might hope to search for a dark-matter force in the tidal features of other recently discovered satellite galaxies in addition to the Sgr dwarf. ",
        "introduction": "\\label{S:intro} At the dawn of the era of precision cosmology, the nature of the dark sector remains the greatest mystery in cosmology and perhaps all of physics.  By definition, dark matter and dark energy cannot be observed directly; we must infer their presence and behavior through their influence on the visible components of the Universe.  Dark matter is ``detected'' through its gravitational contribution to the rotation curves of spiral galaxies and the velocity dispersion of galaxies in clusters.  Interactions purely within the dark sector that leave dark-matter density profiles largely intact are particularly difficult to constrain by observations of this kind.  Yet extensions to the standard model of particle physics might incorporate such dark-matter self-interactions, and an opportunity to constrain them in an astrophysical context could be invaluable.  In a recent {\\it Letter} \\cite{PRL}, we proposed that the tidal tails of a disrupting satellite galaxy such as the Sagittarius (Sgr) dwarf offer just such an opportunity.  Here we discuss the numerical techniques used in our simulations of tidal disruption, and show through a new series of simulations that our proposed signature of dark-matter interactions is robust to changes in our Galactic and Sgr dwarf models.  The signature we are seeking essentially tests the weak equivalence principle, the universality of free-fall between baryons and dark matter in the Galactic gravitational field.  The luminosity $L_{\\rm Sgr} = 1.4 \\times 10^7 L_{\\odot}$ of the Sgr dwarf attests to its stellar content, while its large velocity dispersion suggests the presence of considerable amounts of dark matter \\cite{2MassTT}.  If the stars and dark matter fall differently, the centers of mass of the stars and dark matter will separate and stars will no longer be disrupted at roughly equal rates from the tidal bulges on the near and far sides of the satellite.  Because the typical velocities imparted to stars during tidal disruption are much less than the satellite's orbital velocity about the Galactic center, these stars will still faithfully trace the satellite's orbit at the time of disruption \\cite{TSdyn}.  An attractive dark-matter interaction would cause the dark matter in the Sgr core to fall faster in the Galactic potential than its stars, leading to a relative enhancement of the trailing stream formed from stars disrupted on the satellite's far side.  This is a twenty-first century generalization of Galileo's possibly apocryphal test of the equivalence principle.  Galileo allegedly dropped objects of different mass off the leaning tower of Pisa at the same time and watched to see whether the heavier object left the lighter behind.  We hope to watch tidally disrupted stars and a dark-matter-dominated satellite race around the Galactic center, and see whether dark matter with a greater effective value of Newton's constant $G$ will leave the visible matter behind.  Detection of a dark-matter interaction of even a few percent the strength of gravity should be feasible with observations already in hand.  Surface brightnesses for both the leading and trailing streams of the Sgr dwarf have been observed by both the Two-Micron All-Sky Survey (2MASS) \\cite{2MassObs} and the Sloan Digital Sky Survey (SDSS) \\cite{SDSS}.  As stars disrupted from the Sgr dwarf have been identified over the full $360^{\\circ}$ of its orbit, an extremely long lever arm exists over which to make observations.  In this paper, we perform a series of N-body simulations of a satellite with similar mass and orbit to the Sgr dwarf, and show that a dark-matter interaction 1\\% the strength of gravity could reduce the ratio of leading to trailing stars by 33\\% or more, while a 4\\% dark-matter force could reduce this ratio by over 85\\%.  While these results are far from conclusive, the leading-to-trailing ratio is a crude statistic compared to what could be constructed from the distances and radial velocities of stream stars measured along the Sgr orbit.  The distances \\cite{2MassObs} and radial velocities \\cite{2MassVel} to hundreds of M-giant stars in the Sgr streams have already been determined from photometric parallaxes and spectroscopy.  Furthermore, future missions such as the Space Interferometry Mission (SIM) and Gaia will measure these quantities to unprecedented precision, and provide proper motions as well.  As this paper is only intended to draw attention to the tight limits that could be set on a dark-matter force through a more detailed comparison to observations, we restrict our attention here to the leading-to-trailing ratio.  An outline of our theory, methodology, and results is given below.  In Section~\\ref{S:DM} we consider how fundamental physics might give rise to a dark-matter self-interaction, and survey previous attempts to constrain such an interaction in astrophysical settings.  We focus on the case of tidal streams in Section~\\ref{S:tidal}, which have been used in previous research to probe the depth and shape of the host's gravitational potential.  We extend this work by incorporating a dark-matter self-interaction into pre-existing Galactic models, and propose that its effects are not degenerate with changes to other more conventional model parameters.  This hypothesis is tested using simulations whose methodology is described in Section~\\ref{S:sim}.  The results of these simulations are presented in Section~\\ref{S:res}, with some concluding remarks as to their cosmological implications and hopes for observational verification given in Section~\\ref{S:conc}. ",
        "conclusions": "\\label{S:conc}  \\begin{figure} \\scalebox{0.90}{\\includegraphics{Fig13.ps}} \\caption{The ratio of leading to trailing stars as a function of charge-to-mass ratio $\\beta$ for different models of the host-satellite system.  The black (solid) curve is our default best-fit model (runs 1a-1d). The magenta (long-short dashed) curve doubles the initial mass of the satellite (runs 2a-2c).   The two red (long-dashed) curves have satellites with spin angular momentum: the top curve has prograde spin (runs 3a-3b), while the bottom curve has retrograde spin (runs 3c-3d).  The blue (dotted) curves have satellites with different orbits: the top curve has the satellite on a more circular orbit (runs 4a-4b), while the bottom curve has the satellite on a planar orbit rather than the polar orbit of Sgr (runs 5a-5c).  The cyan (dot-short dashed) curve uses a Milky Way model with lighter halo and heavier disk (runs 6a-6c).  The green (short-dashed) curve has a satellite where 25\\% of the most bound particles represent stars (runs 7c-7d).  The yellow (dot-long dashed) curves have satellites with lower $M/L$ ratios (higher stellar mass fractions): the top curve has $M/L = 4.5$ while the bottom curve has $M/L = 10$.} \\label{F:rat} \\end{figure}  In a recent {\\it Letter} \\cite{PRL}, we introduced the idea that an asymmetry in the leading and trailing streams of a tidally disrupting satellite could be a signature of a dark-matter-only force.  Such a force cannot be detected in a system in isolation, as the stars we observe are only sensitive to the dark-matter density profile which can be maintained in the presence of an attractive (repulsive) dark-matter force by increasing (decreasing) the velocities of dark-matter particles.  However for satellites orbiting a larger host system, a dark-matter force will displace the center of mass of the satellite's dark matter with respect to that of its stars leading to several potentially observable effects.  Previous work suggested that this displacement will produce asymmetric rotation curves or warping in a satellite galactic disk \\cite{EPlong}.  Unfortunately, not all satellites have well-defined disks and we would certainly expect to observe warping in a strong tidal gravitational field.  Tidal streams provide a uniquely sensitive test of a dark-matter force because in a system {\\it already} disrupting due to gravitational effects, the dark-matter force tips the balance between the tidal bulges on the near and far side of the satellite.  Our primary concern in this paper is to make sure that some other more conventional change to our models does not produce the same effect.  To accomplish this, we explored the specifics of tidal disruption in much greater detail and examined what factors besides a dark-matter force influence the extent of tidal streams.  Features of galaxy morphology such as tidal streams have traditionally been viewed as the messy products of nonlinear dynamics, a far cry from clean probes of cosmology like the cosmic microwave background (CMB).  While it may be difficult to derive the detailed structure of an individual stream from first principles, tidal streams can be understood in general using analytic arguments supported by numerical simulations.  With the discovery of several new stellar tidal streams possibly associated with Milky Way dwarf satellites \\cite{Willman,G&O,Grill}, and the prospect of further discoveries by SDSS and future survey missions like Gaia or SIM, we believe that serious consideration of tidal streams as probes of long-range dark-matter interactions is warranted.  The wide variety of dark-matter candidates and possible dark-matter interactions makes it difficult to establish precise constraints, but any effective violation in the universality of free fall between baryons and dark-matter on scales of $\\mathcal{O}$(1 kpc) should conceivably be reflected in the morphology of tidal streams.  Many other factors besides a dark-matter force affect the amount of tidal disruption, and in Section~\\ref{S:res} we sought to survey some the most important of these factors.  The results of this survey are summarized in Fig.~\\ref{F:rat}.  We have counted the number of leading stellar particles $N_{\\rm lead}$ and trailing stellar particles $N_{\\rm trail}$ and plotted their ratio as a function of $\\beta$ for the simulations listed in Table~\\ref{T:runs}.  In most of the simulations, we have counted the stellar particles in $100^{\\circ}$ slices of angular separation $\\theta$ from the satellite core as shown on the $x$-axis of Fig.~\\ref{F:SD1}.  The position of these slices was chosen to include the stars piled up at the apocenters furthest along the leading and trailing streams.  There were no prominent accumulations near apocenter in the simulations with a nearly circular orbit (runs 4a-4b) and those in which the stellar particles were 25\\% of the the most bound particles (runs 7c-7d), so we counted all stellar particles further than $10^{\\circ}$ from the satellite core.  Runs 7a and 7b were not presented on this plot as there was virtually no tidal disruption in these simulations.  At first glance, Fig.~\\ref{F:rat} confirms our worst suspicions of tidal streams as a probe of cosmology.  In the absence of a dark-matter force $(\\beta = 0.0)$, the leading-to-trailing ratio still varies by almost a factor of two between simulations.  Changes like doubling the satellite mass or adding prograde spin dramatically increase the total amount of tidal disruption and increase the leading-to-trailing ratio to some extent as well.  In contrast, retrograde spin and a planar orbit reduce the leading-to-trailing ratio compared to the default simulations shown by the black solid curve. However, in spite of all these changes, the leading-to-trailing ratio exceeds 0.5 for {\\it all} simulations without a dark-matter force and {\\it never} exceed 0.2 for $\\beta = 0.2$, a dark-matter force only 4\\% the strength of gravity.  While it would be premature to conclude that a dark-matter force exists based on a single number, measurements of the satellite core and stellar distances and radial velocities along the streams should tell us about what leading-to-trailing ratio to expect.  An observed leading-to-trailing ratio significantly below (above) this value would alert us to the possible presence of an attractive (repulsive) dark-matter force.  Even if observations of a single tidal stream did not convince us of the existence of a dark-matter force, several tidal streams have already been discovered about the Milky Way.  A consistent asymmetry observed in different satellites on different orbits might be quite persuasive.  Before concluding, we wish to make a brief comment in the implications of modified Newtonian dynamics (MOND) for tidal streams.  The model explored in this paper is not an example of MOND; it includes copious amounts of dark matter and explains all conventional dark-matter signatures (e.g. galactic rotation curves, cluster velocities dispersions) through the Newtonian gravitational influence of this dark matter.  Nonetheless, one might imagine that the Sgr tidal streams could place interesting constraints on MOND, and indeed this was the subject of a recent paper \\cite{ReadMOND}.  That work focused on the precession of the Sgr orbital plane; naively one would expect MOND to overpredict the amount of precession as the potential would be less spherically symmetric in the absence of a dark-matter halo.  However, the potential sourced by an infinitely thin disk is much more spherical in MOND than Newtonian gravity, implying an amount of precession similar to that of a mildly oblate dark-matter halo $(q = 0.9)$. MOND could not explain the discrepancy between Sgr leading-stream velocity data, which favors a prolate halo, and the orbital precession which prefers an oblate halo \\cite{2MassTT}.  We considered MOND using the analytical argument presented in Section~\\ref{S:tidal}, with an acceleration proportional to $1/r$ below the MOND scale $a_0 = 1.2 \\times 10^{-10}$ m/s$^2$.  This flatter potential led to a slightly larger natural asymmetry $\\Delta r_{\\rm nat}$ despite the smaller ratio $m_{\\rm sat}/M_R$ in the absence of dark matter.  MOND does not provide a mechanism for the strong suppression of the leading (trailing) tidal stream we expect for an attractive (repulsive) dark-matter force for $\\beta \\gtrsim 0.2$.  Our model also does not seem to offer an explanation for the halo-oblateness discrepancy for acceptable values of $\\beta$, though we have not explored what a careful conspiracy of parameters might allow.  Although we believe using tidal streams to constrain a dark-matter force is a promising approach, significant observational and theoretical challenges remain before it can be practically implemented.  Accurate measurements of the satellite core are essential to determining its mass, mass-to-light ratio, and possible spin.  We must also find a reliable method of identifying stream stars at large separations $\\theta$ from the satellite core, and must take care that this method does not introduce a bias between trailing and leading stars.  On the theoretical front, more work needs to be done to establish a correspondence between data and simulations.  It may not be feasible to evolve streams backwards into a bound satellite, so we need to develop techniques to guess the appropriate initial conditions more systematic than trial and error.  We hope to confront these issues in the course of searching for a dark-matter force in actual observations of the Sgr tidal stream.  While neither theory nor observation offer as yet a compelling reason to believe in the existence of a dark-matter force, the tremendous payoff the discovery of such a force could provide in terms of learning physics beyond the Standard Model makes the search well worth the effort."
    },
    "0608/astro-ph0608576_arXiv.txt": {
        "abstract": "The Pierre Auger Observatory aims to determine the nature and origin of the ultra-high energy cosmic rays (UHECR). The Auger hybrid detector combines fluorescence observations of extended air showers, initiated in the atmosphere by these most energetic particles, with measures of the shower front at the ground level by its large array of Cherenkov water tanks. This allows to improve considerably the precision on reconstructed primary cosmic ray parameters, and to make important cross-calibrations between two techniques at these energies, unreachable with accelerator experiments. The Southern Auger site in Argentina is close to completion. The first measure of the primary cosmic rays energy spectrum, the anisotropy search results, and the limit on the photon fraction in the UHECR are discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608606_arXiv.txt": {
        "abstract": "The Laser Interferometer Gravitational-wave Observatory (LIGO) has performed the fourth science run, S4, with significantly improved interferometer sensitivities with respect to previous runs. Using data acquired during this science run, we place a limit on the amplitude of a stochastic background of gravitational waves. For a frequency independent spectrum, the new limit is $\\Omega_{\\rm GW} < 6.5 \\times 10^{-5}$. This is currently the most sensitive result in the frequency range 51-150 Hz, with a factor of 13 improvement over the previous LIGO result. We discuss complementarity of the new result with other constraints on a stochastic background of gravitational waves, and we investigate implications of the new result for different models of this background. ",
        "introduction": "A stochastic background of gravitational waves (GWs) is expected to arise as a superposition of a large number of unresolved sources, from different directions in the sky, and with different polarizations. It is usually described in terms of the GW spectrum: \\begin{equation} \\Omega_{\\rm GW}(f) = \\frac{f}{\\rho_c} \\; \\frac{d \\rho_{\\rm GW}}{d f}\\,, \\end{equation} where $d\\rho_{\\rm GW}$ is the energy density of gravitational radiation contained in the frequency range $f$ to $f+df$ \\citep{allenromano}, $\\rho_c$ is the critical energy density of the Universe, and $f$ is frequency (for an alternative and equivalent definition of $\\Omega_{\\rm GW}(f)$ see, for example, \\citep{baskaran}). In this paper, we will focus on power-law GW spectra: \\begin{equation} \\Omega_{\\rm GW}(f) = \\Omega_{\\alpha} \\Bigg( \\frac{f}{100 {\\rm \\; Hz}} \\Bigg)^{\\alpha}. \\label{template} \\end{equation} Here, $\\Omega_{\\alpha}$ is the amplitude corresponding to the spectral index $\\alpha$. In particular, $\\Omega_0$ denotes the amplitude of the frequency-independent GW spectrum.  Many possible sources of stochastic GW background have been proposed and several experiments have searched for it (see \\citep{maggiore,allen} for reviews). Some of the proposed theoretical models are cosmological in nature, such as the amplification of quantum vacuum fluctuations during inflation \\citep{PA1}, \\citep{PA2}, \\citep{star}, pre-big-bang models \\citep{pbb}, \\citep{pbbrep}, \\citep{BMU}, phase transitions \\citep{PT1}, \\citep{PT2}, and cosmic strings \\citep{CS1}, \\citep{CS2}, \\citep{CS3}. Others are astrophysical in nature, such as rotating neutron stars \\citep{RFP}, supernovae \\citep{coward} or low-mass X-ray binaries \\citep{cooray}.  A number of experiments have been used to constrain the spectrum of GW background at different frequencies. Currently, the most stringent constraints arise from large-angle correlations in the cosmic microwave background (CMB) \\citep{cobe1,cobe2}, from the arrival times of millisecond pulsar signals \\citep{pulsar}, from Doppler tracking of the Cassini spacecraft \\citep{doppler}, and from resonant bar GW detectors, such as Explorer and Nautilus \\citep{bars}. An indirect bound can be placed on the total energy carried by gravitational waves at the time of the Big-Bang Nucleosynthesis (BBN) using the BBN model and observations \\citep{BBN,maggiore,allen}. Similarly, \\citep{smith} used the CMB and matter spectra to constrain the total energy density of gravitational waves at the time of photon decoupling.  Ground-based interferometer networks can directly measure the GW strain spectrum in the frequency band 10 Hz - few kHz, by searching for correlated signal beneath uncorrelated detector noise. LIGO has built three power-recycled Michelson interferometers, with a Fabry-Perot cavity in each orthogonal arm. They are located at two sites, Hanford, WA, and Livingston Parish, LA. There are two collocated interferometers at the WA site: H1, with 4km long arms, and H2, with 2km arms. The LA site contains L1, a 4km interferometer, similar in design to H1. The detector configuration and performance during LIGO's first science run (S1) was described in \\citep{S1}. The data acquired during that run was used to place an upper limit of $\\Omega_0 < 44.4$ on the amplitude of a frequency independent GW spectrum, in the frequency band 40-314 Hz \\cite{S1stoch}. This limit, as well as the rest of this paper, assumes the present value of the Hubble parameter $H_0 = 72$ km/s/Mpc or, equivalently, $h_{100} = H_0 / (100 {\\rm \\; km/s/Mpc}) = 0.72$ \\citep{hubble}. The most recent bound on the amplitude of the frequency independent GW spectrum from LIGO is based on the science run S3: $\\Omega_0 < 8.4 \\times 10^{-4}$ for a frequency-independent spectrum in the 69-156 Hz band~\\citep{S3paper}. \\begin{figure} \\epsscale{.80} \\plotone{f1.eps} \\caption{Typical strain amplitude spectra of LIGO interferometers during the science run S4 (solid curves top-to-bottom at 70 Hz: H2, L1, H1). The black dashed curve is the LIGO sensitivity goal. The gray dashed curve is the strain amplitude spectrum corresponding to the limit presented in this paper for the frequency-independent GW spectrum $\\Omega_0 < 6.5 \\times 10^{-5}$.} \\label{strain} \\end{figure}  In this paper, we report much improved limits on the stochastic GW background around 100 Hz, using the data acquired during the LIGO science run S4, which took place between February 22, 2005 and March 23, 2005. The sensitivity of the interferometers during S4, shown in Figure \\ref{strain}, was significantly better as compared to S3 (by a factor $10\\times$ at certain frequencies), which leads to an order of magnitude improvement in the upper limit on the amplitude of the stochastic GW background: $\\Omega_0 < 6.5 \\times 10^{-5}$ for a frequency-independent spectrum over the 51-150 Hz band.  This limit is beginning to probe some models of the stochastic GW background. As examples, we investigate the implications of this limit for cosmic strings models and for pre-big-bang models of the stochastic gravitational radiation. In both cases, the new LIGO result excludes parts of the parameter space of these models.  The organization of this paper is as follows. In Section 2 we review the analysis procedure and present the results in Section 3. In Section 4, we discuss some of the implications of our results for models of a stochastic GW background, as well as the complementarity between LIGO and other experimental constraints on a stochastic GW background. We conclude with future prospects in Section 5. ",
        "conclusions": ""
    },
    "0608/astro-ph0608356_arXiv.txt": {
        "abstract": "New NTT/SOFI imaging and spectroscopy of the Wolf-Rayet population in Westerlund~1 are presented. Narrow-band near-IR imaging together with follow up spectroscopy reveals four new Wolf-Rayet stars, of which three were independently identified recently by Groh et al., bringing the confirmed  Wolf-Rayet content to 24 (23 excluding source S) -- representing 8\\% of the known Galactic Wolf-Rayet population -- comprising 8 WC stars and 16 (15) WN stars. Revised coordinates and near-IR photometry are presented, whilst a quantitative near-IR spectral classification scheme for Wolf-Rayet stars is presented and applied to members of Westerlund~1. Late subtypes are dominant, with no subtypes earlier than WN5 or WC8 for the nitrogen and carbon sequences, respectively. A qualitative inspection of the WN stars suggests that most ($\\sim$75\\%) are highly H-deficient. The Wolf-Rayet binary fraction is high ($\\geq$62\\%), on the basis of dust emission from WC stars, in addition to a significant WN binary fraction from hard X-ray detections according to Clark et al. We exploit the large WN population of Westerlund~1 to reassess its distance ($\\sim$5.0~kpc) and extinction ($A_{K_{S}} \\sim$ 0.96 mag), such that it is located at the edge of the Galactic bar, with an oxygen metallicity $\\sim$60\\% higher than Orion. The observed ratio of WR stars to red and yellow hypergiants, N(WR)/N(RSG+YHG) $\\sim$ 3, favours an age of $\\sim$4.5--5.0 Myr, with individual Wolf-Rayet stars descended from progenitors of initial mass $\\sim 40-55 M_{\\odot}$. Qualitative estimates of current masses for non-dusty, H-free WR stars are presented, revealing $10-18 M_{\\odot}$, such that $\\sim$75\\% of the initial stellar mass has been removed via stellar winds or close binary evolution. We present a revision to the cluster turn-off mass for other Milky Way  clusters in which Wolf-Rayet stars are known, based upon the latest temperature calibration for OB stars. Finally, comparisons between the observed WR population and subtype distribution in Westerlund~1 and instantaneous burst evolutionary synthesis models are presented. ",
        "introduction": "Several hundred Wolf-Rayet (WR) stars -- the evolved descendants of the most massive O stars -- have been identified within the Milky Way (van der Hucht 2006). In principle, studies of WR stars in open clusters provides excellent observational constraints upon their ages and initial masses (e.g. Schild \\& Maeder 1984). In practice, this has proved challenging, due to the small number of WR stars observed within individual Milky Way or Magellanic Cloud clusters. This is unsurprising, given their short lifetimes and small number of suitably massive stars in typical open clusters ($\\approx 10^{3} M_{\\odot}$), for which empirical studies indicate a maximum stellar mass of $\\sim 30-35 M_{\\odot}$ (e.g. Weidner \\&  Kroupa 2006). Hitherto, solely the Arches, Quintuplet and Galactic Centre clusters have offered the potential for studying large numbers of WR stars in suitably massive ($10^{4} M_{\\odot}$) clusters, for which significantly higher mass stars are observed, albeit hindered by exceptionally high interstellar extinction.  Fortunately, Westerlund~1 (Westerlund 1961) offers the possibility of undertaking a comprehensive study of WR stars at Solar or moderately super-Solar metallicity, since its mass has been estimated at 10$^{5} M_{\\odot}$ (Clark et al. 2005). Instantaneous burst evolutionary synthesis models at Solar metallicity (e.g. Starburst99, Leitherer et al. 1999) predict 20--30 WR stars at an age of 4.5~Myr for such a high cluster mass. Observationally,  Clark \\&  Negueruela (2002), Negueruela  \\& Clark (2005) and Hopewell et al. (2005)  have identified 20   WR stars within 4 arcmin ($\\sim$6 pc at a distance of 5~kpc) of the central cluster, most of which were discovered  serendipitously.  \\begin{table*} \\caption{Catalogue of Wolf-Rayet stars in Westerlund~1, including WR nomenclature following van der Hucht (2006), following previous studies by Clark \\& Negueruela (2002, CN02), Negueruela \\& Clark (2005, NC05), Hopewell et al. (2005, H05), Negueruela (2005, priv comm, N05) and Groh et al. (2006, G06). Photometry is obtained primarily from NTT/SOFI or 2MASS (Skrutskie et al. 2006). Coordinates are obtained from our NTT/SOFI images, except where noted, based upon astrometry of A--G from 3.6cm radio images (see CN02). Previously published coordinates for I and J were in error due to spatial crowding.} \\label{WRcat} \\begin{tabular}{ c@{\\hspace{2mm}} l@{\\hspace{3mm}} l@{\\hspace{2mm}} l@{\\hspace{3mm}} l@{\\hspace{2mm}} r@{\\hspace{3mm}} r@{\\hspace{3mm}} r@{\\hspace{3mm}} c@{\\hspace{3mm}} l@{\\hspace{3mm}} l@{\\hspace{3mm}} l} \\hline Source & Alias & WR & $\\alpha$ & $\\delta$  & J & H & K$_{S}$ & Ref & Previous & Ref & This\\\\ &       &    & \\multicolumn{2}{c}{J2000} & mag & mag & mag & &  Sp Type  & & Study\\\\ \\hline A & Wd1-72 & WR77sc & 16 47 08.32 & --45 50 45.5 & 10.34 & 9.11: & 8.37: & SOFI & WN4--5, $<$WN7 & CN02, NC05 & WN7b \\\\ B  &  & WR77o & 16  47 05.36 & --45 51 05.0 & 10.91 & 9.79 & 9.18 & SOFI & WNL, WN8? & CN02, NC05 & WN7o \\\\ C &  & WR77m & 16 47 04.40 & --45 51 03.8 & 11.26 & 9.51 & 8.23 & SOFI & WC8, WC8.5 & CN02, NC05 & WC9d \\\\ D &  & WR77r & 16 47 06.24 & --45 51 26.5 & 11.63 & 10.31 & 9.61 & SOFI & WN6--8  & CN02 & WN7o \\\\ E & Wd1-241 & WR77p & 16 47 06.05 & --45 52 08.2 & 10.12 & 9.09 & 8.29: & SOFI & WC9  & CN02 & WC9 \\\\ % F & Wd1-239 & WR77n & 16 47 05.22 & --45 52 25.0 & 9.85 &  7.97 & 7.28 & SOFI & WC9  & CN02 & WC9d \\\\ G &   & WR77j & 16 47 04.01 & --45 51 25.2 & 11.36 & 9.97 & 9.28 & SOFI &  WN6--8  & CN02 & WN7o \\\\ H &  & WR77l & 16 47 04.22 & --45 51 20.2 & 10.31 & 8.56 & 7.38 & SOFI & WC9  & CN02 & WC9d \\\\ I &     & WR77c & 16 47 00.88 & --45 51 20.8 & 10.89 & 9.57 & 8.86 & SOFI& WN6--8   & CN02 & WN8o \\\\ J &      & WR77e & 16 47 02.47 & --45 51 00.1 &11.7: & 10.3: & 9.7: & SOFI & WNL   & CN02 & WN5h \\\\ K &       & WR77g & 16 47 03.25 & --45 50 43.8 & 11.81 &  10.40 & 9.53 & SOFI & WC, WC7 & CN02, NC05 & WC8 \\\\ L & Wd1-44 & WR77k & 16 47 04.19 & --45 51 07.4 & 9.08 & 7.72 & 7.19 & SOFI & WN9   & NC05 & WN9h:\\\\ M & Wd1-66 & WR77i & 16 47 03.96 & --45 51 37.8 & 10.13 & 7.64 & 6.9: & SOFI& WC9  &  NC05 & WC9d \\\\ % N &     & WR77b   & 16 46 59.9 & --45 55 26 &9.69 & 7.84 & 6.41 &2MASS & WC8  & NC05 & WC9d\\\\ O &    & WR77sb & 16 47 07.66 & --45 52 35.9 & 11.00 & 9.98 & 9.45 & SOFI & WN6   & NC05 & WN6o \\\\ P & Wd1-57c & WR77d & 16 47 01.59 & --45 51 45.5 & 11.06 & 9.83 & 9.26 & SOFI & WN8  & NC03 & WN7o \\\\ Q &      & WR77a  & 16 46 55.55 &--45 51 35.0&11.72& 10.67&10.00& SOFI & WN6--7  & NC05 & WN6o\\\\ R & Wd1-14c & WR77q & 16 47  06.07 & --45 50 22.6 & 11.92 & 10.84 & 10.26 & SOFI & WN6--7   & NC05 & WN5o \\\\ % S & Wd1-5 & WR77f & 16 47 02.98 & --45 50 20.0 & 9.81 & 8.80: & 8.29 & SOFI & WNVL  & NC05 & WN10-11h or B0--1Ia$^{+}$\\\\ % T & HBD4 & WR77aa & 16 46 46.3 &--45 47 58 &10.04&8.21&6.72 & 2MASS & WC9d   & H05 & WC9d\\\\ U & \\#1  & WR77s & 16 47 06.55 & --45 50 39.0 & 10.77 & 9.72 & 9.20 & SOFI & WN4, WN5--7  & N05, G06 & WN6o \\\\ V &  & WR77h & 16 47  03.81 & --45 50 38.8 &  10.75 & 9.42 & 8.76 & SOFI & WN8  & N05 & WN8o \\\\ W & \\#3 & WR77sa & 16 47 07.58 & --45 49 22.2 & 12.11 & 10.75 & 10.04 & SOFI & WN5--6   & G06 & WN6h \\\\ X & \\#2 & WR77sd & 16 47 14.1 & --45 48 32 & 12.36 & 11.08 & 10.25 & 2MASS & WN4--5  & G06 & WN5o \\\\ \\hline \\end{tabular} \\end{table*}  \\begin{figure} \\centerline{\\psfig{figure=crowther-f1a.eps,width=8.8cm,angle=0}} \\centerline{\\psfig{figure=crowther-f1b.eps,width=8.8cm,angle=0}} \\caption{(upper panel) {\\it Spitzer} IRAC (3.6$\\mu$m) image of the 10$\\times$10 arcmin (15 $\\times 15$ pc at a distance of 5~kpc) region surrounding Westerlund~1 surveyed with NTT/SOFI, taken from the GLIMPSE survey (Benjamin et al. 2003). Identifications of the four WR stars located beyond the cluster centre (shown as a box); (lower panel) ESO/NTT SOFI Br$\\gamma$ image of the central  2.5$\\times$2.5 arcmin (3.75$\\times$3.75 pc)  of  Westerlund~1, for which the  Wolf-Rayet stars have been indicated, plus W9 (sgB[e]) and W243 (LBV, Clark et al. 2005). North  is up, with east to the left in both images.} \\label{Brg} \\end{figure}  The present study examines the Wolf-Rayet content of Westerlund~1 based upon near-infrared narrow-band imaging and follow up spectroscopy. High interstellar extinction ($A_{\\rm  V}$= 11.6 mag, Clark et al. 2005) prohibits the standard optical approach of using a narrow-band He\\,{\\sc ii} $\\lambda$4686 filter and adjacent continuum filter to identify WR candidates (e.g. Hadfield et al. 2005). Homeier et al. (2003ab) have previously used interference filters within the K-band to identify WR candidates within the inner Milky Way. However, in contrast to the optical technique, the K-band is observationally more challenging since (i) WR line fluxes are much weaker; (ii) there is no single pair of interference filters that can applied to identify all WR subtypes; (iii) dust emission in WC stars may heavily dilute the WR stellar emission line fluxes.   The present paper is structured as follows. New K-band and Y-band imaging of Westerlund~1 is presented in Sect.~\\ref{obs} together with near-IR spectroscopy. Near-IR spectral classification of WR stars in Westerlund~1 is presented in Sect.~\\ref{class}. The reddening and distance to  Westerlund~1 are reassessed from its WR population in Sect.~\\ref{Wd1}, and comparisons made with other Milky Way clusters hosting WR stars. The observed WR content of Westerlund~1 is compared to predictions from single and binary models in Sect~\\ref{evol}. Finally, brief conclusions are drawn in Sect.~\\ref{conc} ",
        "conclusions": "\\label{conc}  We present interference filter  Y- and K-band imaging and follow up near-IR spectroscopy  in order to identify the complete WR population within Westerlund~1. Our results are as follows:  \\begin{enumerate} \\item Four WR stars in Westerlund~1 are newly confirmed here, with nomenclatures of U, V, W, X following Clark \\& Negueruela (2002), Negueruela \\& Clark (2005) and Hopewell et al. (2005), bringing the total number of WR stars to 24. Of these, U and V were previously discovered from optical spectroscopy (Negueruela, 2005, priv. comm.),  whilst U, W and X have been independently reported by Groh et al. (2006). Indeed, our photometric  approach  is similar to Groh et al. (2006) except that  we probe  deeper in the Y-band and cover a larger field, but neglect the He\\,{\\sc ii} 1.0124$\\mu$m filter.  \\item Intermediate resolution JHK spectroscopy  for all 24 WR stars in Westerlund~1 are presented. We present a quantitative near-IR classification scheme for WN and WC stars, calibrated using line strengths measured in optically visible field WR stars. From our sample, 16 WN stars are fairly evenly divided between mid (WN5--6)  and late ($\\geq$WN7) subtypes, whilst all 8 WC stars are late subtypes (7/8 are WC9 stars). The majority ($\\sim$75\\%) of early and late WN stars are H-poor, as estimated from the strength of Paschen/Brackett lines. Amongst the WN stars, solely star A shows a strong, broad line spectrum (WN7b), which is indicative of a dense, fast stellar wind.  \\item Dominant late WC populations, such as those in Westerlund~1 are well known in metal-rich regions such as the  inner Milky Way (Hopewell et al. 2005) and M83 (Hadfield et al. 2005). Early subtypes dominate in metal-poor regions. The reason for this dichotomy appears to be, at least in part, due to  the sensitivity of WC spectral classification to (metallicity-dependent) wind strength (Crowther et al. 2002), rather than elemental abundance differences.  \\item From a total of eight WC stars within Westerlund 1, six contain hot dust, which has the effect of strongly diluting the K-band spectral features with respect to shorter wavelengths - e.g. C\\,{\\sc iii} 2.11$\\mu$m/0.971$\\mu$m = 0.42 for the only apparent non-dusty WC9 star (E) versus 0.02 for star F. 2$\\mu$m  narrow-band surveys  presently underway to identify visibly obscured WR stars in the Milky Way (e.g. Homeier et al. 2003ab) will not be so sensitive to such dusty WC stars due to the dilution of their emission lines by dust continuum. Instead, dust forming WC stars will most easily be detected at shorter wavelengths where their stellar signature is much more prominent, or via their hot dust signatures from near and mid-IR photometry.  \\item It is widely believed that dust formation in WR stars requires the presence of an OB companion, such that most WC stars within Westerlund~1 are binaries. The presence of hard X-ray emission from a subset of WN stars (Skinner et al. 2006; Clark et al. 2006) indicates a minimum WR binary fraction of $\\sim$38\\% in Westerlund~1, or more likely 62\\%. Indeed, our present observations do not exclude a WR binary fraction of 100\\%.  \\item We apply an absolute K-band calibration for our (non-dusty) WR stars to estimate the distance and extinction, revealing $\\sim$5.0 kpc and $A_{K_{S}}$=0.96 $\\pm$0.14 mag, in reasonable agreement with previous studies based upon yellow hypergiants and pre-main sequence stars (Clark et al. 2005; Brandner et al. 2006).  WC9 stars were excluded from our calibration since non-dusty WC9 stars were not previously known within optically visible clusters (van der Hucht \\& Williams 2006). Applying our Westerlund~1 distance to the only non-dusty WC9 star (star E), suggests $M_{K_{S}}$ = --6.2 mag (the remainder span --7.0 to --9.1 mag).  \\item For an assumed Solar galactocentric distance of 8.5 kpc, Westerlund~1 lies at a galactocentric distance of 4.2$^{+0.75}_{-0.35}$ kpc, i.e.  similar to the outer edge of the Galactic bar (Benjamin et al. 2005). This location may prove significant for the formation of a 10$^{5} M_{\\odot}$ cluster within the Milky Way. Based on the Esteban et al. (2005) Milky Way oxygen metallicity gradient, Westerlund~1 is anticipated to be moderately oxygen-rich, exceeding that of the Orion Nebula by 60\\%, i.e.  $\\log$ (O/H) + 12 $\\sim$8.87.  \\item For the hydrogen-deficient, non-dusty WR stars we have applied model atmosphere bolometric corrections and intrinsic visual to JHK$_{S}$ colours to provide estimates of stellar luminosities, from which masses follow using the mass-luminosity relation of Schaerer \\& Maeder (1992). We estimate current masses of 10--18$M_{\\odot}$, typical of binary orbit deduced masses for Milky Way Wolf-Rayet stars (e.g. Crowther 2007).  \\item  We provide revised mass limits for WR progenitors in other optically visible Milky Way clusters, updated from Massey et al. (2001). We adopt revised OB spectral type--temperature calibrations and apply Solar metallicity isochrones from Meynet et al. (2004). Westerlund~1 more than doubles the known statistics, and includes subtypes not previously covered. The observed ratio N(WR)/N(RSG + YHG) $\\sim$3 for Westerlund~1 suggests an age of $\\sim$4.5--5.0~Myr, from comparison with evolutionary models, such that the WR stars are descended from stars of initial mass $40-55 M_{\\odot}$ with potential implication for a higher progenitor mass for the pulsar recently observed in Westerlund~1 (Muno et al. 2006). Consequently, approximately 75\\% of the WR progenitor mass has been removed due to stellar winds or close binary mass-transfer.  \\item Comparisons between the observed WR population in Westerlund~1 and evolutionary models for rotating/non-rotating Solar metallicity stars is rather poor, in the sense that the observed ratios of N(WC)/N(WN) and N(WN H-rich)/N(WN H-poor) are rather lower than predicted, for both single and binary models. \\end{enumerate}  We intend to follow up the present observational study with quantitative analysis of the WR population in Westerlund~1 from which we will be able to extract more reliable abundances, stellar luminosities and current masses, plus wind properties, with which to better test current evolutionary predictions for single and binary massive stars."
    },
    "0608/astro-ph0608483_arXiv.txt": {
        "abstract": "Unique among neutron stars, \\ns\\ is an X-ray pulsar with a spin period of 424 ms that contains at least two strong absorption features in its energy spectrum.  This neutron star is positionally coincident with the supernova remnant \\snr\\ and has been identified as a member of the growing class of radio-quiet  compact central objects in supernova remnants. From previous observations with \\cxo\\ and \\xmm, it has been found that the \\ns\\ is not spinning down monotonically as  is common for young, isolated pulsars.  The spin frequency history requires either  strong, frequent glitches, the presence of a fall-back disk,  or a binary companion.  Here, we report on a sequence of seven \\xmm\\ observations of \\ns\\ performed during a 40 day window between  2005 June 22 and July 31.  Due to unanticipated variance in the phase measurements during the observation period that  was beyond the statistical uncertainties, we could not identify a unique phase-coherent timing solution.  The three most probable timing solutions give frequency time derivatives of $+$0.9, $-$2.6, and $+$1.6 $\\times$ 10$^{-12}$ Hz s$^{-1}$ (listed in descending order of significance).  We conclude that the local frequency derivative during our \\xmm\\ observing campaign differs from the long-term spin-down rate by more than an order of magnitude.  This measurement effectively rules out glitch models for \\ns.  If the long-term spin frequency variations are caused by timing noise, the strength of the timing noise in \\ns\\ is much stronger than in other pulsars  with similar period derivatives.   Therefore, it is highly unlikely that the spin variations are caused by the same physical process that causes timing noise in other isolated pulsars.  The most plausible scenario for the observed spin irregularities is the presence of a binary companion  to \\ns. We identified a family of orbital solutions that are consistent with our phase-connected timing solution, archival frequency measurements, and constraints on the companions mass imposed by deep IR and optical observations. ",
        "introduction": "There exist a handful of enigmatic X-ray point sources, very likely young neutron stars, positionally coincident with supernova remnants (SNRs) whose nature remains uncertain.  These objects are commonly referred to as central compact objects (CCOs) that are characterized by soft, thermal X-ray  spectra and an absence of ordinary pulsar activity such as radio pulsations, $\\gamma$-ray emission and pulsar wind nebulae (Pavlov et al.\\ 2002a, 2004).  The CCO \\ns\\ (\\nss\\ hereafter) in the PKS 1209$-$51/52 SNR is a particularly interesting member of this class in that it is both an X-ray pulsar (with a 0.424 s period; Zavlin et al.\\ 2000)  and the only CCO found to possess prominent  absorption lines in its  spectrum (Sanwal et al.\\ 2002). The X-ray energy spectrum is best modeled with a continuum blackbody component of temperature $kT \\approx$ 0.14 keV and at least two  broad absorption lines centered at 0.7 and 1.4 keV. The strength of these lines depend upon the rotational phase of the pulsar (Mereghetti et al.\\ 2002).  Two additional features, at 2.1 and 2.8 keV, have been reported (Bignami et al.\\ 2003); however,  their validity has been questioned (Mori, Chonco \\& Hailey 2005).  The physical origin for the spectral lines remains unknown.   Sanwal et al.\\ (2002) have concluded that these lines cannot be associated with transitions in Hydrogen atoms and argued that neither electron nor proton cyclotron resonance could cause these features. These authors suggest that the lines could be due to absorption by once-ionized Helium in a magnetic field $B\\sim 2\\times 10^{14}$ G (see also Pavlov \\& Bezchastnov 2005), while Hailey \\& Mori (2002) and Mori \\& Hailey (2006) argue that the lines could be formed in an Oxygen atmosphere with $B\\sim 10^{11}$--$10^{12}$ G. As different interpretation imply very different magnetic field strengths, measuring the field strength of \\nss\\ would be most important for understanding the nature of the spectral lines. For isolated pulsars, the most straightforward method for estimating  dipole magnetic field strengths is to measure the spin frequency $\\nu$ and its time derivative (spin-down rate) $\\dot{\\nu}$.  Zavlin, Pavlov \\& Sanwal  (2004) studied the spin evolution of \\nss\\ using a compilation of \\cxo\\ and \\xmm\\ observations covering 3.5 years.  They found that the spin frequency of this pulsar is not steadily decreasing as one would expect for a magnetically-braking dipole.   Instead, the frequency evolution was quite erratic, leading Zavlin et al.\\ to consider three possible explanations: ($i$) the star is undergoing frequent  glitches, ($ii$) the star is surrounded by a debris disk that influences its spin evolution through accretion and propeller torques, or ($iii$) the star is a member of a (non-accreting) binary system.  Another possibility is that the spin-down of the star is influenced by timing noise, a ubiquitous property of isolated neutron stars  whose physical origin is unclear.  The sparse frequency history of \\nss\\ could not distinguish between these models.  Here, we report on a sequence of seven \\xmm\\ observations of \\nss\\ that were designed to perform phase-coherent timing in order to precisely measure the pulse frequency and frequency derivative of the source.  A measurement of the local frequency derivative  would  help us to distinguish between the possible scenarios.  Below, we describe  the observation (\\S2) and our timing analysis of the \\xmm\\ data set (\\S3),  and discuss the resulting constraints on the physical mechanisms for the spin-frequency evolution in \\nss\\ (\\S4). ",
        "conclusions": "We observed \\nss\\ with \\xmm\\ seven times during the course of a 40 day interval in an effort to measure the local pulse frequency and frequency derivative with high precision.  Due to unanticipated phase noise, we were unable to phase-connect the full data set.  From systematic pulse ephemeris searches, we identified a number of possible timing solutions, none of which had spin-down rates close to the average spin-down rate of \\nss\\ over the last 5 years.   The interpretation of the erratic long-term timing behavior of \\nss\\ in terms of the usual pulsar timing noise would require timing noise levels much higher than seen in isolated neutron stars with comparable spin-down rates.  The most plausible explanation for the erratic long-term pulse frequency evolution of \\nss\\ is the presence of a binary companion, although the current data do not allow us to place strong constraints on system parameters.  Further efforts at phase-coherent pulse timing observations of \\nss\\ are required to ($i$) unambiguously identify the nature of the long-term pulse frequency variations, and ($ii$) confirm and further investigate the observed short-term phase noise.  Just one additional 40-day observing sequence with higher sampling would yield far better constraints on orbital parameters should \\nss, in fact, possess a binary companion."
    },
    "0608/astro-ph0608160_arXiv.txt": {
        "abstract": "{ We present the detailed spectroscopic analysis of 72 evolved stars, which were previously studied for accurate radial velocity variations. Using one Hyades giant and another well studied star as the reference abundance, we determine the [Fe/H] for the whole sample. These metallicities, together with the \\Teff\\ values and the absolute $V$-band magnitude derived from Hipparcos parallaxes, are used to estimate basic stellar parameters (ages, masses, radii, \\bvo\\ and \\logg) using theoretical isochrones and a Bayesian estimation method. The \\bvo\\ values so estimated turn out to be in excellent agreement (to within $\\sim0.05$~mag) with the observed \\bv, confirming the reliability of the \\Teff--\\bvo\\ relation used in the isochrones. On the other hand, the estimated \\logg\\ values are typically 0.2~dex lower than those derived from spectroscopy; this effect has a negligible impact on \\feh\\ determinations. The estimated diameters $\\theta$ have been compared with limb darkening-corrected ones measured with independent methods, finding an agreement better than 0.3~mas within the $1<\\theta<10$~mas interval (or, alternatively, finding mean differences of just $6$~\\%). We derive the age-metallicity relation for the solar neighborhood; for the first time to our knowledge, such a relation has been derived from observations of field giants rather than from open clusters and field dwarfs and subdwarfs. The age-metallicity relation is characterized by close-to-solar metallicities for stars younger than $\\sim4$~Gyr, and by a large \\feh\\ spread with a trend towards lower metallicities for higher ages. In disagreement with other studies, we find that the \\feh\\ dispersion of young stars (less than 1~Gyr) is comparable to the observational errors, indicating that stars in the solar neighbourhood are formed from interstellar matter of quite homogeneous chemical composition. The three giants of our sample which have been proposed to host planets are not metal rich; this result is at odds with those for main sequence stars. However, two of these stars have masses much larger than a solar mass so we may be sampling a different stellar population from most radial velocity searches for extrasolar planets. We also confirm the previous indication that the radial velocity variability tends to increase along the RGB, and in particular with the stellar radius. }  \\authorrunning{L. da Silva  et al.} \\titlerunning{Basic physical parameters of Giant stars} ",
        "introduction": "It has recently become evident from high accuracy radial velocity (RV) measurements that late-type (G and K) giant stars are RV variables (see e.g. Hatzes \\& Cochran 1993, 1994; Setiawan et al. 2003, 2004). These variations occur on two greatly different timescales: short term variability with periods in the range 2--10 days, and long term variations with periods greater than several hundreds of days. The short term variations are due to stellar oscillations. The cause of the long term variability is not clear and at least three mechanisms, i.e., low mass companions, pulsations and surface activity, have been proposed. The first results of our long term study of the nature of the variability of K giants have shown that all three mechanisms are likely contributors to long term RV variability, although their dependence on the star's fundamental characteristics remain unknown (Setiawan et al. 2003, 2004).  The purpose of this paper is to accurately determine the radii, temperatures, masses, and chemical composition of the stars of our sample, in order to understand better how RV variability and stellar characteristics are related. We are particularly interested in determining to what extent the stars found to host planetary system class bodies (Setiawan et al. 2003a, 2005) exhibit high metallicity, as has been found for the dwarfs hosting giant exoplanets ( Gonzales 1997; Santos et al. 2004).  The present sample is composed of the stars published by Setiawan et al. (2004). The selection criteria and the observations were explained there. We recall that observations were obtained with the FEROS spectrograph at the ESO 1.5m telescope (Kaufer et al. 1999), with a resolving power of 50000 and a signal-to-noise ratio exceeding 150 in the red part of the spectra. ",
        "conclusions": "\\subsection{RV variation }  Setiawan et al. (2004) have studied the trend of radial velocity (RV) variability -- determined by the standard deviation $\\sigma_{\\rm RV}$ -- with $M_V$, detecting an increase of variability along the RGB. We can now do the same for other stellar parameters. Figure~\\ref{hrd} shows our star sample in the HR diagram obtained from the information of Table~\\ref{tabelao}. We note that the $\\sigma_{RV}$ represents the total standard deviation without regard to the timescales involved (i.e. short- versus long term variations). The 11 binaries of the sample which are reported by Setiawan et al. (2004) are excluded from this plot. There is a trend of increasing $\\sigma_{\\rm RV}$ with position along the RGB, which becomes more evident if we exclude the stars that host sub-stellar companion candidates (marked with different symbols).  A similar trend is seen in the plot of $\\sigma_{\\rm RV}$ against stellar radius in Fig.~\\ref{rv_r}. The two stars suspected to host brown dwarf companions, HD~27256 and HD~224533, stand out as showing too large a $\\sigma_{\\rm RV}$ for their radii. If these two stars are excluded, an unweighted least squares fit to the data results in the relation \\[ \\sigma_{\\rm RV} ({\\rm m\\,s}^{-1}) = (1.76\\pm0.31)\\,R(R_\\odot) + (38.1\\pm6.0) \\,. \\] The Spearman rank correlation coefficient is 0.45. This result, although supporting the increase of $\\sigma_{\\rm RV}$ with $R$, is not significant because $\\sigma_{\\rm RV}$ is limited by the long-term precision of our RV measurements (about 25~${\\rm m\\,s}^{-1}$, see Setiawan et al. 2003, 2004).  \\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics{5105_17.ps}} \\end{center} \\caption{ HR diagram of our star sample, excluding 11 binaries. Stellar parameters were taken from Table~\\protect\\ref{tabelao}. Each circle diameter is proportional to $\\sigma_{\\rm RV}$, ranging from $\\sim30$ to 280~m\\,s$^{-1}$. Stars that host substellar companions are marked with full squares for suspected brown dwarfs and full triangles for suspected giant planets. } \\label{hrd} \\end{figure}  \\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics{5105_18.ps}} \\end{center} \\caption{ $\\sigma_{\\rm RV}$ versus stellar radius diagram for our sample, excluding 11 binaries. Refer to Fig.~\\protect\\ref{hrd} for the meaning of the symbols. } \\label{rv_r} \\end{figure}   \\subsection{Distribution of metallicity }  Another major result of this study is the derivation of the age--metallicity relation for the Solar Neighbourhood (see Fig.~\\ref{amr}). Its behaviour agrees, in general, with most previous determinations in the literature, except for the very low \\feh\\ spread that we find for the youngest ages, which is comparable to the observational error. The main novelty with respect to previous results is that we derive the AMR {\\em using data for field red giants only}, whereas the majority of present-day determinations have used samples of field dwarfs (including just a small fraction of subgiants), or giants belonging to open clusters. It worth noticing that we used {\\em very simple data} in our determinations, namely the $BV$ photometry together with Hipparcos parallaxes and measured values for \\Teff\\ and \\feh\\ .  Of course, the same work can be extended to all giants in Hipparcos catalog, once we have obtained homogeneous \\Teff\\ and \\feh\\ determinations for them. This opens the possibility of improving considerably the statistics and reliability of the local age--metallicity relation, simply by acquiring spectroscopic data for a larger sample of bright giants, and performing the same abundance and parameter analysis as in the present work.  Moreover, similar methods can be applied to nearby galaxies with well-known distances, once we have available both the photometry and spectroscopy for a sufficiently large number of their red giants. Zaggia et al. (in preparation) use a procedure similar to ours to derive the age--metallicity relation of the Sgr dSph galaxy.   \\subsection{Stars hosting low mass companions and planets }  In the course of earlier studies we have identified three stars as candidates to host low mass companions: two with planets (HD~47536 and HD~122430) and one with either a brown dwarf or a planet companion (HD~11977). We investigate a range of stellar masses larger than the range of masses usually investigated with radial velocity techniques. Not only can we provide better constraints on the companion mass, but we can also investigate to what extent the conditions for companion formation differ within the mass range. Only one of our three stars, HD~122430, has nearly solar metallicity ($\\feh=-0.05$), while HD~11977 is slightly sub-solar ($\\feh=-0.21$) and HD~47536 ($\\feh=-0.68$) is the most metal poor star of the sample. Schuler et al. (2005) derived a metallicity of [Fe/H] = -0.58 for the K giant star hosting planet HD 13189. Although we are considering a small number of objects, this result seems at odds with what has been found for dwarf stars hosting giant exoplanets, which are preferentially metal rich (e.g. Santos 2004). However, most RV planet search programs have concentrated on solar mass stars and two of our planet hosting giant stars have masses considerably larger than solar. In the case of HD 13189 the host star has a mass of 3.5 M$_\\odot$ (Schuler et al. 2005). At the present time is is unknown what role stellar mass plays in the process of planet formation and for massive stars this may be a more dominant factor than the metallicity. Any investigation of the metallicity-planet relation among giant stars should focus on those in the same mass range. It may be that for a given mass range stars with higher metal abundances still tend to host a higher frequency of giant planets."
    },
    "0608/astro-ph0608210_arXiv.txt": {
        "abstract": "Millisecond pulsars represent an evolutionarily distinct group among rotation-powered pulsars. Outside the radio band, the soft X-ray range ($\\sim 0.1$--10 keV) is most suitable for studying radiative mechanisms operating in these fascinating objects. X-ray observations revealed diverse properties of emission from millisecond pulsars. For the most of them, the bulk of radiation is of a thermal origin, emitted from small spots (polar caps) on the neutron star surface heated by relativistic particles produced in pulsar acceleration zones. On the other hand, a few other very fast rotating pulsars exhibit almost pure nonthermal emission generated, most probably, in pulsar magnetospheres. There are also examples of nonthermal emission detected from X-ray nebulae powered by millisecond pulsars, as well as from pulsar winds shocked in binary systems with millisecond pulsars as companions. These and other most important results obtained from X-ray observations of millisecond pulsars are reviewed in this paper, as well as results from the search for millisecond pulsations in X-ray flux of the radio-quite neutron star RX~J1856.5--3754. ",
        "introduction": "Millisecond pulsars (MSPs) significantly differ in properties from other (ordinary) radio pulsars. First of all, MSPs possess very short and stable spin periods, $P\\lapr 50$ ms, with extremely small period derivatives, $\\dot{P}\\lapr 10^{-18}$ s s$^{-1}$. These two main parameters separate MSPs from the majority of other pulsars, as illustrated in the $P$--$\\dot{P}$ diagram\\footnote{ Based on the pulsar catalog provided by the Australia Telescope National Facility (Manchester et al. 2005) and available at {\\tt http://www.atnf.csiro.au/research/pulsar}.} shown in Figure~1. According to the conventional pulsar magnetic-braking model, MSPs are very old neutron stars, with characteristic ages $\\tau=P/2\\dot{P}\\sim 0.1$--10 Gyr, and low surface magnetic fields $B_{\\rm surf}\\simeq 3.2\\times 10^{19}\\,(P\\dot{P})^{1/2}\\lapr 10^{10}$ G (see Fig.~1). They were presumably spun up by angular momentum transfer during a mass accretion phase in binary systems\\footnote{It is why MSPs are often called ``recycled'' pulsars.} (e.g., Alpar et al. 1982), and their low values of $B_{\\rm surf}$ could be explained by the Ohmic and/or accretion-induced decay of magnetic field (see Cumming 2005 and references therein). Discoveries of several accretion-driven binary MSPs, first of all the famous pulsar SAX J1808.4--3658 with $P\\simeq 2.5$ ms, support this hypothesis (see, e.g., Wijnands 2004 for a review on accreting MSPs).  Since the discovery of the first fast rotating pulsar B1937+21 by Backer et al. (1982), MSPs have been extensively searched for and studied in radio domain. Currently, about 130 MSPs are known (Manchester et al. 2005). Outside the radio band, as MSPs are intrinsically faint at optical wavelengths and most of them ($\\sim 80$\\%) reside in binary systems with optically brighter white dwarf companions, the soft X-ray energy range ($\\sim 0.1$--10 keV) is the main source of information on these pulsars. The detection of pulsed X-ray emission from the brightest (and nearest) MSP J0437--4715 with $ROSAT$ (Becker \\& Tr\\\"umper 1993) initiated a series of dedicated X-ray observations of these intriguing objects in 90's with this satellite and also $ASCA$, Beppo$SAX$ and $RXTE$. Later on this observational ``relay'' has been continued with \\chan\\ and \\xmm. So far, firm X-ray detections have been reported for about three dozens of isolated (solitary, or non-accreting if in binaries) MSPs. The majority of these objects are located in globular clusters (mostly 47~Tuc --- see Bogdanov et al. 2006). This paper mainly concentrates on eight MSPs, listed in Table~1, with available detailed information on properties of detected X-ray emission. Besides $P$ and $\\tau$, Table~1 gives estimates on  the distances to these objects, $d$ (inferred from either pulsar parallaxes or dispersion measures and the NE2001 galactic electron density model by Cordes \\& Lazio 2003)\\footnote{ For the MSPs of the second group in Table~1 (with $d<1$ kpc), the distance estimates obtained from the pulsar parallaxes and those inferred from their dispersion measures agree well with each other (see Hotan, Bails \\& Ord 2006 and Lommen et al. 2006).}, and on their spin-down energies, $\\ed=4\\pi^2 I P^{-3}\\dot{P}$ (assuming a standard neutron star moment of inertia, $I=10^{45}$ g cm$^2$). Note that for PSR J1024--0719 the intrinsic period derivative is not well determined, $\\dot{P}<3\\times 10^{-21}$ s s$^{-1}$ (Hotan, Bails \\& Ord 2006), that results in the lower and upper limits on $\\tau$ and $\\ed$, respectively.  Generally, X-ray emission from radio pulsars consists of two different components, thermal and nonthermal, generated on the neutron star surface or in its vicinity. The nonthermal component is usually described by a power-law (PL) spectral model and attributed to radiation produced by synchrotron and/or inverse Compton processes in the pulsar magnetosphere, whereas the thermal emission can originate from either the whole surface of a cooling neutron star or small hot spots around the magnetic poles (polar caps; PCs) on the star surface, or both. As predicted by virtually all pulsar models, these PCs can be heated up to X-ray temperatures ($\\sim 1$ MK) by relativistic particles generated in pulsar acceleration zones. A conventional assumption about the PC radius is that it is close to the radius within which open magnetic field lines originate from the pulsar  surface, $\\sim [2\\pi R_{\\rm NS}^3/cP]^{1/2}\\simeq 2\\,[P/5\\,{\\rm ms}]^{-1/2}$ km (for a neutron star radius $R_{\\rm NS}=10$ km). In case of MSPs, the entire surface at a neutron star age of $\\sim 1$ Gyr is too cold, $\\lapr 0.1$ MK, to be detectable in X-rays (although it may be seen in the $UV/FUV$ band --- see Sec.~6). Therefore, only nonthermal (magnetospheric) and/or thermal PC radiation is expected to be observed in X-rays from these objects. In addition to the radiation produced by MSPs themselves, nonthermal emission from pulsar-wind nebulae (PWNe) associated with MSPs moving at supersonic velocities ($\\gapr 100$ km s$^{-1}$) through interstellar medium may be detected. Another source of nonthermal X-ray radiation generated in binary systems can be an intrabinary shock formed where the pulsar wind and matter from the stellar component collide (Arons \\& Tavani 1993), although such a component would be hardly separated from radiation produced by the pulsar itself (unless properties of the nonthermal emission varies with orbital phase).  The MSPs in Table~1 belong to two distinct groups: those which emit almost pure nonthermal radiation and those with a predominantly thermal PC component. The next two Sections provides details on the X-ray properties of these objects. Section~4 is devoted to three X-ray emitting MSPs whose properties remain uncertain. X-ray PWNe powered by MSPs are briefly discussed in Section~5. A summary is given in Section~6.  \\begin{figure} \\centering \\includegraphics[height=7cm,angle=0]{zavlin_f2.eps} \\caption{X-ray spectrum of PSR B1821--24 as detected with the \\chan\\ ACIS-S instrument (crosses) fitted with a PL model of $\\Gamma=1.2$ (dotted curves) plus a possible PC component (dashes). See also Becker et al. (2003). } \\label{fig:2} \\end{figure}  \\begin{figure} \\centering \\includegraphics[height=7cm,angle=0]{zavlin_f3.eps} \\caption{Same as in Fig.~2 for PSR B1937+21 and a PL model of $\\Gamma=1.2$ (see also Kuiper et al. 2006). } \\label{fig:3} \\end{figure}  \\begin{figure} \\centering \\includegraphics[height=7cm,angle=0]{zavlin_f4.eps} \\caption{Same as in Fig.~2 for PSR J0218+4232 (as detected with the \\xmm\\ EPIC-MOS instruments) and a PL model of $\\Gamma=1.1$ (see also Webb, Olive \\& Barret 2004). } \\label{fig:4} \\end{figure}  \\begin{figure} \\centering \\includegraphics[height=7cm,angle=0]{zavlin_f5.eps} \\caption{Same as in Fig.~2 for PSR B1957+20 and a PL model of $\\Gamma=1.9$. (Note that the detected flux can contain a contribution from a possible intrabinary shock --- see Sec.~2 and Stappers et al. 2003 for more details.) } \\label{fig:5} \\end{figure} ",
        "conclusions": ""
    },
    "0608/astro-ph0608026_arXiv.txt": {
        "abstract": "We give an overview of our current understanding of the structure of gas giant planets, from Jupiter and Saturn to extrasolar giant planets.  We focus on addressing what high-pressure laboratory experiments on hydrogen and helium can help to elucidate about the structure of these planets. ",
        "introduction": "In order to understand the formation of giant planets, and hence, the formation of planetary systems, we must be able to determine the interior structure and composition of giant planets.  Jupiter and Saturn, our solar system's gas giants, combine to make up 92\\% of the planetary mass of our solar system.  Interestingly, knowledge of only a few key quantities allows us to gain important insight into their interior structure.  The equation of state of hydrogen, together with measurements of the mass, radius, and oblateness of Jupiter and Saturn is sufficient to show that these planets are hydrogen-helium rich objects with a composition similar to that of the Sun \\citep{Demarcus58}.  Furthermore, estimates of the transport coefficients of dense metallic hydrogen and the observation that Jupiter emits more infrared radiation than it absorbs from the Sun \\citep{Low66}, is sufficient to show that gas giant planet interiors are warm, fluid, and convective, not cold and solid \\citep{Hubbard68}.  It has also been clear for some time that the composition of Jupiter and Saturn is not exactly like that of the Sun---these planets are enhanced in ``heavy elements'' (atoms more massive than helium), compared to the Sun \\citep{Podolak74}.  An understanding of how these planets attained these heavy elements, and their relative ratios, can give us a wealth of information on planetary formation and the state of the solar nebula.  Looking beyond Jupiter and Saturn, we now have 200 extrasolar giant planets (EGPs) that have been found to orbit other stars.  A subclass of these planets are the ``hot Jupiters'' that orbit their parent stars at around 0.05 AU.  To date, ten planets (with masses from 0.36 to 1.5 \\mj) have been seen to transit their parents stars.  All of these objects are hot Jupiters, with orbital periods of only a few days \\citep[see][]{Charb06}.  These transiting planets are important because we can measure their masses \\emph{and} radii, thereby allowing us access to information on their interior structure \\citep{Guillot05}.  While our understanding of the interiors of these planets will never be as detailed as that for Jupiter and Saturn, we will eventually have a very large sample of these transiting objects at various masses, compositions, and orbital distances, which will allow for an understanding of the mass--radius relation for giant planets under a variety of conditions.  By far the most important physical input into giant planet structural models is the equation of state (EOS) of hydrogen.  The decade of pressure that is most important for understanding the interiors of giant planets is 1-10 Mbar (100-1000 Gpa) \\citep{Saumon04}.  In the past decade experiments have been able to probe into the lower end of this pressure range \\citep{Weir96, Collins98, Knudson01, Boriskov05}.  In this paper, instead of focusing on equation of state physics we will focus on key questions for understanding the structure and composition of giant planets.  As we discuss giant planet interiors we will investigate how high pressure laboratory experiments have and will continue to allow us to better answer these questions. ",
        "conclusions": ""
    },
    "0608/astro-ph0608557_arXiv.txt": {
        "abstract": "Using UV spectra obtained with {\\it FUSE}, {\\it HST}, and/or {\\it IUE} together with higher resolution optical spectra, we determine interstellar column densities of $^{12}$CO, $^{13}$CO, and/or C$_2$ for ten Galactic sight lines with $E(B-V)$ ranging from 0.37 to 0.72. The $N$(CO)/$N$(H$_{2}$) ratio varies over a factor of 100 in this sample, due primarily to differences in $N$(CO). For a given $N$(H$_2$), published models of diffuse and translucent clouds predict less CO than is observed. The $J$ = 1--3 rotational levels of $^{12}$CO are sub-thermally populated in these sight lines, with $T_{ex}$ typically between 3 and 7 K. In general, there appears to be no significant difference between the excitation temperatures of $^{12}$CO and $^{13}$CO. Fits to the higher resolution CO line profiles suggest that CO (like CN) is concentrated in relatively cold, dense gas.  We obtain C$_2$ column densities from the F-X (0-0) band at 1341 \\AA\\ (three sight lines; $J$ = 0--12), the F-X (1-0) band at 1314 \\AA\\ (one sight line; $J$ = 0--12), the D-X (0-0) band at 2313 \\AA\\ (four sight lines; $J$ = 0--18), and the A-X (3-0) and (2-0) bands at 7719 and 8757 \\AA\\ (seven sight lines; $J$ = 0--12). Comparisons among those column densities yield a set of mutually consistent band $f$-values for the UV and optical C$_2$ bands, but also reveal some apparent anomalies within the F-X (0-0) band. Both the kinetic temperature $T_{\\rm k}$ inferred from the C$_2$ rotational populations and the excitation temperature $T_{02}$(C$_2$) are generally smaller than the corresponding $T_{01}$(H$_2$) --- suggesting that C$_2$ is concentrated in colder, denser gas than H$_2$.  Incorporating additional column density data for \\ion{K}{1}, HD, CH, C$_2$, C$_3$, CN, and CO from the literature (for a total sample of 74 sight lines), we find that (1) CO is most tightly correlated with CN; (2) the ratios $^{12}$CO/H$_2$ and $^{13}$CO/H$_2$ both are fairly tightly correlated with the density indicator CN/CH (but C$_2$/H$_2$ is not); and (3) the ratio $^{12}$CO/$^{13}$CO is somewhat anti-correlated with both CN/CH and $N$(CO). Sight lines with $^{12}$CO/$^{13}$CO below the average Galactic value of $^{12}$C/$^{13}$C appear to sample colder, denser gas in which isotope exchange reactions have enhanced $^{13}$CO, relative to $^{12}$CO. ",
        "introduction": "\\label{sec-intro}  Apart from H$_2$ and (in some cases) HD, CO is the most abundant interstellar molecule. Determinations of the abundance and rotational excitation of the dominant $^{12}$C$^{16}$O and its isotopomers thus can provide significant information about the physical conditions in Galactic and extragalactic molecular clouds. The CO/H$_{2}$ abundance ratio is often assumed to be roughly constant in dense molecular gas, with a ``canonical'' value of order 10$^{-4}$ derived primarily from millimetric and radio observations of CO and inferred abundances of H$_{2}$ in Galactic molecular clouds (Young \\& Scoville 1991). Under that assumption, measurements of CO are used to estimate the total H$_2$ content when no data are available for H$_2$. In more diffuse gas, both H$_2$ and CO were detected directly via their absorption bands in the far-UV using spectra from the {\\it Copernicus} satellite; longer wavelength UV bands of CO were observed with {\\it IUE}. Analysis of the {\\it Copernicus} spectra showed that the CO/H$_{2}$ ratio varied by two orders of magnitude (from about 5 $\\times$ 10$^{-8}$ to 5 $\\times$ 10$^{-6}$) (Federman et al. 1980) --- so that CO is not a good predictor for H$_{2}$ in more diffuse material. Theoretical models of cloud chemistry suggested that such large variations in the CO/H$_{2}$ ratio could be expected in diffuse ($A_{\\rm V}$ $\\la$ 1 mag) and translucent (1 mag $\\la$ $A_{\\rm V}$ $\\la$ 5 mag) regions still permeated by dissociating photons (e.g., van Dishoeck \\& Black 1986a, 1988) as well as in regions perturbed by the recent passage of dissociative shocks (Neufeld \\& Dalgarno 1989). The relatively low sensitivity and resolution of {\\it Copernicus} and {\\it IUE}, however, allowed only limited investigation of more heavily reddened sight lines, where the CO/H$_2$ ratio might be expected to approach the dense cloud value.  Determining observationally the population distribution in the individual rotational levels of any molecular species can yield constraints on the physical properties characterizing the molecular gas (e.g., van Dishoeck et al. 1991; Federman et al. 1997b; Wannier et al. 1997). Observations of the rotational excitation of C$_2$, for example, have yielded estimates for the temperatures and densities in diffuse molecular clouds (e.g., van Dishoeck \\& de Zeeuw 1984); data for CN have been used both to constrain the temperature of the cosmic microwave background (e.g., Roth \\& Meyer 1995) and to estimate local hydrogen densities (e.g., Black \\& van Dishoeck 1991). The pervasive abundance of the CO and H$_{2}$ molecules in the interstellar gas make those species especially important tools for investigations of the gas properties. Unfortunately, however, the spectrographs on {\\it Copernicus} and {\\it IUE} could not resolve the rotational structure of the CO bands.  Isotopic ratio measurements provide another means to investigate the chemical processes occurring in the interstellar medium. Carbon, whose isotopes $^{12}$C and $^{13}$C are products of primary and secondary nucleosynthesis processes (Langer \\& Penzias 1990), respectively, is well suited for such studies. Within about 1--2 kpc of the Sun, the current average $^{12}$C/$^{13}$C ratio is about 70 (e.g., Wannier et al. 1982; Stahl \\& Wilson 1992; Lucas \\& Liszt 1998), with some variations; there also is an apparent systematic decrease in the ratio toward the Galactic center (e.g., Wilson 1999; Milam et al. 2005). When determined from molecules unaffected by chemical fractionation (e.g., CH$^{+}$), the $^{12}$C/$^{13}$C ratio is an indicator of Galactic stellar evolution. The (optical) absorption lines from $^{13}$CH$^+$, however, are typically very weak. Because chemical fractionation can occur for CO, the $^{12}$CO/$^{13}$CO ratio gives additional information on chemical processes in the molecular gas (e.g., Bally \\& Langer 1982; Sheffer et al. 1992). Again, however, the limited sensitivity of {\\it Copernicus} and {\\it IUE} made it difficult to obtain reliable measures of the generally weak $^{13}$CO absorption.  The higher far-UV sensitivity of the {\\it FUSE} instrument (Moos et al. 2000), together with the higher sensitivity and resolution of the {\\it HST} GHRS and STIS echelle modes in the near-UV, however, has allowed us to revisit these issues. The {\\it FUSE} spectra contain numerous absorption lines from the Lyman and Werner bands of molecular hydrogen, from which accurate total H$_{2}$ column densities can be derived even toward fainter, more heavily reddened stars (e.g., Rachford et al. 2002). While the B-X (0-0) (1150 \\AA), C-X (0-0) (1087 \\AA), and E-X (0-0) (1076 \\AA) bands arising from the $X ^{1}\\Sigma^{+}$ ground electronic state of CO are also present within the {\\it FUSE} wavelength coverage, an accurate determination of the CO column densities from the {\\it FUSE} data alone is often hindered by line blending, saturation, and unresolved structure in the absorption profiles. With GHRS and STIS, however, one could observe weaker bands in the CO A-X system between 1240 and 1550 \\AA\\ at high enough resolution and S/N to discern and measure the individual rotational ($J$) levels of the bands --- allowing accurate estimates of their individual column densities. Moreover, some of the weaker CO spin-forbidden (triplet-singlet) intersystem bands between 1300 and 1550 \\AA\\ can be detected once the permitted bands become saturated (Morton \\& Noreau 1994) --- enabling more reliable determinations of higher CO column densities. Several recent studies, for example, have used UV spectra from {\\it HST} and/or {\\it FUSE} to determine accurate $N$(CO) toward the moderately reddened stars $\\zeta$ Oph (Lambert et al. 1994), $\\rho$ Oph and $\\chi$ Oph (Federman et al. 2003), X Per (Sheffer, Federman, \\& Lambert 2002a; Sheffer, Lambert, \\& Federman 2002b), and HD 203374A (Sheffer, Federman, \\& Andersson 2003); Pan et al. (2005) have reported $N$(H$_2$) and $N$(CO) toward a number of stars in Cep OB2 and Cep OB3. High-resolution GHRS spectra of the C$_2$ D-X (Mulliken) (0-0) band toward $\\zeta$ Oph have enabled measurements of the C$_2$ rotational levels up to $J$ = 24 (Lambert, Sheffer, \\& Federman 1995) --- higher $J$ than had been obtainable from the weaker A-X (Phillips) bands in the optical and near-IR.  In this paper, we present a study of the $^{12}$CO, $^{13}$CO, and/or C$_{2}$ molecules toward ten moderately reddened [$E(B-V)$ = 0.37--0.72] Galactic stars, based on UV spectra obtained with {\\it FUSE}, {\\it IUE}, and/or {\\it HST}; on higher resolution optical spectra of CH, CN, and \\ion{K}{1}; and on very high S/N optical spectra of C$_2$. This combination of UV and optical data has allowed us to perform detailed analyses of the CO and C$_{2}$ profiles and to derive column densities for the individual rotational levels of the ground electronic and vibrational states --- up to $J$ = 3 for $^{12}$CO, $J$ = 2 for $^{13}$CO, and $J$ = 18 for C$_2$. In \\S~\\ref{sec-obsred}, we discuss the observational data and the data reduction procedures. In \\S~\\ref{sec-analysis}, we describe the methods used to determine column densities from the spectra. In \\S~\\ref{sec-phys}, we combine our new results with existing literature data for H$_2$, HD, CH, C$_2$, C$_3$, CN, and CO (for a total sample of 74 stars) and discuss the physical conditions characterizing the diffuse molecular gas using this expanded star sample. Column densities of H$_2$, CH, C$_2$, CN, $^{12}$CO, and $^{13}$CO for that total sample are tabulated in the Appendix. In \\S~\\ref{sec-summ}, we summarize the main results of this study. ",
        "conclusions": "\\label{sec-summ}  Archival UV spectra from {\\it HST} (GHRS and STIS), {\\it FUSE}, and/or {\\it IUE}, together with high and medium resolution optical spectra, were analyzed to study the molecular abundances and physical conditions in the molecular gas along ten translucent sight lines. Absorption from $^{12}$CO was detected in the far-UV E-X (0-0) (1076 \\AA), C-X (0-0) (1087 \\AA), and B-X (0-0) (1150 \\AA) bands, in a number of bands in the A-X system (between 1240 and 1550 \\AA), and/or in several of the stronger intersystem bands (between 1360 and 1550 \\AA); corresponding A-X bands of $^{13}$CO were detected in six of the sight lines. Accurate column densities were determined for $^{12}$CO and $^{13}$CO using a combination of curve of growth, apparent optical depth, and profile fitting methods. High-resolution optical spectra of \\ion{K}{1}, CH, and CN were used to determine the velocity structure of the gas along the sight lines in order to account for potential saturation effects in the UV spectra.  The UV C$_2$ F-X (0-0) absorption band at 1341 \\AA\\ was detected in high-resolution STIS echelle spectra of HD~24534, HD~27778, and HD~206267; the weaker F-X (1-0) band at 1314 \\AA\\ was detected toward HD~24534; and the stronger D-X (0-0) band at 2313 \\AA\\ was detected toward HD~24534, HD~27778, HD~147888, and HD~207198. These constitute nearly all of the reported high-resolution spectra of the UV bands of C$_2$. Absorption from the optical C$_2$ A-X (3-0) and/or (2-0) bands at 7719 and 8757 \\AA\\ was detected toward seven stars. Comparisons among the various C$_2$ bands have yielded a (reasonably) mutually consistent set of band $f$-values; the value adopted for the A-X (2-0) band is consistent with several recent theoretical and experimental determinations. There appear to be some discrepancies between the observed and predicted wavelengths and/or $f$-values for (at least) some individual lines in the F-X (0-0) band.  Reliable column densities (or limits) are now available for $^{12}$CO, CN, CH, C$_2$, and H$_2$ for at least 30 diffuse or translucent sight lines (11 with detections of $^{13}$CO; 13 with detections of HD; 7 with detections of C$_3$; 15 with $A_{\\rm V}$ between 1.0 and 2.7 mag) --- enabling investigations of correlations between the various species and providing more stringent tests for models of the structure and chemistry of translucent clouds. Of those molecular species, CN shows the best correlation with CO --- both in overall column density and in component structure. Strong correlations between the ratios $^{12}$CO/H$_2$ and $^{13}$CO/H$_2$ and the density indicator CN/CH suggest that the abundances of both CO isotopomers increase with increasing local density. Differences in the strength of the ambient radiation field (as inferred from H$_2$ rotational excitation) also are likely to affect the CO abundances. For our sample of ten translucent sight lines, the $N$(CO)/$N$(H$_{2}$) ratio varies over a factor of 100, due primarily to differences in $N$(CO). For the total sample, the $^{12}$CO/H$_2$ ratio ranges from 2.5 $\\times$ 10$^{-8}$ to 2.8 $\\times$ 10$^{-5}$; CO is not yet the dominant carbon-containing species in any of the sight lines. The models of van Dishoeck \\& Black (1988, 1989) predict relationships with H$_2$, CN, CH, and C$_2$ that have slopes similar to those obtained in fits to the observed data, but the predicted $N$(CO) generally are lower than the observed values by factors of 2--4. The models of Warin et al. (1996) and Le Petit et al. (2006) predict much less CO than is observed, for a given $N$(H$_2$) --- though more of the parameter space available in the Meudon PDR code should be explored. Inclusion of the effects of inhomogeneous/clumpy structure (e.g., Spaans 1996) and increasing the CO self-shielding and (for the Warin et al. models) the gas-phase carbon abundance should yield better agreement with the observed CO abundances.  Rotational level populations and corresponding excitation temperatures were determined up to $J$ = 3 for $^{12}$CO (nine sight lines) and up to $J$ = 2 for $^{13}$CO (four sight lines). The $^{12}$CO excitation temperatures $T_{01}$ typically are between 3 and 7 K (well below the likely kinetic temperature of the gas); the corresponding values for the higher rotational levels are generally equal to $T_{01}$, within the uncertainties. Some of these excitation temperatures are slightly higher than the values previously reported for eight other comparably reddened sight lines, but they are similar to the values found via mm-wave CO absorption by Liszt \\& Lucas (1998) and to the values predicted for diffuse clouds with $A_{\\rm V}$ $\\sim$ 1.1 mag by Warin et al. (1996). The excitation temperatures $T_{01}$ for $^{13}$CO are equal (within the mutual uncertainties) to the corresponding values for $^{12}$CO --- suggesting that the radiative excitation mechanism proposed by Wannier et al. (1997) may not be significant, in many cases.  Rotational level populations and excitation temperatures were also determined for C$_2$ --- up to $J$ = 12 from the UV F-X (0-0) and (1-0) bands (three sight lines), up to $J$ = 18 from the near-UV D-X (0-0) band (four sight lines), and up to $J$ = 12 from the optical A-X (2-0) and (3-0) bands (seven sight lines). The excitation temperature for the lowest rotational level ($T_{02}$) ranges from about 20 to 45 K for those sight lines where it is reasonably well determined; corresponding values for the higher rotational levels increase with $J$ (especially for $J$ $\\ga$ 8), due to radiative excitation of the upper states. In most cases, the C$_2$ rotational excitation seems consistent with model predictions for kinetic temperatures of 35--55 K and local total hydrogen densities of 200--400 cm$^{-3}$, assuming the typical near-IR interstellar radiation field.  Of the twelve sight lines with measured $N$($^{13}$CO), three have $^{12}$CO/$^{13}$CO ratios significantly above the average Galactic $^{12}$C/$^{13}$C value of 70, five have $^{12}$CO/$^{13}$CO ratios consistent with the average value, and four have $^{12}$CO/$^{13}$CO ratios significantly below the average value. There appears to be an anti-correlation between the $^{12}$CO/$^{13}$CO ratio and the CN/CH ratio --- suggestive of a decrease in $^{12}$CO/$^{13}$CO with increasing local density, as predicted by the models of Warin et al. (1996) for densities in the range 100-1000 cm$^{-3}$. The $^{12}$CO/$^{13}$CO ratio is also anti-correlated with the column density of CO, as found (for a smaller sample) by Liszt \\& Lucas (1998). Sight lines with below average $^{12}$CO/$^{13}$CO ratios also appear to have kinetic temperatures (from analysis of the C$_2$ rotational excitation) less than about 35 K; sight lines with average or above average $^{12}$CO/$^{13}$CO ratios appear to be somewhat warmer. While $^{12}$CO/$^{13}$CO ratios above the average value likely result from selective photodissociation (Federman et al. 2003), which favors the self-shielded $^{12}$CO, the below-average ratios probably are due to isotope exchange reactions, which favor $^{13}$CO in cold gas ($T$ $\\la$ 35 K).  While the H$_2$ excitation temperature $T_{01}$ ranges from about 50 K to 110 K for our ten sight lines, there are several indications that C$_2$ and (especially) CO trace colder, denser gas in many cases. The excitation temperature $T_{02}$ of C$_2$ (which provides an upper limit to the local kinetic temperature for the gas containing the C$_2$) and the $T_{\\rm k}$ inferred from comparisons with models of the excitation are generally lower than $T_{01}$ for H$_2$ (which should equal the kinetic temperature of the gas traced by H$_2$). Fits to the CO line profiles indicate that smaller $b$-values (0.3--0.7 km~s$^{-1}$) are needed than for fits to the profiles of \\ion{K}{1}, CH, and CN; the CO structure also appears more similar to that of CN (which is concentrated in denser regions). The column densities of both $^{12}$CO and $^{13}$CO are most tightly correlated with that of CN, and their abundances, relative to H$_2$, are correlated with the density indicator $N$(CN)/$N$(CH).  These results illustrate the utility of combining information from multiple species sensitive to different combinations of physical parameters in order to probe regions characterized by different properties (along a sight line or within a single ``cloud''). They also emphasize the need for developing models capable of incorporating ranges and/or mixtures of physical conditions in order to accurately assess and predict the abundances of various atomic and molecular species in the diffuse ISM."
    },
    "0608/astro-ph0608411_arXiv.txt": {
        "abstract": "We construct a simple model for radioisotopic enrichment of the protosolar nebula by injection from a nearby supernova, based on the inverse square law for ejecta dispersion. This idealized model concisely and explicitly connects the observational data to the key astrophysical parameters; as such, it provides a useful tool for parameter studies that can inform the computationally expensive numerical (magneto)hydrodynamic studies essential for a full solution to this problem. We find that the presolar radioisotopes abundances (i.e., in solar masses) demand a nearby supernova: its distance $D$ can be no larger than $D/R_{\\rm SS} \\le 66$ times the size $R_{\\rm SS}$ of the protosolar nebula, at a 90\\% confidence level, assuming 1 M$_\\odot$ of protosolar material. The relevant size of the nebula depends on its state of evolution at the time of radioactivity injection. In one scenario, a collection of low-mass stars, including our sun, formed in a group or cluster with an intermediate- to high-mass star that ended its life as a supernova while our sun was still a protostar, a starless core, or perhaps a diffuse cloud. Using recent observations of protostars to estimate the size of the protosolar nebula constrains the distance of the supernova to $D \\sim 0.02$ pc to 1.6 pc. The supernova distance limit is consistent with the scales of low-mass stars formation around one or more massive stars, but it is closer than expected were the sun formed in an isolated, solitary state. Consequently, if {\\em any} presolar radioactivities originated via supernova injection, we must conclude that our sun was a member of such a group or cluster that has since dispersed, and thus that solar system formation should be understood in this context. In addition, we show that the timescale from explosion to the creation of small bodies was on the order of 1.8 Myr (formal 90\\% confidence range of 0 to 2.2 Myr), and thus the temporal choreography from supernova ejecta to meteorites is important. Finally, we can not distinguish between progenitor masses from 15 to 25 M$_\\odot$ in the nucleosynthesis models; however, the 20 M$_\\odot$ model is somewhat preferred. ",
        "introduction": "The picture of star formation that has been developed over the last few decades provides a good broad-brush explanation of isolated, low-mass, Sun-like star formation \\citep[e.g.][]{shuppiii}. However, is that picture valid for the formation of the Sun? The scenario is probably unlikely as there is good evidence that the majority of Sun-like stars form in clusters \\citep[e.g.,][]{carpenter2000,ladalada2003}. In addition, the theoretical view that protostars form in relative isolation from their molecular cloud neighbors is also weakening \\citep[e.g.,][]{clark2005}.  To better explain the origin of the Sun and Planets, it is important to first understand the initial conditions in which the Sun formed.  Did the Sun form in isolation or in a small group or cluster of stars that has since dispersed? The concept of a lonely Solar System birth has been challenged by the idea that the Sun was born near at least one massive star, in a group or cluster of stars, based on the study of short-lived radioisotopes in meteorites \\citep[e.g.][]{hester2004}.  The main evidence for short-lived radioactive species is from isotopic anomalies of daughter species in primitive meteorites \\citep[for recent reviews, see][]{goswami2000,meyer2000,wadhwa2006}. The isotopes are now extinct, but they must have been live, and in relatively high abundances,  during the early stages of planet formation in our Solar System. Possible origins for the measured isotopic abundances are that the Sun formed in a location that was enriched either by a stellar wind or a recent core-collapse supernova \\citep[e.g.][]{lee1977}, hence in or near a cluster or stellar group environment. Other possibilities for the short-lived radioactive isotopes include irradiation in the circumstellar disk \\citep[e.g.,][]{lee1998} or mixing of ejecta from core-collapse supernova in the interstellar medium (ISM) \\citep[e.g.,][]{meyer2000,hester2005,oue2005}. However, new results reveal that the early Solar System contained significant amounts of the short-lived $^{60}$Fe radionuclide \\citep[e.g.][]{tachibana2003,most2005} that can not be explained with the irradiation models or ISM mixing models \\citep[e.g.][]{tachibana2003,gounelle2006}.  In this paper, we make the assumption that a supernova event occurred near the early solar nebula, such that short-lived radioactive species were injected into the solar nebula \\citep[e.g.][]{cameron1977,goswami2000}. Although, the $^{60}$Fe abundance could also be explained by ejecta from Asymptotic Giant Branch (AGB) stars, novae, or Type I supernovae, the likelihood of such late stellar stage objects overlapping with a star formation region is much smaller than a supernova that was born in the same cluster \\citep[e.g.][]{kastner1994,hester2005}.  With the assumption of a nearby supernova, we construct a simple model that directly and explicitly relates the observed pre-solar radioisotope abundances to the supernova distance and the size of the protosolar nebula into which they were injected. Our model, while idealized, clearly illustrates that the radioisotope abundances directly constrain (and in particular, place an upper limit on) the ratio of the supernova distance to the size of the nebula. This type of model has been discussed before \\citep[e.g.][]{vanhala2000,vanhala2002,oue2005} with more emphasis on the detailed hydrodynamics of the injection of supernova material into the nebula. In this paper, we focus on a statistical parameter study, with emphasis on the supernova distance $D$. Our simple model allows for a rapid search of the parameter space and shows that in most scenarios, a nearby supernova (i.e., $D \\le$  a few pc, and $\\ll$ 1 kpc) is required if {\\em any} radioisotopes came from a supernova. This has important qualitative implications for the nature of the Sun's formation, and quantitatively our results can inform and calibrate the more computationally expensive hydrodynamics calculations needed to solve the problem in all of its detail. ",
        "conclusions": "We present a simple model relating the pre-solar abundances of short-lived radioisotopes to the properties of the injecting supernova event. While our model is idealized, it is also explicit and computationally simple, yielding a parameter study whose general aspects are intriguing. In general, we are driven to conclude that a very nearby supernova event must be invoked if {\\em any} presolar radioactivities are due to injection from a single explosion. Specifically, to a formal confidence level of 90\\%, the distance from the Solar Nebula and the supernova was in the range of 0.1 to 1.6 pc, for siblings and up to approximately 4 pc (in the most conservative, and unlikely, case of perfect injection) for triggered birth, using the \\cite{rauscher02} nucleosynthesis models for 15 to 25 M$_\\odot$ progenitors. This sounds surprisingly close, but it is consistent with typical distances found for low-mass stars clustering around one or more massive stars. We posit that our Sun was a member of such a cluster that has since dispersed.  It is important to note that our distance estimates imply that the Sun was directly influenced by the massive star's photoionization effects, much like proplyds in Orion.  To understand the Sun's formation, one must include the large role played by such a complicated environment. The minimum distance is also interesting, as it is similar to the distance at which the disk would probably be destroyed $<$ 0.25 pc \\citep{chevalier2000}. Our estimated distance range includes the expected uncertainties of the nucleosynthesis models and the radioisotopic measurements. Although the physical parameters of the solar nebula also play a large role in these estimates, we use nearly all protostellar evolutionary states to compile our range.  In addition to the distance of the supernova, we constrain the time scale of the explosion to the creation of small bodies to a 90\\% confidence range of 0 to 2.2 Myrs. The majority of this time was probably spent in our solar system during the building of small objects from the circumstellar dust; in fact, it is consistent with the estimates for chondrule production. It is important to note that the temporal choreography from supernova ejecta to meteorites is crucial. The upper limits of our fit decay times make it possible for the supernova to have triggered the fast collapse of our Sun, seeded it with short-lived isotopes, and created CAIs, but it seems more likely that the Sun was already a protostar during the explosion \\citep[e.g.,][]{hester2005}.  In addition to the inherent uncertainties in our model, there are additional assumptions that would imply a smaller or larger supernova distance. There are two clear cases that would cause our deduced distances to be an upper limit: (1) an inclined circumstellar disk (Class I/II source) would have a lower cross-sectional area and not collect the full supernova contaminates as a face-on disk, and (2) the ejecta may not be effectively stopped by the circumstellar disk as assumed ($\\finj \\ll 1$) \\citep[i.e., effectiveness of the injection process; see][]{vanhala2002}. Similarly, there are two opposing cases that would make our deduced distances lower limits: (1) highly inhomogeneous ejecta \\citep[i.e. bullets as seen in Cas A,][]{willingale2003} would not allow a predictable distance, and (2) although the material is injected uniformly onto the circumstellar disk, the inner regions of the disk has a higher density than the outer regions \\citep[e.g.,][]{mundy1996}, so if there is no mixing, the measured meteorites are produced {\\it in situ} with abundances lower than expected. These crucial issues lie within the domain of sophisticated (magneto)hydrodynamical simulations that will yield a fuller understanding of this problem; our simple model serves as a tool that highlights and quantifies the issues that the full numerical models address.  In closing, we note that the analysis presented here for extinct presolar meteoritic radioactivities closely parallels considerations of live geological radioactivities as signatures of relatively recent nearby supernovae \\citep[e.g.,][]{efs,knie}. Both analyses leverage the uniqueness of radioisotopes as a signature of supernova activity, and both use the inverse square law in the same way to gauge distance. Moreover, \\citet{fhe} emphasize that given a confirmed nearby supernova event, one can then turn the problem around, reading the presence or absence of, and detailed abundances of, all radioisotopes as direct probes of supernova ejecta and thus nucleosynthesis. This adds another way that meteoritic anomalies serve as fossil testimony to our massive sibling."
    },
    "0608/astro-ph0608627_arXiv.txt": {
        "abstract": "{} {A huge multiwavelength campaign targeting the blazar \\object{AO 0235+164} was organized by the Whole Earth Blazar Telescope (WEBT) in 2003--2005 to study the variability properties of the source.} {Monitoring observations were carried out at cm and mm wavelengths, and in the near-IR and optical bands, while three pointings by the XMM-Newton satellite provided information on the X-ray and UV emission.} {We present the data acquired during the second observing season, 2004--2005, by 27 radio-to-optical telescopes. The $\\sim 2600$ data points collected allow us to trace the low-energy behaviour of the source in detail, revealing an increased near-IR and optical activity with respect to the previous season. Increased variability is also found at the higher radio frequencies, down to $\\sim 15$ GHz, but not at the lower ones. While the X-ray (and optical) light curves obtained during the XMM-Newton pointings reveal no significant short-term variability, the simultaneous intraday radio observations with the 100 m telescope at Effelsberg show flux-density changes at 10.5 GHz, which are more likely due to a combination of intrinsic and extrinsic processes.} {The radio (and optical) outburst predicted to peak around February--March 2004 on the basis of the previously observed 5--6 yr quasi-periodicity did not occur. The analysis of the optical light curves reveals now a longer characteristic time scale of variability of $\\sim 8$ yr, which is also present in the radio data. The spectral energy distributions corresponding to the XMM-Newton observations performed during the WEBT campaign are compared with those pertaining to previous pointings of X-ray satellites. Bright, soft X-ray spectra can be described in terms of an extra component, which appears also when the source is faint through a hard UV spectrum and a curvature of the X-ray spectrum. Finally, there might be a correlation between the X-ray and optical bright states with a long time delay of about 5 yr, which would require a geometrical interpretation.} ",
        "introduction": "Among the active galactic nuclei, blazars (i.e.\\ flat-spectrum radio quasars and BL Lac objects) exhibit the most extreme properties: violent variability at all wavelengths, high optical (and radio) polarization, superluminal motion of radio components, and brightness temperatures often exceeding the inverse Compton limit. These features are explained by assuming that the emission comes from a relativistic plasma jet closely aligned with the line of sight \\citep[see e.g.][]{bla78,bla79}.  One of the most intriguing blazars is AO 0235+164 (J0238+1636) at $z=0.94$ which, in addition to the above characteristics, also shows a complex environment, where a couple of foreground galaxies at $z=0.524$, possibly interacting, can both significantly contaminate its emission in the optical--UV band and absorb its radiation from the near-IR to the soft X-rays. Moreover, the stars of these objects might act as microlenses, producing at least part of the observed variability \\citep[e.g.][]{ost85,sti88}.  The analysis of the radio and optical light curves of AO 0235+164 from 1975 to 2000 revealed correlated variability and a quasi-periodicity of the main radio (and optical) outbursts on a $\\sim 5.7$ yr time scale \\citep{rai01}. The latter feature made this object a good candidate for the binary black hole scenario \\citep{rom03,ost04}, as in the case of OJ 287, another famous blazar \\citep[e.g.][]{sil88,vil98}.  In 2003--2005 AO 0235+164 has been the target of a huge multiwavelength observing effort chiefly aimed at verifying its possible quasi-periodicity, since a radio outburst had been predicted to peak around February--March 2004 (with an uncertainty of 6 months). The Whole Earth Blazar Telescope\\footnote{{\\tt http://www.to.astro.it/blazars/webt/}; see e.g.\\ \\citet{vil04b,boe05,rai05}.} (WEBT) consortium organized a radio, near-IR, and optical long-term monitoring campaign, during which three observations by the XMM-Newton satellite were performed, with simultaneous dense radio monitoring at the 100 m Effelsberg telescope. A spectroscopic monitoring was also carried out at the VLT, TNG, and NTT telescopes to study the relationship between the \\ion{Mg}{ii} line flux and the continuum flux density, and 15 VLBA epochs were granted to study the source structure variability. The results of these observations will be presented elsewhere.  The results of the first observing season 2003--2004 were reported by \\citet{rai05}, while a detailed analysis of the X-ray observations was performed by \\citet{rai06a}. Here we present the radio-to-optical data acquired by 27 telescopes during the second observing season, 2004--2005. The data of this extensive WEBT campaign 2003--2005, added to the historical light curves, allow one to study the behaviour of this source over the last 30 yr.  This paper is organized as follows: the optical, near-IR, and radio light curves obtained during the 2004--2005 observing season are presented in Sect.\\ 2. In Sect.\\ 3 we derive optical colour indices over a longer time interval, including the 1989--1991 and 1997--1998 outbursts, to study the spectral variability of the source. The intraday observations performed with the Effelsberg radio telescope during the August 2004 and January 2005 XMM-Newton pointings are presented and discussed in Sect.\\ 4. The long-term optical and radio behaviour of AO 0235+164 is presented in Sect.\\ 5. In Sect.\\ 6 we analyse 11 SEDs of the source corresponding to different epochs of X-ray observations, trying to understand how many emission components are involved. Finally, the conclusions are drawn in Sect.\\ 7. ",
        "conclusions": "Among blazars, AO 0235+164 presents several peculiar features. It has shown dramatic flux changes at all wavelengths on a wide range of timescales, including IDV. Its main outbursts have repeated quasi-periodically, at least for some time, with clear correlation of the radio and optical flux variations with very short time delays. When considering this along with the presence of an intervening galactic system, also gravitational microlensing becomes a viable interpretation for the source variability. At milliarcsecond scales it shows an extremely compact core, with faint extensions occasionally appearing at various position angles. Several observations suggest that the source radiation is unusually highly Doppler-boosted.  In this paper we presented the results of the radio, near-IR, and optical monitoring during the second observing season of the WEBT campaign 2003--2005 targeting AO 0235+164. The results of the first observing season and a detailed analysis of the three XMM-Newton observations performed during the campaign have already been reported in \\citet{rai05,rai06a}.  The big radio outburst, with a possible optical counterpart, which was foreseen to peak around February--March 2004 on the basis of the quasi-periodicity of $5.7 \\pm 0.5$ yr discussed by \\citet{rai01}, did not occur. The radio flux remained rather low during the whole WEBT campaign, with some mild increased activity at the higher radio frequencies only in the last observing season 2004--2005. A mild increased activity in the last season also characterized the optical and near-IR light curves. The comparison of the optical and radio data taken during the WEBT campaign with the long-term light curves of this object shows how this last, long period of suppressed activity especially at radio wavelengths is unusual for this source. The analysis of the historical light curves reveals now a longer characteristic time scale of variability of $\\sim 8$ yr, more evident in the optical, but also present in the radio data along with the $\\sim 5.7$ yr one.  The colour-index analysis reveals noticeable spectral changes in the optical band, which nonetheless do not show correlation with the source magnitude. The long-term trend appears essentially achromatic, with a possible progressive reddening in the last years.  All the three XMM-Newton observations performed during the WEBT campaign found the source in a faint state at all wavelengths. Neither the optical nor the  X-ray light curves show any significant variability. In contrast, the simultaneous intraday radio observations at the 100 m Effelsberg telescope reveal some flux changes at 10.5 GHz. If they are ascribed to intrinsic processes, a Doppler factor up to $\\sim 100$ can in principle be derived, confirming previous findings that this is an extremely beamed source. However, on the basis of these data we cannot rule out that interstellar scintillation was responsible for at least part of the variability. Actually, a reduced Doppler factor would agree with recent findings pointing to a reduced beaming during the current relatively quiescent period.  We compiled the SEDs corresponding to the X-ray observations since the ROSAT pointing in summer 1993. Radio and optical data simultaneous with the X-ray observations have been looked for in our database in order to shed light on the various emission components involved. The aim was mainly to find clues to the presence of the UV--soft-X-ray bump already discussed by \\citet{jun04} and \\citet{rai05,rai06a}. The existence of this extra component, besides the classical synchrotron and inverse-Compton ones, is mainly supported by the UV data, when available, and by the soft X-ray spectra detected by ROSAT in summer 1993 and by XMM-Newton in February 2002.  A discussion on the possible interpretations of this bump can be found in \\citet{rai06a}. Among the various pictures, these authors focused mainly on the presence of a thermal accretion disc and on an additional synchrotron component. In the latter case, which was favoured by the authors, we can envisage two possibilities. Either there are actually two synchrotron components, coming from distinct and separate jet regions without any simple relationship, or we are in the presence of a continuous inhomogeneous jet whose radiation is for some reason suppressed in the optical--near-UV region. Here we have again two possibilities: i) the break in the spectrum is due to a discontinuity in the opacity (as in the case of Mkn 501, \\citealt{rai03p}), ii) the suppressed emission is caused by a misalignment of the corresponding jet region whose radiation is less beamed towards the line of sight\\footnote {Although at other wavelengths (and hence spatial resolution), this scenario has indeed been suggested by \\citet{bac06} to interpret the region of suppressed flux density in the intensity profiles of the VLBA images of BL Lacertae.}. The observed strong variability in both these SED components might favour a scenario where both the emission from the near-IR to the soft X-rays and the spectral gap change in time as the strongly bent jet precesses or rotates. This view may be supported by the large variations of the VLBI jet components position angles observed by many authors \\citep[see e.g.][]{fre00,fre06,jor01,pin06}.  We noticed how among the ten satellite pointings in 1993--2005, only two of them (ROSAT in summer 1993 and XMM-Newton in February 2002) found the source in a bright and soft X-ray state, possibly of synchrotron origin. The All Sky Monitor X-ray light curve of AO 0235+164\\footnote{\\tt http://heasarc.nasa.gov/docs/xte/ASM/sources.html} from January 1996 to November 2005 shows that X-ray bright states are indeed very uncommon for this source: only in the first half of 2002 a significant increase of the count rate was observed, in correspondence to the first XMM-Newton pointing.  The two bright X-ray states are separated by about 8.5 yr, and a $\\sim 9$ yr interval elapsed between the ROSAT and the 1984 EXOSAT observation, when a high (but very uncertain) X-ray flux was measured. A similar time separation is found when considering the last strongest optical outbursts of 1990 and 1998. Thus, both of them seem to follow an X-ray brightening by about 5 yr. It will be interesting to see whether an optical flux increase will occur in the next observing seasons, i.e.\\ about 5 yr after the 2002 X-ray bright state. It would support both the possible quasi-periodicity of $\\sim 8$ yr of the optical (and radio) outbursts and a correlation between the optical and X-ray emission with a long time delay. A correlation of this type would be explained by a scenario where a curved inhomogeneous jet precesses or rotates, thus subsequently aligning its different emitting portions with the line of sight.  Several WEBT members are continuing to monitor the source in the radio, near-IR, and optical bands so that we will be able to check whether the source activity continues to increase and eventually leads to a forthcoming outburst."
    },
    "0608/nucl-th0608019_arXiv.txt": {
        "abstract": "The bulk viscosity of three-flavor color-superconducting quark matter originating from the nonleptonic process $u+s\\leftrightarrow u+d$ is computed. It is assumed that up and down quarks form Cooper pairs while the strange quark remains unpaired (2SC phase). A general derivation of the rate of strangeness production is presented, involving contributions from a multitude of different subprocesses, including subprocesses that involve different numbers of gapped quarks as well as creation and annihilation of particles in the condensate. The rate is then used to compute the bulk viscosity as a function of the temperature, for an external oscillation frequency typical of a compact star $r$-mode.  We find that, for temperatures far below the critical temperature $T_c$ for 2SC pairing, the bulk viscosity of color-superconducting quark matter is suppressed relative to that of unpaired quark matter, but for $T\\gtrsim 10^{-3}\\,T_c$ the color-superconducting quark matter has a higher bulk viscosity. This is potentially relevant for the suppression of $r$-mode instabilities early in the life of a compact star. ",
        "introduction": "\\label{intro}  Compact stars consist of extremely dense matter, part or all of which may be deconfined quark matter \\cite{perry}. Quark matter that is sufficiently dense and cold is expected to be in a color-superconducting state \\cite{bailin}, characterized by Cooper pairing between the quarks, mediated by their attractive strong interaction. This is analogous to the formation of electron Cooper pairs in a superconducting metal \\cite{bcs}. However, unlike electrons, quarks occur in many varieties: there are three colors and in a compact star core we expect to find three flavors, so a multitude of pairing (and thus symmetry breaking) patterns and hence different phases is theoretically possible \\cite{reviews}. It is an open problem of current research to determine the ground state out of these many possibilities for quark matter inside a compact star.  One approach to this problem is purely theoretical, using perturbative QCD as well as more phenomenological models in order to compare different possibilities and find the one with the lowest free energy. This approach yields a definite answer for very large densities (more precisely for quark chemical potentials $\\mu$ much larger than the masses of the up, down, and strange quarks). In this situation, quark matter is in the color-flavor locked (CFL) state \\cite{Alford:1998mk}. However, in nature, i.e., in the interior of a compact star, the quark chemical potential is expected to be of the order of 500 MeV or less, which requires the mass of the strange quark $m_s$ to be taken into account. Moreover, further constraints on the matter are imposed by electric and color charge neutrality as well as by equilibrium with respect to the weak interaction. All these conditions impose a stress on the pairing, no matter which pairing pattern is chosen \\cite{Rajagopal:2005dg}. Therefore, besides the CFL phase, many unconventional color-superconducting phases have been proposed, e.g., phases that break rotational \\cite{spin1} or translational \\cite{LOFF} symmetries.  The second approach to determining the ground state of dense quark matter in a compact star is the study of possible astrophysical effects of color superconductors. In particular, the ultimate goal is to compare observational data with theoretical predictions for color superconductors in order to exclude or confirm proposed phases. Moreover, one would like to use observations to distinguish between quark matter and more conventional matter inside the star, such as nuclear or hyperonic matter.  Promising physical properties that are distinct for different color-superconducting phases are related to transport properties such as thermal conductivity \\cite{Shovkovy:2002kv}, neutrino emissivity \\cite{Carter:2000xf,Jaikumar:2005hy,Schmitt:2005wg,Wang:2006tg}, magnetic fields \\cite{Alford:1999pb,Ferrer:2005vd}, and viscosity.  In this paper, we calculate the bulk viscosity of color-superconducting quark matter in the 2SC phase. (For a recent study of the bulk viscosity in spin-one color superconductors, see Ref.~\\cite{Sa'd:2006qv}; for hybrid/CFL matter see Ref.\\ \\cite{Drago:2003wg}. ).  In 2SC quark matter, up and down quarks form Cooper pairs with each other while the strange quarks remain unpaired. Although constraints of charge neutrality tend to disfavor the 2SC phase \\cite{Alford:2002kj,Steiner:2002gx}, it is found in phase diagrams calculated from NJL models, for example in Ref.~\\cite{Ruster:2005jc} (see their Figs.~2 and 6) and in Ref.~\\cite{Abuki:2005ms} (see their Fig.~7). In our calculations we shall neglect the effect of electric charge neutrality, so we treat the up and down quarks as having the same chemical potential and hence the same Fermi momentum. This has a quantitative effect, which we expect to be of order $m_s^2/\\mu^2$, on the available phase space for the quark momenta. We therefore expect our results to deviate only slightly from the results for electrically neutral 2SC matter.  Bulk viscosity, as well as shear viscosity \\cite{Manuel:2004iv}, is relevant to vibrational and rotational modes of a compact star, whose frequencies are typically in the kHz range \\cite{Kokkotas:2001ze}. In particular, for a rapidly rotating star there is the possibility of rotational ``$r$-modes'' that are unstable with respect to gravitational radiation \\cite{Andersson:1997xt,Andersson:2001bz,Lindblom:2000jw,Weber:2004kj}. This would lead to a drastic decrease of the rotation frequency of the star, which is incompatible with observations of rapidly rotating pulsars with periods of the order of ms. One mechanism to damp the instabilities is a nonzero viscosity of the fluid \\cite{cutler}.  In general, bulk viscosity characterizes the response of the system to an externally imposed oscillating volume change.  If the volume compression (and expansion) induces a change in the chemical composition of the matter, i.e., if it drives the system out of chemical equilibrium, then microscopic processes will try to re-equilibrate the system. If these processes occur on a timescale comparable to that of the driving oscillation then power dissipation will occur. Because the strange quark is massive, a volume oscillation causes a difference between the chemical potential of the $s$ quark and that of the $u$ and $d$ quarks.  The relevant microscopic re-equilibration process is the weak process $u+d\\leftrightarrow u+s$ \\cite{Wang:1985tg} which produces (annihilates) strangeness. (Leptonic processes are suppressed by their smaller phase space. However, we shall discuss the possibility that they might still yield a significant contribution to the bulk viscosity.) Strong processes can be neglected since they occur on much shorter time scales than the scale given by the oscillation of the star. This is in contrast to the context of heavy ion collisions where gluonic processes are important for the bulk viscosity \\cite{Arnold:2006fz}. Moreover, we do not consider a potential contribution from the Goldstone boson associated with the spontaneous breaking of $U(1)_A$ because this symmetry is explicitly broken at densities relevant for compact stars. Other (approximate) global symmetries are not broken in the 2SC phase, hence there are no other (pseudo) Goldstone bosons.  In order to compute the rate of the process $u+d\\leftrightarrow u+s$ one has to consider the corresponding collision integral. In the case of a superconductor/superfluid this integral is considerably more complicated than in the unpaired phase. Some of the complication originates from the broken number conservation, allowing not only processes that consist of two ingoing and outgoing particles but also processes with three (four) particles coalescing into one (zero) particles and one (zero) particles decaying into three (four) particles. In these processes, particles are created from or deposited into the condensate. This effect is also present in more conventional superfluids such as $^3$He \\cite{bhatta,vollhardt}. Starting from the kinetic equation, we present a detailed derivation of the collision integral which contains all these possible processes.  An additional complication arises from the nature of the 2SC phase. The process $u_1+d\\leftrightarrow u_2+s$ actually contains several subprocesses according to whether the incoming and outgoing states are gapped (up or down quarks that are red or green) or ungapped (all the other species).  All these subprocesses are included naturally in our formal derivation. We encounter processes where zero, one, two, or three of the quarks $u_1$, $d$, $u_2$, $s$ are gapped. Therefore, the 2SC phase is an interesting system to study the more general question of the quantitative effects of the energy gap for the rate of the microscopic process. We compute the rates for arbitrary temperatures $T$ (up to the superconducting transition temperature), i.e., we go beyond the low-temperature regime, where the effect of the energy gap $\\Delta$ is an exponential suppression of the rate.  A final subtlety is that the bulk viscosity is not a monotonic function of the rates of the equilibrating processes. The bulk viscosity is highest when those rates are closest to the time scale of the oscillation that is being damped. As we will see, this means that, depending what frequencies are physically relevant, a phase in which equilibration proceeds more slowly may have a higher bulk viscosity than one in which it proceeds quickly. Thus the 2SC phase, whose equilibration rate is always slower than that of unpaired quark matter, will turn out to have a larger bulk viscosity at kHz frequencies over a wide range of temperatures.  The paper is organized as follows. In Sec.~\\ref{definebulk} we define the bulk viscosity and point out an analogy to the power dissipation in an electric circuit with a capacitance and a resistance. The derivation of the rate of the process $u+d\\leftrightarrow u+s$ in 2SC quark matter is presented in Sec.~\\ref{deriverate}. At the end of this section, in Sec.~\\ref{subprocesses} we discuss the physical content of the derived expression. In Sec.~\\ref{evaluaterate}, the rate is evaluated in full generality and several limit cases (unpaired phase, zero-temperature limit) as well as several subprocesses and their contribution to the full result is discussed. We use these results to compute the bulk viscosity in Sec.~\\ref{calcbulk} and conclude with Sec.~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We have computed the bulk viscosity of three-flavor color-superconducting quark matter in the 2SC phase originating from the nonleptonic weak process $u+d\\leftrightarrow u+s$. To this end, we have presented a detailed calculation of the rate for this process for arbitrary temperatures. Since 2SC quark matter contains ungapped strange quarks as well as gapped and ungapped up and down quarks, this process contains several subprocesses. Each of these subprocesses that contains at least one gapped quasiparticle mode can be further divided into ``subsubprocesses'' which are similar to the ones in ordinary superfluids such as $^3$He. In both color superconductors and $^3$He, these subsubprocesses allow for creation and annihilation of quasiparticles from the condensate. Therefore, the process $u+d\\leftrightarrow u+s$ has contributions from decay and coalescence processes. We have presented a formal derivation of the total rate, starting from a kinetic equation and making use of the microscopic fermion propagators, resulting in an expression that naturally contains all these contributions. Our results are not only interesting for the special case of the 2SC phase but also, because of the multitude of subprocesses, provide a general quantitative analysis for the effect of a fermion condensate on the considered nonleptonic process.  We have evaluated the rate numerically, observing the expected exponential suppressions for small temperatures. For subprocesses with three gapped modes we find a suppression of the form $e^{-2\\Delta/T}$, while two or one gapped mode(s) yield $e^{-\\Delta/T}$. Upon evaluating the collision integral for temperatures up to the superconducting transition temperature, we have also studied the temperature regime where the exponential functions are no longer valid approximations for the rate.  The results for the rate, which translate into a typical time scale of the weak process, have been used to compute the bulk viscosity as a function of temperature for a given external oscillation frequency, chosen to be of order 1~kHz, which corresponds to the fastest rotation rate observed for pulsars. We have pointed out that the bulk viscosity is the response of the quark system to this external oscillation just as an electric circuit responds to an external alternating voltage by an induced current.  In this sense, we have identified a ``resistance'' and a ``capacitance'' of the quark system in terms of the microscopic process. Using this analogy, it is easy to understand that the behavior of the bulk viscosity assumes a maximum when the microscopic time scale matches that of the external frequency.   Our calculations show that, for a transition temperature of $T_c=30$ MeV and oscillation frequencies in the kHz range, the bulk viscosity of 2SC quark matter is larger than that of unpaired quark matter over a wide temperature range. This range is approximately given by $10^{-3}\\,T_c < T < 2T_c/3$ for an oscillation period of one millisecond. For longer oscillation periods, this range is even larger. As discussed above, this is because in that temperature range the slower equilibration rate of 2SC matter resonates better with kHz frequency oscillations than does the faster equilibration rate of unpaired quark matter. This result has potential physical implications for the behavior of the star early in its life, before its internal temperature drops below about 0.1 MeV. It is also interesting to note that at these early times neutrinos are trapped, and the resultant lepton number chemical potential can favor a 2SC phase with strange quarks \\cite{Steiner:2002gx} (but see also \\cite{Ruster:2005ib}). Moreover, bulk viscosity is likely to be the dominant factor for damping of potentially unstable $r$-modes in this very temperature regime \\cite{Kokkotas:2001ze,Weber:2004kj}. Although the star cools down below 0.1 MeV within days, this short time might potentially be long enough for instabilities to grow and spin down the star drastically \\cite{Lindblom:2000az}. For hot newly born stars, we have thus found that quark Cooper pairing actually works in favor of damping of unstable modes. For much smaller temperatures, $T\\lesssim 0.1$ keV, shear viscosity effects become important, and Cooper pairing is expected to reduce the shear viscosity and thus the damping of the unstable modes \\cite{Madsen:1999ci}.  We have found a similar effect for the bulk viscosity: for sufficiently small temperatures, $T<10^{-3}T_c$, the bulk viscosity of 2SC quark matter is smaller than that of unpaired quark matter.  Besides the damping of unstable rotational modes, bulk viscosity is also an important quantity for time scales of radial pulsation of compact stars \\cite{Sahu:2001iv}. Our results can serve as an ingredient for calculations addressing the viscous damping of these pulsations.  Finally, we emphasize that the results for the bulk viscosity in the 2SC phase cannot be extrapolated to the CFL phase. In this phase, all fermions acquire a gap in their energy spectrum and thus the rate of the process $u+d\\leftrightarrow u+s$ will be smaller than in the 2SC phase. In particular, this rate will be exponentially suppressed for small temperatures. However, kaons (which have nonzero strangeness) serve as a more effective source for re-equilibration, and their contribution to the bulk viscosity has to be taken into account \\cite{preparation}."
    },
    "0608/astro-ph0608141_arXiv.txt": {
        "abstract": "We present simulated observations of the Doppler shifts of the solar Mg~I Fraunhofer line scattered by asteroidal, cometary, and trans-Neptunian dust particles.  The studies are based on the results of integrations of orbital evolution of particles under the gravitational influence of planets, the Poynting-Robertson drag, radiation pressure, and solar wind drag. The derived shifts in the centroid and profile of the line with solar elongation are different for different sources of dust. A comparison of the velocities of zodiacal dust particles based on these numerical integrations with the velocities obtained from WHAM observations shows that the fraction of cometary dust particles among zodiacal dust particles is significant and can be dominant. A considerable fraction of trans-Neptunian dust particles among zodiacal dust particles also fits different observations. The mean eccentricity of zodiacal dust particles is estimated to be about 0.5. ",
        "introduction": "Ground-based spectroscopic observations allow one to study Doppler shifts of scattered solar Fraunhofer lines in the zodiacal light (e.g., James 1969, Hicks et al. 1974, Fried 1978, and East \\& Reay 1984).  Analysis of these shifts thus provide an opportunity to explore velocities of interplanetary dust in the inner solar system.  Using the Wisconsin H-Alpha Mapper (WHAM) spectrometer, Reynolds et al. (2004) were the first to obtain accurate measurements of the centroid velocities and line profiles of the scattered solar Mg I $\\lambda$5184 absorption line in the zodiacal light, both along the ecliptic equator and at high ecliptic longitudes.  The models for migration of dust particles are discussed in Section 2. These results are based on integrations of the orbital evolution of dust particles developed by Ipatov et al. (2004) and Ipatov \\& Mather (2006a-b) (all our recent papers can be found at astro-ph and at \\url{http://www.astro.umd.edu/$\\sim$ipatov}). In Sections 3 and 4, based on the results of our runs, we consider our model calculations of the radial velocity profile of the scattered Mg~I line.  We compare plots of centroid velocities and line widths of the Mg~I line versus solar elongation with WHAM observations obtained by Reynolds et al. (2004).  Studies of typical eccentricities and inclinations of zodiacal dust particles that fit the WHAM observations are presented in Section 5. In previous papers, observations have been compared with analytical models, but this is the first comparison with numerical integrations of the particles. ",
        "conclusions": "Our study of velocities and widths of the scattered Mg~I line in the zodiacal light is based on the distributions of positions and velocities of migrating dust particles originating from various solar system sources. These distributions were obtained from our integrations of the orbital evolution of particles started from asteroids, comets, and trans-Neptunian objects. A comparison of our models with the observations of velocities and widths of the zodiacal Mg~I line made by Reynolds et al. (2004) shows that asteroidal dust particles alone cannot explain these observations, and that particles produced by comets, including high-eccentricity comets (such as Comet 2P/Encke and long-period comets), are needed. A considerable fraction of trans-Neptunian dust particles among zodiacal dust particles also fits different observations. The mean eccentricity of zodiacal dust particles that best fit the WHAM observations is estimated to be about 0.5."
    },
    "0608/astro-ph0608188_arXiv.txt": {
        "abstract": "New visible polarization data combined with existing IR and FIR polarization data are used to study how the magnetic field threading the filamentary molecular cloud GF 9 connects to larger structures in its general environment. We find that when both visible and NIR polarization data are plotted as a function of extinction, there is no evidence for a plateau or a saturation effect in the polarization at $A_{\\rm V} \\approx$ 1.3 as seen in dark clouds in Taurus. This lack of saturation effect suggests that even in the denser parts of GF 9 we are still probing the magnetic field. The visible polarization is smooth and has a well-defined orientation. The IR data are also well defined but with a different direction, and the FIR data in the core region are well defined and with yet another direction, but are randomly distributed in the filament region. On the scale of a few times the mean radial dimension of the molecular cloud, it is as if the magnetic field were `blind' to the spatial distribution of the filaments while on smaller scales within the cloud, in the core region near the IRAS point source PSC 20503+6006, polarimetry shows a rotation of the magnetic field lines in these denser phases. Hence, in spite of the fact that the spatial resolution is not the same in the visible/NIR and in the FIR data, all the data put together indicate that the field direction changes with the spatial scale. Finally, the Chandrasekhar and Fermi method is used to evaluate the magnetic field strength, indicating that the core region is approximately magnetically critical. A global interpretation of the results is that in the core region an original poloidal field could have been twisted by a rotating elongated (core$+$envelope) structure. There is no evidence for turbulence and ambipolar diffusion does not seem to be effective at the present time. ",
        "introduction": "With the measurement of visible linear polarization of field stars, followed by developments of Near and Far Infrared (respectively NIR and FIR) and submillimetric polarimetry, it is now possible to trace magnetic field structures from the scale of the Milky Way to that of star forming regions and cores. In the diffuse Interstellar Medium (ISM), the general assumption that elongated dust grains usually have their smallest axis aligned in a direction comparable to that of the local magnetic field has lead to a toroidal representation of the Galactic magnetic field. A similar topology of the magnetic field can also sometimes be observed in other galaxies (e.g. \\citet{Scarrott1990}). On the smaller scales of molecular clouds, a similar relation exists between grain alignment and magnetic field directions (see \\citet{Lazarian2003} for an exciting review on the subject) and many topologies of the magnetic field are suspected to exist. Among others, helicoidal, toroidal, poloidal and hourglass configurations have been proposed to interpret observed polarization patterns, taking care of the physical conditions in these environments (e.g., Orion A,  M17). However, one of the most important limitations of linear polarimetry is that we only get a two-dimensional projected average representation of the magnetic field topology because of the integration on the plane of the sky of the contribution to the polarized emission due to each aligned grain, or its dichroic absorption for visible/NIR observations. This leads to a degeneracy problem since many three dimensional (3D) topologies of the field can produce identical polarization patterns. A method to get rid of this degeneracy problem is the use of models. On the scale of relatively nearby filamentary molecular clouds and cores, this method has been used and allows in some cases to select one field topology among many possible ones (e.g., \\citet{Matthews2001}, \\citet{Vallee2003}).  In addition to models, other diagnostics of magnetic fields can be used to lift the degeneracy due to the plane-of-sky integration. When column densities are high enough, observations of the Zeeman effect can yield the line-of-sight component of the magnetic field strength (e.g. \\citet{Crutcher1999}). Finally, by determining the ion-to-neutral molecular line-width ratio in non turbulent regions with a linear flow, it is possible to determine the angle between the field direction and the line of sight (\\citet{Houde2000a}, \\citet{Houde2000b}, \\citet{Houde2001}). By combining all three methods together, we get the 3D field topology (e.g., \\citet{Houde2002} for M17, \\citet{Houde2004} for Orion A). The observation of Zeeman splitting of spectral lines requires bright dense cores, but even in other regions, it is possible to get the 3D geometry of the field even if the field strength is unknown.  These methods can inform us about how magnetic fields can contribute to cloud dynamics, but our ultimate goal is to understand the competition between gravitational, thermal, magnetic, and turbulent forces. More specifically, magnetic fields and turbulence are both competing to slow down the star formation process, otherwise stars would form much faster than observations show. Combining polarization modeling techniques (e.g., \\citet{Fiege2000}) and predictions for magnetized simulations of turbulence (e.g., \\citet{Heitsch2001}, \\citet{Padoan2001}), one can hope to be able to determine the relative contributions of magnetic fields and turbulence.  To reach this goal, and to understand how magnetic fields thread the dark filaments, and to put observational constraints on modeling we propose to make comparisons of magnetic fields as seen in and around individual objects. This approach, namely a multi-scale analysis of the magnetic field, has begun to be done via comparisons of visible, NIR, FIR and/or submm polarization data in the OMC1 (\\citet{Schleuning1998}), OMC3 (\\citet{Matthews2001}) and GF 9 regions (\\citet{Jones2003}). In these works, aspects of the galactic magnetic field extending on tens of degrees are compared with aspects of the clouds probed in the densest regions of the clouds with resolutions $\\sim 10''$, but there is generally a lack of information between these two scales. Here we present observations at an intermediate scale, as much as the presence of field stars allows it, to have a clearer view of the topology of the field at the interface ISM/filamentary molecular clouds. From a general point of view, this approach should provide new insights on the way magnetic fields and clouds are dynamically coupled and shaped. In this paper we address this problem for the GF 9 filamentary molecular cloud region. A forthcoming paper will address the more complex situation for the Orion A region.  A description of data available about GF 9 is given in section \\ref{GF9REGION}. Observations and observational results are presented in section \\ref{OBSRES}. An analysis and a discussion follow in section \\ref{DISCUSS}, and our conclusions are given in section \\ref{CONCLU}. ",
        "conclusions": "\\label{CONCLU}  Visible linear polarization data taken at the Mont-M\\'egantic Observatory on 78 stars in the vicinity of the dark filamentary cloud GF 9 were presented and compared with available NIR and FIR polarization data covering part of the same region. Our main conclusions are:  On a scale of a few times the mean radial dimension of the molecular clouds, the magnetic field is smooth and well-ordered, and appears to be `blind' to the spatial distribution of the filaments.  On smaller scales in the core region, NIR polarimetry shows a rotation of the magnetic field lines when reaching into the denser parts.  There is no evidence for saturation of polarization as a function of extinction for both the visible and NIR data, suggesting that even in the denser regions the same grain properties and grain alignment mechanisms apply, and that the magnetic field is still being probed adequately there.  Finally, the magnetic field strength was estimated with the \\citet{Chandrasekhar1953} method. The results would suggest that the core region is approximately magnetically critical.  All these facts put together suggest that {\\it in GF 9 the magnetic field changes direction with the spatial scale.} The field differs by $35^{\\circ}$ in the GF 9 core compared to the surrounding region. Ambipolar diffusion does not appear to be supported given the current evidence. The effects expected from turbulence are definitely missing. A global interpretation of the results is that an original poloidal field could have been twisted into a toroidal configuration by a rotating elongated structure during its collapse.  Future work should check if the trend in polarization position angles to rotate counterclockwise with increasing wavelength still continues in the submm. Also, as the lack of saturation of polarization as a function of extinction suggests, if dust grains are magnetically aligned there should be no polarization hole in the submm in the core and the filament regions. As an additional benefit, submm polarization would afford a higher spatial resolution than the FIR data discussed here.  We thank the Conseil de recherche en sciences naturelles et en g\\'enie du Canada for supporting this research. The authors thank J. S. Greaves for providing information about observations of GF 9 with SCUBAPOL at the JCMT, as well as B. Malenfant and G. Turcotte for their helpful and friendly support during observations at Mont-M\\'egantic Observatory. We also thank an anonymous referee for his constructive comments and helpful suggestions. This work made extensive use of the SIMBAD database at the Canadian Astronomy Data Center, which is operated by the Dominion Astrophysical Observatory for the National Research Council's Herzberg Institute of Astrophysics.  \\clearpage"
    },
    "0608/astro-ph0608694_arXiv.txt": {
        "abstract": "{} {In this paper we  study  the effect of the quadrupole-term in the gravitational potential of a rotationally deformed central Be star on one armed density waves in the circumstellar disc. The aim  is to explain the observed long-term violet over red (V/R) intensity variations of the double peaked Balmer emission-lines, not only in cool Be star systems, but also in the hot systems like $\\gamma$ Cas.} {We have carried out  semi-analytic  and numerical  studies of  low-frequency one armed global oscillations in near Keplerian discs around Be stars. In  these we have  investigated surface density profiles for the circumstellar disc which have inner narrow low surface density or  gap regions, just interior to  global maxima close to the rapidly rotating star, as well as the mode inner boundary conditions.} {Our results indicate that  it is not necessary to invoke extra forces such as caused by line absorption from the stellar flux in order to explain the long-term V/R variations in the discs around massive Be stars. When there exists a narrow gap between the star and its circumstellar disc, with the result that the radial velocity perturbation is non-zero at the inner disc boundary, we find oscillation (and V/R) periods in the observed range for plausible  magnitudes for  the rotational quadrupole term.} {} ",
        "introduction": "Many Be stars exhibit long-term periodicities in the relative intensity of the double peaked Balmer emission lines, the so-called V/R variations, which occur on timescales of typically a few years to a few decades, see e.g. \\citet{Dachs87}, \\citet{Hub94} and \\citet{Ok97} and references therein. \\citet{Ok91} has suggested that a one armed (m=1) oscillation mode in the circumstellar disc about the Be star is responsible for the V/R variations. In a near Keplerian disc a one-armed mode shows up as a very slowly revolving perturbation pattern: the locally enhanced temperature and density causes stronger emission in one part of the disc. This results in either the violet- or red component of the double Balmer emission peaks to be brighter than normal when the compressed enhanced emission region is located in the approaching, respectively receding part of the disc. \\citet{Ok91} applied a point mass potential for the Be star and obtained a slowly precessing global m=1 pattern by the invoked pressure perturbations. However, in this way one generates {\\it retrograde} modes which tend to propagate throughout the disc, while observations showed the V/R variations to occur only near the central star. \\citet{Pap92} improved on this by considering a natural extension of this model by including the quadrupole-term in the gravitational potential of the rapidly rotating and thus flattened Be star. Deviation from a $1/r$ point mass potential is known to give rise to precessing elliptic particle orbits. In a gaseous disc the deviation from a point mass potential leads naturally to slowly revolving {\\em prograde} m=1 oscillation modes (if excited) which are confined to a region within a few stellar radii from the star. \\citet{Pap92} showed that the propagation region extends in fact from the inner disc boundary up to about the point where the modal pattern speed equals the local particle precession speed. The radial size of this natural oscillation cavity is of order the stellar radius $R_s$ as the quadrupole component decays rapidly with increasing distance from the star.  Meanwhile observations have shown that the perturbation pattern in several Be star-disc systems with V/R variations revolves indeed in a prograde sense, e.g. in $\\beta^1$ Mon \\citep{Telt94}. \\cite{HV95} and \\cite{HH97} calculated line profiles from discs with revolving m=1 modal patterns and showed that these eccentric modes can explain the observed variation of the asymmetric line profiles quite well. The existence of prograde m=1 modes close to the central star has also been directly detected by spectrally resolved interferometry of the Be stars $\\zeta$ Tau \\citep{Vak98} and $\\gamma$ Cas \\citep{Berio99}.  \\cite{Ok97} argued that, although the observed density patterns appear indeed to revolve in the prograde sense, the proposed mechanism by \\cite{Pap92} does not work for the hot early type Be star systems and he introduced a radiative force with the aim of confining the m=1 density waves to the inner part of these hot discs and induce prograde motions. Okazaki is indeed correct that in calculations with a rigid inner disc boundary \\citep{SH93} unrealistically low disc temperatures had to be adopted. In a more recent study of m=1 modes \\cite{FH06} note sensitivity of model predictions to disc model parameters, but comment that the predictive power of the model is better for the cooler Be stars where radiative forces are not added. In the latter work power law disc density profiles and rigid boundaries were also assumed.  However, in this paper we show, using both semi-analytic and numerical approaches, that if there is a narrow low  density region or gap between the star and its circumstellar disc, so that  the disc no longer has a rigid inner boundary but one that is essentially freely moveable, and there is also a significant single surface density maximum not far away, prograde m=1 modes with periods of a few years  can be obtained even for hot Be star systems. Although the mechanism by which a Be star ejects matter into the circumstellar disc is poorly understood, we note that for the mechanisms suggested in recent studies, such as non-radial oscillations in the rapidly rotating Be star \\citep{Riv01,Sav05}, or magnetic effects, one may indeed expect such a gap region to form when material with specific angular momentum larger than that corresponding to equatorial break up rotation is impulsively ejected from the surface during an outburst.  The plan of the paper is as follows: in Sect. 2 we give the basic equations for the oscillation modes and consider them in the low frequency limit. In that limit we use variational principles to show that rigid wall conditions are the least favourable for prograde modes and derive conditions on the stellar apsidal motion constant for their existence. These are found explicitly for a class of surface density profiles, with inner narrow  low surface density regions lying interior to nearby global surface density maxima that are also considered numerically. The semi-analytic and numerical approaches were found to be consistent.  In Sect.~3 we go on to apply the model to a hot   Be star: $\\gamma$ Cas (B0.5 IVe) and a cool Be star:  88 Her (B7 Ve). In both these cases prograde modes with reasonable periods are found for an appropriate degree of rotatonal distortion and apsidal motion constant, without the need to introduce an extra radiation force. Finally in Sect.~4, we summarize our conclusions. ",
        "conclusions": "We have investigated the effect of the rotational flattening of the rapidly spinning Be stars on low-frequency global m=1 oscillations in their circumstellar discs using both semi-analytic and numerical approaches.  We considered both a hot, early type Be star system like $\\gamma$ Cas and a cool late type system like 88 Her and found that if it is assumed there exists a (narrow) low surface density region or  gap between the star and the disc, such that the disc  inner boundary is free, low-frequency global m=1 oscillations can be supported with periods in the observed range of the long-term V/R variations of the Balmer emission lines for plausible values of the rotational factor $k_2 f^2(\\le 0.008)$. Therefore it is not necessary to introduce radiation forces  \\citep{Ok97} to obtain prograde m=1 modes that are contained in a oscillation cavity within a few stellar radii from a hot Be star surface.  Our calculations suggest that the more compact circumstellar discs exhibit m=1 modes with smaller oscillation periods than the more extended discs for which the density maximum occurs slightly further out in the disc and for which the density decreases more slowly with increasing distance from the star.  Of course in its present  linear form our simple dynamical model cannot explain the full complexity of the observed variations in Be stars, in particular  non-linear effects need to  be included \\citep{FH06}, especially in enigmatic systems like $\\gamma$ Cas \\citep{Mot05}, but it appears to lead to a  natural explanation for the long term V/R variations in Be star systems, given the observational fact that there are circumstellar discs with slowly revolving m=1 patterns. The model naturally accounts for  the observed \\citep{Berio99,Vak98} confinement and the prograde character of these patterns. How the circumstellar discs are formed and how the low frequency one armed density waves therein are  excited still remains a puzzle."
    },
    "0608/astro-ph0608377_arXiv.txt": {
        "abstract": "We study primordial nucleosynthesis abundance yields for assumed ranges of cosmological lepton numbers, sterile neutrino mass-squared differences and active-sterile vacuum mixing angles. We fix the baryon-to-photon ratio at the value derived from the cosmic microwave background (CMB) data and then calculate the deviation of the $^2$H, $^4$He, and $^7$Li abundance yields from those expected in the zero lepton number(s), no-new-neutrino-physics case. We conclude that high precision ($< 5\\%$ error) measurements of the primordial $^2$H abundance from, {\\it e.g.}, QSO absorption line observations coupled with high precision  ($< 1\\%$ error) baryon density measurements from the CMB could have the power to either: (1) reveal or rule out the existence of a light sterile neutrino if the sign of the cosmological lepton number is known; or (2) place strong constraints on lepton numbers, sterile neutrino mixing properties and resonance sweep physics. Similar conclusions would hold if the primordial $^4$He abundance could be determined to better than $10\\%$. ",
        "introduction": "\\label{sec-1} Recent developments in observational cosmology may allow the primordial abundances of the light elements, including deuterium, to become novel probes of the mass and mixing properties of light sterile neutrinos.  The high precision inference of the baryon-to-photon ratio $\\eta$ from the observed relative acoustic peak amplitudes in the cosmic microwave background (CMB) \\cite{WMAP, WMAP1,3yrwmap} suggests a new way to employ Big Bang Nucleosynthesis (BBN) considerations.  Historically, the comparison of the observationally-inferred light element abundances with calculated BBN abundance yields has been carried out with the intent of obtaining the baryon density. Indeed, on account of the near-exponential dependence of the deuterium yield on $\\eta$ in BBN, the observations of isotope-shifted hydrogen absorption lines in Lyman limit systems along lines of sight to high redshift QSO's provide another, independent high precision measure of the baryon content of the universe \\cite{Tytler,OMeara}. This value of $\\eta$ currently is in good agreement, within errors, with the CMB-derived value \\cite{burles,cyb03}.  However, these two independent determinations of the baryon density depend on new neutrino physics in different ways. In particular, the BBN deuterium yield depends, albeit weakly, on the neutron-to-proton ratio and the expansion rate at the BBN epoch \\cite {CF} and these quantities, in turn, can depend on the mass/mixing properties of sterile neutrinos.  It has been shown recently that the $^4$He abundance yield in Big Bang Nucleosynthesis (BBN) can be dramatically sensitive to medium-enhanced active-sterile neutrino flavor transformation in the presence of a significant lepton number \\cite{abfw,kfs}. This sensitivity comes about through (post-neutrino-decoupling) neutrino flavor mixing-induced alterations in the $\\nu_e$ and $\\bar\\nu_e$ energy spectra. These alterations cause changes in the weak interaction rates governing the inter-conversion of neutrons and protons, and so ultimately  they cause changes in the neutron-to-proton ratio in BBN. We show here that active-sterile neutrino mixing likewise can induce modest changes in the deuterium and $^7$Li abundance yields.  This sets up a potentially new avenue for probing or constraining the active-sterile neutrino mixing parameter space: comparison of the value of $\\eta$ derived from observationally-inferred deuterium on the one hand and the CMB-derived value on the other. Though the $^4$He BBN yield is far more sensitive to alterations of the neutron-to-proton ratio than is the $^2$H yield, at present the prospects for reliable and precise determination of the primordial deuterium abundance might be better than for helium.  The primordial helium abundance is likely between $23\\%$ and $26\\%$ by mass \\cite{OS,Olive,IT}. It may be possible to do much better than this by adroit attention to issues of radiative transfer and compact blue galaxy morphology \\cite{cyburt,Steigman,Steigman1}. However, as more QSO lines of sight become available, the statistics for deuterium abundance determinations in quasar absorption line systems will improve. Arguably, we may already know the primordial deuterium abundance at least as well as we know helium\\cite{Tytler,OMeara,pettini}. In any case, it is worth exploring how much leverage deuterium measurements have in constraining the parameter space of sterile neutrino mass/mixing values and lepton number(s).  The LSND anomaly is being re-investigated in the mini-BooNE experiment \\cite{LSND,miniB}. A positive signal in that experiment would indicate active neutrino coherent flavor transformation at a mass-squared scale significantly different from the atmospheric and solar neutrino mass-squared differences, $\\delta m^2_{\\rm atm} \\approx 3\\times{10}^{-3}\\,{\\rm eV}^2$ and $\\delta m^2_{\\odot} \\approx 8\\times{10}^{-5}\\,{\\rm eV}^2$, respectively. Given the $Z^0$-width limit, this would immediately imply the existence of a light sterile neutrino.  If this light sterile neutrino (and its helicity-flipped partner, or \\lq\\lq sterile antineutrino\\rq\\rq ) were to completely thermalize in the early universe, there could be both an increased $^4$He yield, which is possibly unwelcome, and trouble with CMB- and large scale structure-derived bounds on the sterile neutrino rest mass closure contribution \\cite{WMAP, WMAP1,abfw,Sloan,hannestad,murayama,aba, dib,dibad,barger,Kneller:2001cd,dod06}. Invocation of a significant net lepton number can suppress the production of sterile neutrinos in the epoch when neutrinos scatter frequently ({\\it i.e.,} prior to weak decoupling), thereby easing these constraints \\cite{FV,fv97,cirelli}. However, this lepton number will drive active-sterile resonant production of the sterile neutrino (or sterile antineutrino) after weak decoupling \\cite{abfw}. Post-weak-decoupling resonant sterile neutrino production would leave the active neutrinos and the sterile neutrino with distorted, non-thermal energy spectra which can have a significant impact on the neutron-to-proton ratio and the $^4$He yield \\cite{abfw}.  To investigate the effects of these spectral distortions and resonant transformation scenarios on the $^4$He, $^2$H and $^7$Li BBN yields, we follow the evolution of the neutrino distribution functions in various resonance sweep scenarios and self-consistently couple this with a calculation of the light element abundances performed with the full BBN nuclear reaction network code. (We employ a modified version of the Kawano/Wagoner code described in Ref.\\ \\cite{Smith:1992yy}.) In Section II we briefly outline the physics of active-sterile resonance sweep in the early universe. We describe our nucleosynthesis calculations in Section III. Results are given in Section IV. Discussion and conclusions are given in Section V. ",
        "conclusions": "We have for the first time self consistently and simultaneously coupled the full BBN nuclear reaction network with medium-enhanced active-sterile neutrino flavor transformation. One conclusion from our work is straightforward: the existence of light sterile neutrinos $\\nu_s$ which mix with active neutrinos $\\nu_e$, $\\nu_\\mu$, $\\nu_\\tau$ could alter significantly the relationship between primordial lepton numbers and the BBN light element ($^2$H, $^4$He, $^7$Li) abundance yields. Our work also shows that precision predictions of primordial nucleosynthesis yields for given neutrino properties and lepton numbers likely will require accurate treatment of the evolution of neutrino spectral distortion. Depending on neutrino mixing parameters, this may require a full $4\\times 4$ quantum kinetic equation neutrino flavor transformation scheme.  However, the calculations we have performed here allow us to point out some intriguing trends that eventually may provide a means for constraining lepton numbers and/or sterile neutrino properties. In particular, though we find that positive (negative) potential lepton number causes the BBN $^4$He, $^2$H, and $^7$Li abundance yields to be suppressed (increased) relative to the standard BBN zero-lepton-number case, these trends are counteracted and even {\\it reversed} when sterile neutrinos exist and medium-enhanced $\\nu_e\\rightleftharpoons \\nu_s$ ($\\bar\\nu_e \\rightleftharpoons \\bar\\nu_s$) takes place.  An underlying theme of our work is that the BBN paradigm, especially as it concerns constraints on neutrino physics, may be changing. We now know to fair precision the baryon density from the CMB acoustic peak amplitude ratios. The uncertainty in baryon-to-photon ratio $\\eta$ likely will improve with future observations. This trend will culminate in the near term in the Planck mission, which is projected to reduce the uncertainty in $\\eta$ to $\\leq 1\\%$. The near elimination of uncertainty in this quantity and the neutron lifetime will leave the leptonic sector as the principle source of uncertainty in conventional, thermal BBN. This is an intriguing development which comes hard on the heels of new experimental/observational revelations of neutrino mass-squared differences, flavor mixing parameters, and CMB- and large scale structure- derived bounds on light neutrino mass contributions to closure.  The active neutrino mixing parameters in particular allow $^4$He-derived constraints on electron lepton number $L_{\\nu_e}$ \\cite{Kneller:2001cd} to be extended to $L_{\\nu_\\mu}$ and $L_{\\nu_\\tau}$ \\cite{abb,wong,dolgov}. Certainly, improvements in precision in the observationally-inferred helium abundance $Y_p$ will translate directly into improved constraints on lepton number \\cite{Kneller:2001cd}. This is obvious in Figures \\ref{figure4}, \\ref{figure5}, \\ref{figure6}, \\ref{figure7}, for the value of $\\mathcal{L}_e = 4L_{\\nu_e}$ at which the \\lq\\lq Lepton Number Only Case\\rq\\rq\\ line crosses the limits of the horizontal band ($^4$He mass fraction $23\\%$). This occurs for $L_{\\nu_e} = \\mathcal{L}_e / 4 \\approx 0.045$, corresponding to $\\nu_e$ degeneracy parameter $\\eta _{\\nu_e} \\approx 0.07$. This is, of course, a crude upper limit, corresponding to our adopted range for the primordial helium abundance.  If the uncertainty in the observationally-inferred deuterium abundance D/H could be improved significantly, it could become competitive with $^4$He as a probe of lepton number and new neutrino physics. This is not the case currently. The statistical uncertainty in D/H is between $15\\%$ and $30\\%$, as derived from the isotope-shifted hydrogen absorption lines observed in Lyman Limit and Damped Lyman Alpha absorption systems along lines of sight to high redshift QSO's \\cite{Tytler,pettini}. As Figures \\ref{figure8}, \\ref{figure9}, \\ref{figure10}, and \\ref{figure11} show, we would need to get the error in D/H down to $\\leq 5\\%$ to enable deuterium to provide constraints on, or signatures of, {\\it e.g.}, sterile neutrinos.  This is problematic in the short term, but not an inconceivable eventuality in the longer term. The current deuterium abundance stems from of order a half dozen QSO absorption systems. With the projected increase in the number of 8m class telescopes and extensions of surveys to the southern hemisphere, we might expect to increase the number of ``clean\" ({\\it i.e.}, no interloper cloud) Lyman-$\\alpha$ absorption systems by a factor of 2 or so. This likely will not be good enough. However, the advent of 30m class telescopes could give many more QSO lines of sight and this might provide for much higher precision deuterium abundance determinations. This could be useful for probes of primordial baryon inhomogeneity or the star formation and chemical evolution history of the early universe. However, in this paper we point out that it also could be useful for constraints on new neutrino physics, especially as regards lepton numbers plus sterile neutrinos.  If, for example, we knew {\\it a priori} the sign of the primordial lepton number, then there may be a signature for sterile neutrinos or at least a means for better constraining their properties. A positive lepton number would be expected to give a suppression of $^2$H and $^4$He relative to the zero-lepton-number, standard BBN predictions based on the CMB-derived baryon density. However, a light sterile neutrino could reverse this and give an {\\it increase} in these abundances over the standard case. In broad brush, it works in an opposite sense for a negative lepton number, though the effect is less dramatic on account of the threshold in $\\bar\\nu_e+p\\rightarrow n+e^+$.  Likewise, if experimental evidence exists for active-sterile neutrino mixing, then the sign of the deviation of $^2$H, $^4$He, or $^7$Li from the standard zero-lepton number BBN case may afford a direct measure of the sign and magnitude of the potential lepton number. At the least, measurements and reasoning along these lines may allow for significantly better constraints on lepton numbers in this case.  There is another possibility.  Suppose that active-sterile mixing properties are measured in the lab and precise primordial light-element abundances are obtained from observation.  Further suppose that the effects (abundance deviations from standard BBN) pointed out here are {\\it not} seen.  This could indicate that there is another mechanism, other than large ($\\vert L_{\\nu_\\alpha} \\vert > 10^{-3}$) lepton number \\cite{abfw,FV,cirelli}, operating to suppress active neutrino scattering-induced de-coherence production \\cite{vol00,lee00,dol81,mck94,dib00,fv97} of seas of $\\nu_s$ and $\\bar\\nu_s$.  We know that some mechanism must suppress this process because otherwise there would be a conflict with CMB- and large scale structure-derived bounds on light neutrino mass contribution to closure (see for example Ref.\\ \\cite{abfw}).  Two alternative means for sterile neutrino production suppression have been suggested:  a low re-heat temperature for inflation \\cite{gel04}; and an alteration of neutrino mass/mixing properties ({\\it i.e.}, no mass or mixing) at early epochs.  In any case, the linkage between the light elements and new neutrino physics which we have pointed out here increases the leverage of new developments in both observational cosmology and terrestrial neutrino oscillation experiments."
    },
    "0608/astro-ph0608007_arXiv.txt": {
        "abstract": "Earlier papers introduced a method of accurately estimating the angular cosmic microwave background (CMB) temperature power spectrum based on Gibbs sampling.  Here we extend this framework to polarized data. All advantages of the Gibbs sampler still apply, and exact analysis of mega-pixel polarized data sets is thus feasible.  These advantages may be even more important for polarization measurements than for temperature measurements.  While approximate methods can alias power from the larger E-mode spectrum into the weaker B-mode spectrum, the Gibbs sampler (or equivalently, exact likelihood evaluations) allows for a statistically optimal separation of these modes in terms of power spectra.  To demonstrate the method, we analyze two simulated data sets: 1)~a hypothetical future CMBPol mission, with the focus on B-mode estimation; and 2)~a Planck-like mission, to highlight the computational feasibility of the method. ",
        "introduction": "Since the first detection of cosmic microwave background (CMB) polarization (DASI; Kovac et al.\\ 2002, Leitch et al.\\ 2002) and subsequent measurement of the temperature-gradient (TE) cross-power spectrum by the Wilkinson Microwave Anisotropy Probe (WMAP; Kogut et al.\\ 2003), emphasis has shifted to the measurement and analysis of the full polarization angular power spectra.  Many experiments \\citep{leitch:2005, readhead:2004, barkats:2005, montroy:2005, page:2006} have improved on those early findings, producing measurements with a considerable gain in raw sensitivity.  However, an important concern with all such measurements is systematic errors, including not only instrumental effects, observing strategy effects, and astrophysical contaminants, but also statistical issues.  It is essential to develop powerful and flexible data analysis tools to extract the desired information from the raw data reliably.  In this paper we progress towards this goal by extending the previously introduced Gibbs sampling framework \\citep{jewell:2004, wandelt:2004, eriksen:2004} to polarization.  The scientific importance of CMB polarization power spectra is high. For example, our current understanding of the optical depth, amplitude, and scalar spectral index hinges on what we know about the magnitude of the the low-$\\ell$ temperature and polarization spectra from the WMAP 3-year data \\citep{page:2006}.  Also, a detection of large scale B~modes would give a very exciting insight into primordial gravitational waves.  Earlier Gibbs analyses of unpolarized CMB data were described by \\cite{wandelt:2004, odwyer:2004, eriksen:2004, eriksen:2006}. These efforts demonstrated that exact analyses are indeed feasible even for such large data sets as the WMAP data, which comprise several million pixels. This is possible due to the very favorable scaling of the Gibbs sampling algorithm.  While brute-force likelihood evaluations scale as $\\mathcal{O}(N_{\\textrm{pix}}^3)$, $N_{\\textrm{pix}}$ being the number of pixels in the data set, the Gibbs sampler scales identically to the map making operation.  For the special case of uncorrelated noise and symmetric beams, this reduces further to $\\mathcal{O}(N_{\\textrm{pix}}^{3/2})$.  Thus, even Planck-sized data may be analyzed using these tools, as will be demonstrated in the present paper.  Gibbs sampling thus provides an efficient route to the exact posterior (or likelihood).  Moreover, it does not rely on any ad-hoc approximations.  Even for the analysis of temperature data, this proved to be both an important and subtle issue \\citep{spergel:2006, eriksen:2006}. However, it is critical for polarization measurements, because well-known approximate methods such as the pseudo-$C_{\\ell}$ methods (e.g., Chon et al.\\ 2004) can lead to aliasing of E-mode power into the much smaller B-mode power spectrum. Although it is possible to construct ways around this problem \\citep{smith:2005}, exact methods such as full likelihood evaluations or Gibbs sampling are clearly preferable solutions.  We start by discussing the algorithms used for polarized Gibbs sampling, extending the signal and power spectrum sampling steps from temperature to polarization.  Then we analyze simulated data to verify that the algorithm works and to determine the computational efficiency of the method.  \\clearpage ",
        "conclusions": "This paper extends the Gibbs sampling technique to polarized power spectrum estimation.  We have detailed the necessary generalization steps relative to the original temperature-only descriptions given by \\citet{jewell:2004}, \\citet{wandelt:2004} and \\citet{eriksen:2004}, and have considered computational aspects of polarized analysis.  The algorithm was demonstrated with two specific examples. First, considering a possible CMBPol type mission, we showed that the Gibbs sampler cleanly separates E- and B-modes, and no special care is required. This is in sharp contrast to approximate methods such as so-called the pseudo-$C_{\\ell}$, for which great care must be taken in order for the larger E-modes not to compromise the minute B-modes.  Second, we analyzed a Planck-sized data set, demonstrating that the algorithm is useful for analyzing the quantity of data which will come from near-future CMB experiments.  The Gibbs sampling results presented here use symmetric beams and noise which is uncorrelated between pixels. However, the Gibbs sampling algorithm has potential to analyze considerably more complicated data sets than these. For Planck, the solution lies in exploiting the very regular scanning strategy, which reduces the computational burden of a time-ordered data analysis. For a future CMBPol mission, the solution lies in the relatively large angular scales required. Since it is possible to invert the noise covariance matrix for multipoles up to several hundreds, one may pre-compute the all-important $\\mathbf{A}^T \\mathbf{N}^{-1}\\mathbf{A}$ matrix. After paying this high one-time cost, efficient and exact analysis is feasible using the methods described in this paper.  Finally, we re-emphasize that the Gibbs sampler provides a direct route to the exact likelihood (and to the Bayesian posterior), and it is much more reliable than approximate methods. This issue has been demonstrated explicitly through the analysis of the three-year WMAP data, where an approximate likelihood between $\\ell=13$ and 30 caused a non-negligible bias in the spectral index $n_{\\textrm{s}}$ \\citep{eriksen:2006}. Using Gibbs sampling, such worries are greatly reduced. Further, this paper demonstrates that the method is in fact capable of analyzing the amount of data that will come from the Planck mission with reasonable computational resources. It therefore seems very likely that this method will play a significant role in the analysis of future Planck data."
    },
    "0608/astro-ph0608231_arXiv.txt": {
        "abstract": "s{I discuss neutrino production in supernovae (SNe) and the detection of both Galactic core collapse events and the diffuse extra-galactic MeV neutrino background expected from the integrated history of star formation.  In particular, I consider what processes might affect our expectations for both. I focus on ``rapid'' rotation, defined as leading to millisecond initial neutron star spin periods.  Rotation affects the neutrino luminosity, the average neutrino energy, the duration of the Kelvin-Helmholtz cooling epoch, and the ratios of luminosities and average energies between neutrino species; it can strongly suppresses the $\\bar{\\nu}_e$ as well as $\\nu_\\mu$, $\\bar{\\nu}_\\mu$, $\\nu_\\tau$, and $\\bar{\\nu}_\\tau$ fluxes relative to $\\nu_e$.  As a result, depending on the prevalence of rapid rotation in SN progenitors through cosmic time, this may affect predictions for the MeV neutrino background and the history of nucleosynthetic enrichment. I emphasize connections between the MeV neutrino background and tracers of the star formation rate density at high redshift in other neutrino and photon wavebands.} ",
        "introduction": "When the iron core of a massive star collapses, the implosion is reversed at nuclear densities when nuclei dissociate into free nucleons.  The equation of state stiffens dramatically, driving a bounce shockwave into the supersonic infalling outer core. The bounce shock stalls almost immediately as a result of neutrino losses, the ram pressure of the infalling material, and  the dissociation of nuclei into free nucleons across the shock.  After the shock stalls, a characteristic post-bounce accretion structure obtains that is quasi-steady-state.   The hot ($\\sim$10\\,MeV) newly-born ``proto''-neutron star (PNS) has a neutrinosphere radius of $\\sim$80\\,km.\\footnote{Where the optical depth to $\\nu_e$ neutrinos $\\sim$2/3, an energy- and time-dependent quantity.}  Overlying the PNS is a subsonic accretion flow, bounded by the stalled stand-off shockwave at $\\sim$$200$\\,km. Just above the PNS is a cooling layer and beyond that a region of net neutrino heating (the ``gain'' region) provided by both accretion luminosity as the matter falling through the bounce shock is incorporated into the PNS and by core neutrino luminosity as the PNS cools and deleptonizes.  The revival of the shock to an energy of order $10^{51}$\\,ergs is the focus of supernova theory.  Any mechanism for this revival must fundamentally rely on the transfer of gravitational binding energy to the post-shock mantle.  The ``neutrino mechanism''\\cite{bethe}$^,$\\cite{janka2001} employs neutrino interactions  --- primarily $\\nu_e n\\leftrightarrow pe^-$ and $\\bar{\\nu}_e p\\leftrightarrow n e^+$ --- to transfer binding energy to the shocked matter in the post-bounce epoch. With the standard suite of microphysics and a solution to the Boltzmann equation for all neutrino species, the neutrino mechanism fails in one spatial dimension; the stalled bounce shock remains trapped forever, accreting the overlying stellar progenitor.\\cite{rampp2000}\\cdash\\cite{sumiyoshi2005} Although some models show that multi-dimensional effects might be necessary for success of the neutrino mechanism,\\cite{herant}\\cdash\\cite{fryer1999} recent calculations employing more sophisticated neutrino transport fail to explode\\cite{buras2003} --- albeit marginally. In some cases, successes are obtained.\\cite{buras2006}  However, the recent calculation of Ref.~\\refcite{burrows2006} suggests that energy deposition via neutrino interactions may be sub-dominant at late times with respect to acoustic heating generated by oscillations of the PNS generated by anisotropic accretion; the systematics of this new ``acoustic'' mechanism have yet to be elaborated.  Many recent complimentary works focus on the stability of the shockwave and the post-shock material, and its importance for the mechanism.\\cite{blondin}\\cdash\\cite{foglizzo}  Quite apart from the details of the explosion mechanism, the total energy budget dictating the character of neutron star birth is set by the gravitational binding energy of the neutron star: \\beq E_{\\rm bind}\\approx3\\times10^{53}M_{1.4}^2R_{10}^{-1}\\,\\,\\,{\\rm ergs}, \\label{binding} \\eeq where $M_{1.4}=M/1.4$\\,M$_\\odot$ and $R_{10}=R/10$\\,km are the neutron star mass and radius, respectively. Theoretical models suggest\\cite{burrows_lattimer}$^,$\\cite{pons1999} and detection of neutrinos from SN\\,1987A confirm that a fraction of order unity of this energy is radiated in $\\sim10$\\,MeV neutrinos on the Kelvin-Helmholtz timescale $\\tau_{\\rm KH}\\sim10-100$\\,s, long with respect to the collapse and explosion timescales.  Using these results, a number of studies have been made of the MeV neutrino background expected from SNe.\\cite{ando2002}\\cdash\\cite{yuksel2005} Indeed, the constraints on the background from SuperK\\cite{malek2003} have become tight enough that, for example, a total energy radiated in $\\bar{\\nu}_e$ neutrinos of $E_{\\rm bind}/3$ per SN with $\\langle\\varepsilon_{\\bar{\\nu}_e}\\rangle>20$\\,MeV is excluded.\\cite{yuksel2005} Estimates of the background must assume the total neutrino energy radiated per neutrino species per SN, the average neutrino energy per species, and neutrino oscillation parameters. For light water detectors like SuperK the dominant detection mode is single-species ($\\bar{\\nu}_e p\\rightarrow n e^+$) and because the cross section for this interaction is proportional to the neutrino energy squared the partitioning of energy, mixing, and the spectral shape are essential. The typical assumption is that the energy is partitioned equally between species, that the spectrum is thermal,  and that the average energy is $\\sim10-20$\\,MeV.  Although the total energy budget (in all channels) per SN is dictated by equation (\\ref{binding}), it is interesting to consider what processes might affect the best estimates of the diffuse neutrino background and the expected neutrino signal from the next Galactic SN at order unity. Potential processes are both microphysical and macrophysical/astrophysical. As an example of the former, it is possible that a change to the neutrino opacities or the equation of state for dense matter may alter the neutrino luminosity  ($L_\\nu$) and average energy ($\\langle\\varepsilon_\\nu\\rangle$) during the cooling epoch, while the total energy radiated per SN is unaltered.\\cite{pons1999}   Alternatively, as an example of astrophysical uncertainty, one may include the assumption of a universal IMF with constant (neutron-star-producing) SN rate per solar mass per year of star formation as a function of redshift, environment, and metallicity. In addition, there are a number of potential macrophysical effects that might modify our expectations at first order, including rapid rotation of the massive progenitor's iron core just before collapse.\\cite{heger2000}  Millisecond rotation of the PNS is interesting for several reasons. First, it changes $L_\\nu$ and $\\langle\\varepsilon_\\nu\\rangle$ and their ratios between species. This strongly effects the neutrino signature of core collapse and it may affect nucleosynthesis in the inner SN ejecta by altering the ratio of the $\\nu_e$ to $\\bar{\\nu}_e$ fluxes during explosion.\\cite{pruet2005}$^,$\\cite{frolich2005} Second, rapid rotation may be accompanied by more significant gravitational wave emission than non-rotating progenitors, thereby opening up a new channel of emission for a fraction of $E_{\\rm bind}$.  Third, rapid rotation may affect the morphology and nucleosynthesis of the remnant.\\cite{tcq}$^,$\\cite{bucciantini2006} Finally, rapid rotation at birth has been theoretically motivated by the existence of magnetars.  Ref.~\\refcite{duncan_thompson} argue that millisecond rotation may be neceesary for production of the very high magnetic field strengths ($\\sim$$10^{15}$G) inferred for the magnetar neutron star subclass. Although uncertain, simple estimates imply that $>10$\\% of all SNe produce magnetars.\\cite{woods_thompson}  In \\S\\ref{section:rotation}, I summarize results\\cite{tqb} from a set of 1D simulations of core collapse, including rotation in a parameterized way, and I briefly discuss the implications of these models for neutrino detection.  In \\S\\ref{section:budget} I discuss the neutrino background and its connection to other observable backgrounds. ",
        "conclusions": ""
    },
    "0608/astro-ph0608225_arXiv.txt": {
        "abstract": "We have conducted a radio pulsar survey of 56 unidentified $\\gamma$-ray sources from the 3rd $EGRET$ catalog which are at intermediate Galactic latitudes ($5^{\\circ} < |b| < 73^{\\circ}$). For each source, four interleaved 35-minute pointings were made with the 13-beam, 1400-MHz multibeam receiver on the Parkes 64-m radio telescope.  This covered the 95\\% error box of each source at a limiting sensitivity of $\\sim 0.2$~mJy to pulsed radio emission for periods $P\\ga 10$~ms and dispersion measures $\\la 50$ pc cm$^{-3}$. Roughly half of the unidentified $\\gamma$-ray sources at $|b| > 5^{\\circ}$ with no proposed active galactic nucleus counterpart were covered in this survey.  We detected nine isolated pulsars and four recycled binary pulsars, with three from each class being new.  Timing observations suggest that only one of the pulsars has a spin-down luminosity which is even marginally consistent with the inferred luminosity of its coincident $EGRET$ source.  Our results suggest that population models, which include the Gould belt as a component, overestimate the number of isolated pulsars among the mid-latitude Galactic $\\gamma$-ray sources and that it is unlikely that Gould belt pulsars make up the majority of these sources. However, the possibility of steep pulsar radio spectra and the confusion of terrestrial radio interference with long-period pulsars ($P \\ga 200$~ms) having very low dispersion measures ($\\la 10$ pc cm$^{-3}$, expected for sources at a distance of less than about 1~kpc) prevent us from strongly ruling out this hypothesis.  Our results also do not support the hypothesis that millisecond pulsars make up the majority of these sources. Non-pulsar source classes should therefore be further investigated as possible counterparts to the unidentified $EGRET$ sources at intermediate Galactic latitudes. ",
        "introduction": "Determining the nature of Galactic $\\gamma$-ray sources with energies above 100 MeV is one of the outstanding problems in high-energy astrophysics. The $EGRET$ telescope on the Compton Gamma-Ray Observatory, which was active from 1991 to 1999, identified about half a dozen of the brightest $\\gamma$-ray sources in the Galactic plane as young pulsars \\citep{tbb+99}.  It also demonstrated that most of the sources at low Galactic latitudes ($|b| \\la 5^{\\circ}$) are associated with star forming regions, and hence may be pulsars, pulsar wind nebulae, supernova remnants, winds from massive stars, or high-mass X-ray binaries \\citep{kc96,yr97,rbt99}. In addition, molecular clouds can either be sources of $\\gamma$-rays or enhance the production of $\\gamma$-rays by particles produced by the source classes mentioned above \\citep{aha01}.  Various targeted multi-wavelength campaigns to identify low-latitude sources have discovered a number of likely counterparts \\citep{rrk01,hcg+01, rhr+02,brrk02,hgc+04}. The recent Parkes Multibeam Survey has also discovered several new pulsars coincident with $EGRET$ $\\gamma$-ray sources; these pulsars have spin characteristics that are similar to those of the known $\\gamma$-ray pulsars \\citep{dkm+01,kbm+03}.  While there are many candidate counterparts to $EGRET$ sources at low latitudes, there are few firm identifications owing to the large positional uncertainties of the sources (typically $\\sim 1^{\\circ}$ across). In general, a timing signature, such as a pulse detection, is necessary to be certain of a source identity.  Since young pulsars tend to be noisy rotators, extrapolating a pulse ephemeris reliably back to the era of the $EGRET$ observation is generally not possible. With the improved resolution and sensitivity of the upcoming $AGILE$ and $GLAST$ missions, the low-latitude $EGRET$ sources should be more easily identified.  There are estimated to be between 50 and 100 sources detected by $EGRET$ at mid-Galactic latitudes which are associated with our Galaxy.  As a class, these sources tend to be fainter and have steeper spectra than those at low latitudes \\citep{hbb+99}.  Their positional uncertainty is therefore on average even greater ($\\sim 1.5^{\\circ}$ across) than it is for the low-latitude sources. These mid-latitude sources have a spatial distribution which is similar to the Gould belt of local regions of recent star formation plus a Galactic Halo component \\citep{gre00,gre01}.  The Gould belt provides a natural birth place for many nearby ($\\la 0.5$~kpc), middle-aged pulsars similar to Geminga \\citep{hh92}. Both the outer gap \\citep{yr95} and polar cap \\citep{hz01} models of pulsar emission suggest that many of these pulsars should be detectable in $\\gamma$-rays but that the majority should have their radio beams missing Earth. However, if predictions from recent models are realistic, then between 25\\% and 50\\% of $\\gamma$-ray pulsars might still be visible to us as radio pulsars \\citep{gvh04,czlj04}.  The mid-latitude $EGRET$ source distribution is also similar to the distribution of recycled pulsars in the Galactic field \\citep{rom01}. The fastest millisecond pulsars (MSPs) can have spin-down luminosities ($\\dot{E} \\propto \\dot{P}/P^{3}$) and magnetospheric potentials similar to those of young pulsars. There has been one possible detection of $\\gamma$-ray pulsations from an MSP \\citep{khv+00} and some preliminary modeling of that emission \\citep{hum05}. If a significant fraction of the mid-latitude sources are MSPs at typical Galactic distances, many should be detectable as radio pulsars \\citep{sgh05}. Since MSPs tend to be in binary systems, $GLAST$ will not be sensitive to them in blind searches (owing to computational reasons associated with the very long integration times and the large number of trials required to search the parameter space).  Here we describe a radio pulsar survey of 56 unidentified sources from the 3rd $EGRET$ catalog (3EG) \\citep{hbb+99} which are at intermediate Galactic latitudes ($5^{\\circ} < |b| < 73^{\\circ}$).  The survey used the 1400-MHz, 13-beam multibeam receiver \\citep{swb+96} on the 64-m radio telescope in Parkes, Australia to search for pulsed emission. This receiver has been used very successfully to find pulsars in a number of recent radio pulsar surveys \\citep{mlc+01, ebs+01, kbm+03, mfl+06, bjd+06}.  Discovery of radio pulsar counterparts to these $EGRET$ sources would not only provide interesting systems for individual study and establish the identifications of the target sources (e.g., Roberts et al. 2002\\nocite{rhr+02}), but it would also help resolve outstanding questions about the pulsar emission mechanism and the physical origin of pulsar radiation at different wavelengths (see, e.g., Harding et al. 2004 and references therein\\nocite{hgg+04}). ",
        "conclusions": "There are now 20 pulsars that are known to lie within 1.5 times the radius of the 95\\% confidence contours of $EGRET$ sources at $|b| > 5^{\\circ}$. Of these, only the Crab pulsar and PSR B1055$-$52 have confirmed associations with the coincident $\\gamma$-ray emission. Of the remaining 18 pulsars, including the 13 detected in our survey and the recently discovered PSR J2243+1518 \\citep{cml05}, none is energetic enough for a clear association. Other than PSR J1614$-$2230, which is at best a marginal candidate, no pulsars from any survey have been found which can be associated with unidentified $EGRET$ error boxes at mid-Galactic latitudes.  Non-pulsar source classes should therefore be investigated further. \\citet{gkr05} discuss the viability of low-mass microquasars as $EGRET$ sources. Recently, there has been the suggestion that much of the $\\gamma$-ray emission at mid-latitudes is due to gas not being included in the models used for calculating the $\\gamma$-ray background maps \\citep{gct05}. In this case, many of the cataloged sources may not be truly point-like. Regardless, as suggested by spectral and variability studies of the population (e.g., Grenier 2003\\nocite{gre03}), the likelihood of pulsars being able to account for a majority of the cataloged unidentified $EGRET$ sources at intermediate Galactic latitudes seems remote."
    },
    "0608/astro-ph0608013_arXiv.txt": {
        "abstract": "Optical and X-ray studies of six nearby galaxies show that the probability a globular cluster will be an X-ray source is consistent with being linearly proportional to its mass.  { We show that this result is consistent with some recent estimates of the velocity kick distributions for isolated radio pulsars -- those which are the sum of two Maxwellians, with the slower distribution at about 100 km/sec -- so long as a large fraction of the retained binaries are in binary systems with other massive stars.}  We confirm that over a large sample of galaxies, metallicity is clearly a factor in determining whether a globular cluster will contain an X-ray binary, and we estimate the transformations between color and metallicity for a large number of optical filter combinations.  We also show that the core interaction rate is roughly linearly proportional to the stellar mass of a globular cluster for the Milky Way when one bins the clusters by mass. ",
        "introduction": "It has long been known that globular clusters contain much larger numbers of X-ray binaries per unit stellar mass than do field populations of galaxies (Katz 1975; Clark 1975).  This overabundance has been ascribed to stellar interactions in globular clusters which allow neutron stars and/or black holes to enter new binary systems through either tidal capture (Clark 1975; Fabian, Pringle \\& Rees 1975) or exchange interactions (Hills 1976) long after the supernovae that produce them.  In the Chandra era, it has been possible to extend this work to show that the clusters with the highest interaction rates are most likely to contain X-ray sources, and probably are most likely to contain accreting neutron stars (Pooley et al. 2003; Heinke et al. 2003; Gendre 2005).  While studies of Galactic globular clusters have thus been fruitful, the Milky Way's globular cluster system is rather small (containing only about 150 clusters), so it is not possible to study certain phenomena due to lack of statistical significance in a small sample of clusters.  Furthermore, some of the Milky Way's globular clusters' parameters, such as metallicity and galactocentric radius, are strongly correlated with one another, making it difficult to isolate the causes of certain effects.  As a result, extragalactic globular cluster systems can be invaluable in complementing the Galactic globular clusters for producing the best possible data on how globular cluster parameters affect X-ray binary production.  While ROSAT was able to observe globular cluster X-ray sources in M~31 (see e.g. Supper et al. 1997), it was not until the advent of the Chandra X-ray Observatory that there was sufficient angular resolution to resolve apart the sources in more distant galaxies, or to obtain sufficiently accurate positions that comparisons could be made with optical counterparts.  Early on, by comparing the positions of X-ray sources with those of optically detected globular clusters, it was found that large fractions of the X-ray binaries in elliptical galaxies were in globular clusters (Sarazin, Irwin \\& Bregman 2001; Angelini, Loewenstein \\& Mushotzky 2001 -- ALM).  Larger samples of galaxies have shown a trend where the fraction of X-ray binaries in globular clusters in a galaxy seems to vary as a function of galaxy type, increasing from spiral galaxies to lenticular to normal ellipticals to cD galaxies (Maccarone, Kundu \\& Zepf 2003; Juett 2005; Irwin 2005).  From the observations of elliptical galaxies, with their large globular cluster systems, it has been possible to find correlations between cluster properties and the probability that a cluster would contain an X-ray binary which are not statistically significant in the Milky Way, or which have ambiguous interpretations in the Milky Way because of the aforementioned correlations between different cluster parameters.  Specifically, it has been shown clearly that metal rich clusters are far more likely to contain X-ray binaries than metal poor clusters (Kundu, Maccarone \\& Zepf 2002), confirming suggestive results from the Milky Way (Grindlay 1993; Bellazzini et al. 1995). These results have been found again in numerous subsequent papers (see e.g. Jordan et al. 2004; Minniti et al. 2004; Xu et al. 2005; Kim et al. 2005).  It has also been shown that more luminous clusters are more likely to contain X-ray binaries than less luminous clusters (see e.g. ALM;KMZ).  We refer the reader also to reviews by Verbunt \\& Lewin (2006) and Fabbiano (2006) for a broader overview of the literature.  Other results have been suggested by the data, but are not as clearly significant.  While in the Milky Way, it is clear that denser globular clusters are more likely to contain large numbers of X-ray sources per unit stellar mass (Pooley et al 2003; Heinke et al 2003; Gendre 2005), the extragalactic clusters have core radii which have not been sufficiently well resolved spatially as to establish this effect clearly.  KMZ found an anti-correlation between cluster half-light radius and LMXB hosting probability in NGC~4472, but noted that it was statistically insignificant.  Jordan et al. (2004) estimated King model parameters of globular clusters in M87.  They then inferred stellar collision rates ($\\Gamma$) from the model fits, and found that the collision rates were strongly correlated with the likelihood that a cluster would contain an X-ray binary, but that this correlation was only marginally better than the correlation with cluster luminosity. The absence of strong signatures of cluster concentration should not be taken as evidence that cluster concentration is uncorrelated with X-ray binary production.  Half-light radius is not expected to be very well correlated with core concentration, making the null result of KMZ unsurprising.  Because the core radii of globular clusters in the Virgo Cluster are typically a small fraction of a pixel on the Advanced Camera for Surveys aboard the Hubble Space Telescope, neither is it surprising that Jordan et al. (2004) were unable to find a strong signature of cluster concentration.  In this paper, we analyze a sample of six galaxies for which there is good optical and X-ray data.  We show unambigously that the metallicity effect is strong.  We then isolate the effect of cluster mass, $M$, on probability that a cluster will contain an X-ray source, and find that the probability a globular cluster will contain an X-ray source scales as $M^{1.03\\pm0.12}$.  { We compare this value with the expected value from various kick velocity distributions, and find that this result is most consistent with double Gaussian distributions for pulsar kick velocities, where there exists a slow mode to the neutron star kick velocity distribution}. ",
        "conclusions": "\\subsection{Implications of the observational results for retention fractions} The number of X-ray binaries in a globular cluster is likely to be correlated most strongly with the stellar collision rate.  The collision rate should be dominated by the collisions which take place in the core, and it is straightforward to compute analytically the collision rate in the core of a globular cluster: \\begin{equation} \\Gamma=\\rho_c^{3/2}r_c^{2}, \\end{equation} where $\\rho_c$ is the central density of the cluster and $r_c$ is the core radius.  { This formula assumes virial equilibrium in the cluster core (Verbunt \\& Hut 1987; Pooley et al. 2003), and follows from the more general $\\Gamma=\\rho_c^{2}r_c^{3}/\\sigma$, where $\\sigma$ is the velocity dispersion, since $\\sigma\\propto$$\\rho_c^{0.5} r_c$.}  The only galaxy where the core radii and central densities of a large number of clusters have been reliably measured (rather than inferred from model fits) is the Milky Way.  The collision rates for the Milky Way are well correlated with the masses; when we bin the Milky Way's globular clusters by mass and fit a power law to the data.  We find that the best fitting power law index varies between 0.9 and 1.3, depending on the number of bins used (see figure \\ref{bin_MW} for a characteristic fit).  { The statistical errors on the fits are always smaller than 0.1 in index.}  We thus adopt 1.1$\\pm$0.2 for the exponent here, but note that this is a systematic error, and really represents a hard bound on the possible values, rather than a $1\\sigma$ error. {\\it This result is intermediate between the two cases where the spatial properties are independent of mass (i.e. where, among the King model parameters, only the central density varies, and the expectation value would be $\\Gamma\\propto$$M^{1.5}$)} and where the concentration index (i.e. the ratio of tidal radius to core radius) is independent of mass (in which case the expectation value would be $\\Gamma\\propto$$M^{2/3})$.  The intermediate result is expected given that cluster concentration increases with increasing cluster mass (Djorgovski \\& Meylan 1994).  { It is possible to find other means of producing various exponents besides those presented above, if one allows for different relations between the fraction of the mass in the cluster core and the core radius with mass.  The key point, though, is that empirically, the relation for Milky Way globular clusters is that $\\Gamma$ scales approximately as $M^{1.1}$.}  \\begin{figure} \\begin{center} \\renewcommand{\\epsfsize}[2]{0.7#1} \\epsfig{file=5298fig3.ps,width=8cm} \\caption{Binned collision rate versus mass for the Milky Way's globular clusters, using data from the Harris catalog with 8 bins.} \\label{bin_MW} \\end{center} \\end{figure}   Jordan et al. (2004) found that the probability a cluster in M~87 will contain an X-ray binary is proportional to $\\Gamma \\rho_c^{-0.4\\pm0.1}$, and suggested that the deviation from a purely linear relationship with $\\Gamma$ might be due to destruction of binaries by dynamical interactions in the globular cluster.  While clusters do undoubtedly destroy binary systems, this should be done preferentially for the longest period systems, perhaps by making them highly eccentric rather than by destroying them outright if they are Roche lobe overflowing systems(e.g. Hut \\& Paczynski 1984; Rasio \\& Heggie 1995; Maccarone 2005).  Binary destruction should have relatively little impact on the short period systems which, due to their higher duty cycles (see e.g. Piro \\& Bildsten 2002; Bildsten \\& Deloye 2004; Ivanova \\& Kalogera 2006) should provide most of the {\\it bright} X-ray binaries which can be seen out to Virgo Cluster distances (see Maccarone et al. 2005 for a discussion of the effects of destruction of X-ray binaries due to dynamical interactions on their orbital period distributions).  In light of our above finding that, { empirically,} for the Milky Way, $\\Gamma$ scales with cluster mass to the power 1.1$\\pm$0.2, it seems simpler just to consider that the estimates of the core radii of the globular clusters in M~87 are likely to be dominated by measurement uncertainties. In the absence of useful radial information, the derived central density of a globular cluster will scale linearly with the cluster's mass; the estimate of $\\Gamma$ will then scale with the cluster mass to the power 3/2.  We have shown that, for the Galactic globular cluster systems, $\\Gamma$ scales linearly with the cluster mass; therefore, one will need to correct this derived value of $\\Gamma$ by a term of order $M^{1/2}$ (or $\\rho^{1/2}$ in the absence, again, or any useful radial information).  The $\\rho^{-0.4}$ term found by Jordan et al. (2004) is thus more likely this ``correction'' term than a term indicating dynamical destruction of accreting binaries. { The only other means of producing the $\\Gamma$-M correlation we show and the relation among $\\Gamma,M,$ and $\\rho_c$ found by Jordan et al. (2004) would be a specific relationship between the fraction of mass in the core and the core radius for clusters, in which the cluster core density is uncorrelated with the cluster mass. This would, in principle, be possible, e.g. if the most massive clusters has progressively lower fraction of their mass in their core but is, in fact, not the case, given the strong observed correlation between cluster central density and cluster absolute magnitude (Djorgovski \\& Meylan 1994).}  Finally, let us consider an additional important point.  The collision rate $\\Gamma$ is a collision rate of typical stars with one another -- mostly main sequence stars with main sequence stars.  { The collision rate for neutron stars will be the collision rate itself, times the fraction of stars in the cluster core which are neutron stars.  This fraction will come from { three} factors -- stellar evolution, which determines what fraction of stellar mass ends up locked up in neutron stars, the retention fraction, which characterizes what fraction of those neutron stars remain in the cluster, { and the effects of mass segregation, which determines how over-represented neutron stars are in the cores of globular clusters, compared to their representation in the clusters on the whole}.  We assume that the stellar evolution effects have no mass dependence, so any difference between the collision rate as a function of cluster mass and the LMXB hosting probability as a function of cluster mass is most naturally explained by the dependence of the retention fraction on cluster mass, and mass segregation variations as a function of mass.  { As a result, we define a quantity, $\\Phi$, which should scale linearly with the expected formation rate of neutron star X-ray binaries in the core of a globular cluster.  We add only a single assumption at this point, which is that mass segregation works sufficiently effectively on timescales of the ages of globular clusters that all the neutron stars are in the cluster cores.  This introduces an additional factor of $M/M_{core}$ to account for the ``overdensity'' of neutron stars among the total stars in a cluster's core.  }Since $M_{core}\\propto$$r_c^3\\rho_c$, this yields: \\begin{equation} \\Phi=\\Gamma M/M_{core} = \\rho^{1/2}M/r_c. \\end{equation} { Following the same procedure as above to estimate the $\\Gamma-M$ relation, we take the binned sum of the Milky Way's globular clusters' values for these same $\\Phi$.  We find that $\\Phi \\propto M^{0.5-0.7}$, again with the range of estimated exponent values dominated by the binning scheme used (i.e. whether 6, 8 or 10 bins are used) rather than by measurement uncertainties (see e.g. Figure {\\ref{phiplot}).  Since $\\Phi f_{ret}$ should give the LMXB hosting probability, that leads to the inference that $f_{ret} \\propto M^{0.4-0.6}$, since the hosting probability scales as $M^{1.1}$.  We also note that mass segregation of neutron stars is not sufficiently fast as to place {\\it all} neutron stars in the cores of globular clusters; there are some millisecond pulsars with well-measured positions that place them outside the cores of their host clusters (see e.g. Camilo \\& Rasio 2005 for a reasonably up-to-date list of globular cluster pulsars, or Paulo Freire's continuously updated list at http://www.naic.edu/~pfreire/GCpsr.html). While, relaxing this assumption leads to a weaker dependence of the retention fraction on mass, the assumption is generally a good one. About half of all misllisecond pulsars are within one core radius of the center of their host clusters, while the fraction of the total mass in the core is less than 10\\% even in 47~Tuc, which has the highest core fraction of any massive cluster.}  \\begin{figure} \\begin{center} \\renewcommand{\\epsfsize}[2]{0.7#1} \\epsfig{file=5298fig4.ps,width=8cm} \\caption{Binned values of $\\Phi$ versus mass for the Milky Way's globular clusters, using data from the Harris catalog with 6 bins.} \\label{phiplot} \\end{center} \\end{figure}  {\\subsection{Scaling relations of the retention fraction}} We have considered a variety of possible kick velocity distributions in order to determine whether any of them { are consistent with the estimated retention fraction relation -- $F_{ret}\\propto M^{0.5}$}. Essentially, two broad classes of velocity kick distributions have been proposed for radio pulsars: ``zero-peaked'' distributions (e.g. Paczynski 1990; Hartman 1997; Hansen \\& Phinney 1997), where the most likely initial velocity for the neutron star is close to zero (although the distribution retains significant probability density out to hundreds of km/sec), and distributions which are Gaussians peaked at the mean pulsar velocity (e.g. Lyne \\& Lorimer 1994) or Maxwellians, which have rather similar characteristics.  Some recent work has modelled the pulsars' velocity distributions as double Gaussians { or double Maxwellians} (Arzoumanian, Chernoff \\& Cordes 2002; Brisken et al. 2003), typically with one peak at about 100 km/sec and containing about 20\\% of the neutron stars and the other peak at about 300-500 km/sec containing about 80\\% of the neutron stars.  In essence, the two approaches have converged to rather similar distributions, since the low velocity { mode} ``fills in'' the gap left at low velocities in the single { mode} approach.  On the other hand, the most recent estimation of the kick velocity distribution argued for a return to a single-peaked Maxwellian distribution with a characteristic { 1D velocity of about 265 km/sec} (Hobbs et al. 2005), so there still exists controversy about the initial velocity distribution of radio pulsars.}   { \\subsubsection{Results of models tried} We have tried a variety of different models with different parameter values to determine whether they can produce retention fractions with negligible mass dependence. We present the methodology and a sampling of the results here.  In the interests of brevity, we do not present every calculation we have done, but note that these are in the Master's degree thesis of the lead author of this paper { albeit with a somewhat different interpretation section} and can be obtained by contacting the authors of this paper.  To calculate reliable retention fractions we use 111 globular clusters from the Harris catalogue as templates. We used the core and tidal radii from the catalogue, converting arcminutes to parsecs using the distances provided in the catalogue. We also converted the V magnitude to mass assuming a constant mass-to-light ratio of 1.8, as found by Piatek et al. (1994). The 111 clusters are all clusters from the catalogue that are not core-collapsed and include the 3 parameters we need for fitting.  We use two different models to describe the early globular clusters. The first is the King model fitted to present-day average GC properties (from the Harris catalogue). All heavy stars are assumed to be in the center of the cluster at the end of their lifetimes and the retention fractions are calculated using the central escape velocity. The second model assumes that no mass segregation has taken place. Stars of all masses are spread equally through the cluster according to the cluster's mass density profile. The Plummer model is used because it puts the GCs with different central potentials on equal footing. GCs evolve to small core radii and large tidal radii as they age. This means that the present-day central potential has no meaning when applied to primordial GCs. Through dynamical friction, the high mass stars may fall to the center before their maximum age is reached or they may end somewhere between where they started and the center of the cluster.  The escape velocity needed for the stellar remnant is estimated accordingly.  In each cases, we simulate the results (i.e. retention or ejection) for 10000 supernova explosions producing neutron stars.  We use the equations from Drukier(1996) to calculate retention fractions. The probability that a star is retained to the cluster after applying the supernova \"kick\" is as follows:\\\\ \\begin{equation} \\label{retention} P(ret|\\hat{v}, \\hat{v_{k}}) = \\left\\{ \\begin{array}{ll} 1 & \\hat{v_{k}} -1 \\leq -\\hat{v};\\\\ \\frac{1 - (\\hat{v} -\\hat{v_{k}})^{2}}{4\\hat{v}\\hat{v_{k}}} & |\\hat{v_{k}} -1| < \\hat{v};\\\\ 0 & \\hat{v_{k}} -1 \\geq \\hat{v} \\end{array} \\right\\}, \\end{equation} where $\\hat{v_{k}} = \\frac{v_{k}}{v_{e}}$ and $\\hat{v} = \\frac{v}{v_{e}}$, meaning both the pre-kick velocity and the kick velocity are divided by the escape velocity. $\\hat{v}$ is taken to be a Maxwellian distribution between 0 and 1. Any $\\hat{v} > 1$ would escape from the cluster, therefore we use a lowered Maxwellian distribution.  The escape velocities are taken from our Plummer and King models. We use the central escape velocity for both models, but we also use a Plummer model with no primordial mass segregation as explained in the previous subsection. The King model escape velocities are the highest (typically 25 km/s), because most of the GC cores are more concentrated than the cores of the Plummer models (typically 15 km/s). The model with no primordial mass segregation has the lowest escape velocities (typically 8 km/s), because the supernova progenitors are not concentrated in the center of the core.  As a simple way to test the method we considered a constant kick velocity distribution. We tried 10 progressively larger kick velocity intervals. The intervals we used are [0,x] km, with x = 500, 600, ...., 1400. We use the 3 different models: King, P-center (Plummer model, all heavy stars in the core) and P-spread (Plummer model with the stars spread through the GC, no primordial mass segregation) in most cases, although since the different models result in qualitatively similar conclusions, we do not exhaustively present all the results. To each we fit a power law for the retention fraction as a function of mass (figure \\ref{ret}).  The plots for all other kick velocity distributions are qualitatively similar, so for the remainder of the paper, we present tables showing the results rather than plots. We obtain the exponent and the average retention fraction. The results are in table \\ref{constant}. The exponents are all approximately 0.50. This is because the escape velocity goes as the square-root of the cluster mass. The King model has a slightly higher exponent because more massive clusters tend to favor a slightly higher central potential and are more concentrated, resulting in a somewhat higher escape velocity. The average retention fraction is linear with x. This is as expected. If a constant probability distribution is spread out over an interval twice as large (0-1000 instead of 0-500), one expects half as many neutron stars to be retained.  \\begin{figure} \\begin{center} \\renewcommand{\\epsfsize}[2]{0.7#1} \\epsfig{file=5298fig5.ps,width=6cm,angle=270} \\caption{The retention fraction as a function of mass, fitted to the King-model data with a constant kick velocity distribution between 0 and 1000 km/s. Chi-square was minimized to obtain the fit.} \\label{ret} \\end{center} \\end{figure}  \\begin{table*}[!h] \\label{constant} \\caption{Retention fraction fit results for a constant kick velocity distribution} \\begin{center} \\begin{tabular}{|c|c|c|c|c|c|c|} \\hline & \\multicolumn{2}{c|}{King} & \\multicolumn{2}{c|}{P-center} & \\multicolumn{2}{c|}{P-spread} \\\\\\hline x &exponent&fraction&exponent&fraction&exponent&fraction\\\\\\hline 500 & 0.53 & 0.041 & 0.50 & 0.018 & 0.50 & 0.0053 \\\\\\hline 600 & 0.53 & 0.034 & 0.50 & 0.015 & 0.50 & 0.0044 \\\\\\hline 700 & 0.53 & 0.029 & 0.50 & 0.013 & 0.50 & 0.0038 \\\\\\hline 800 & 0.53 & 0.026 & 0.50 & 0.011 & 0.50 & 0.0033 \\\\\\hline 900 & 0.53 & 0.023 & 0.50 & 0.010 & 0.50 & 0.0029 \\\\\\hline 1000 & 0.53 & 0.021 & 0.50 & 0.009 & 0.50 & 0.0026 \\\\\\hline 1100 & 0.53 & 0.019 & 0.50 & 0.008 & 0.50 & 0.0024 \\\\\\hline 1200 & 0.53 & 0.017 & 0.50 & 0.008 & 0.51 & 0.0022 \\\\\\hline 1300 & 0.53 & 0.016 & 0.50 & 0.007 & 0.51 & 0.0020 \\\\\\hline 1400 & 0.53 & 0.015 & 0.50 & 0.006 & 0.51 & 0.0019 \\\\\\hline \\end{tabular} \\end{center} \\end{table*} We have made similar studies of the effects of linear kick distributions (i.e. where the probability distribution varies linearly, with a maximum at $v=0$, and a probability of zero above some characteristic velocity) and Gaussian velocity kick distributions (which are taken to be Gaussian probability density functions with the characteristic velocity equal to the spread in the velocity distribution).  The linear distributions are characterized in terms of a characteristic velocity $x$, and are presented for large and small velocities in tables \\ref{linear} and \\ref{linear2} respectively.  The linear distributions, like the constant probability density distributions, generically produce a $M^{0.5}$ retention fraction mass dependence, unless the characteristic velocity is less than the escape velocity of the most massive clusters, in which case the dependence will be flatter.  Both these functional forms are thus consistent with the observations of extragalactic X-ray binaries, but will most likely have trouble explaining the dearth of very low velocity pulsars, and simultaneously explaining the small, but non-neglible numbers of very high velocity pulsars (e.g. those with velocities of more than 800 km/sec) while still producing retention fractions greater than $\\sim$ 1\\%, which is needed to explain the large numbers of X-ray binaries and millisecond pulsars in globular clusters.  Gaussian distributions produce a wide range of mass dependences, since for a non-zero Gaussian, there typically will exist a positive slope in the probability density distribution over the range of globular cluster escape velocities.  The results of Gaussian distribution studies are shown in table \\ref{Gaussian}.  It is clear that to reproduce the observations, a Gaussian distribution would need a characteristic velocity just less than 20 km/sec.  This is, on its face, strongly inconsistent with the proposed existing kick velocity distributions, but, upon consideration of the effects of binaries, it may become consistent with e.g. the results of Arzoumanian et al. (2002).  The presence of a heavy binary companion at the time of the supernova explosion will dilute the velocity kick applied to the system, making retention more likely; this has already been suggested to be a big part of the solution to the retention fraction problem (e.g. Davies \\& Hansen 1998; { Pfahl et al. 2002}).  The maximum possible dilution would occur if the envelope of the proto-neutron star has been completely stripped in the binary, while the companion star is very close to the maximum possible mass in order not to undergo a supernova explosion itself later.  That is to say, the maximum dilution occurs for a supernova explosion of a 1.4 $M_\\odot$ core with an 8 $M_\\odot$ companion, and this dilution will be a factor of 6.7 [i.e. 1.4/(1.4+8.0)].  This would be enough to bring a 90 km/sec kick (the characteristic kick velocity of the slow mode found by Arzoumanian et al. 2002) down to about 13 km/sec {(we do note that the Arzoumanian model is a Maxwellian, rather than a Gaussian, and a 90 km/sec Maxwellian will have slightly few low velocity objects than a 90 km/sec Gaussian, but that this should lead to only a minor difference)}.  Thus we can say that the most extreme dilution factor is more than sufficient to produce a reasonable dependence of the retention fraction on mass.  Additionally, heavy binaries will have lower initial velocities than the stars in the cluster as a whole, which should allow them to accomodate larger kicks (although in no case will this increase the kick velocity allowed by more than $\\sqrt{2}$).  Since mass segregation will bring the most massive stars to the centers of clusters, and these massive stars will have the highest cross-sections for interactions, it is likely that supernova progenitors will find themselves in binaries with other massive stars. Whether this actually happens to the extent required by the data must be addressed through numerical simulations and is beyond the scope of this paper.  Regardless, it seems that in no case will the data be consistent with a single Maxwellian or Gaussian with a post-dilution characteristic velocity greater than 20 km/sec.  This seems to rule out the Hobbs et al. (2005) suggestion that the pulsar velocity kick distribution can be a single Maxwellian at 265 km/sec, unless the slower kicks in globular clusters are related to binary evolution (see e.g. Pfahl et al. 2002b; Dewi et al. 2005 for suggestions that binary evolution may affect pulsar kick velocities).  It also indicates that the contribution of neutron stars from the fast mode in a distribution like that of Arzoumanian et al. (2002) to retention fractions should be negligible.  { These results are thus quite similar to the conclusions of Pfahl et al. (2002) who found that, while a single fast kick mode had trouble producing neutron star retention fractions large enough to match the observed numbers of millisecond pulsars in globular clusters like 47~Tuc, including also a slow velocity kick mode, and the effects of binaries could potentially solve the retention problem.}  With careful numerical simulations, the globular cluster X-ray binaries are likely to present the strongest constraints on the low velocity end of the pulsar kick velocity distribution, since field studies will always be complicated, e.g., by the intrinsic velocity dispersion in the field.  To obtain strong constraints on the actual asymmetric kick velocity applied to neutron stars at the times of their birth will require numerical simulations which can take account of both the distributions of the neutron star progenitors in the cluster at the time the supernovae occur, and what fraction of the neutron stars are in binaries, and what are the masses of those binary companions.  The numerical simulations will also need to be made for ranges of cluster masses and initial concentration indices.  At the present, the best studies of neutron star retention (e.g. Drukier 1996; Davies \\& Hansen 1998; Pfahl, Rappaport \\& Podsiadlowski 2002) have each considered several aspects of this problem; the simulation closest to meeting these requirements was that of Pfahl et al. (2002), in which all the relevant factors except for a range in cluster escape velocities were considered.    \\begin{table*}[!h] \\label{linear} \\caption{Retention fraction fit results for a linear kick velocity distribution with negative slope} \\begin{center} \\begin{tabular}{|c|c|c|c|c|c|} \\hline & \\multicolumn{2}{c|}{King} & \\multicolumn{2}{c|}{P-center}\\\\\\hline x & exponent&fraction&exponent&fraction\\\\\\hline 100 & 0.44 & 0.348 & 0.47 & 0.168 \\\\\\hline 200 & 0.49 & 0.187 & 0.48 & 0.087 \\\\\\hline 300 & 0.50 & 0.129 & 0.49 & 0.059 \\\\\\hline 400 & 0.51 & 0.098 & 0.49 & 0.044 \\\\\\hline 500 & 0.51 & 0.080 & 0.49 & 0.036 \\\\\\hline 600 & 0.51 & 0.066 & 0.49 & 0.030 \\\\\\hline 700 & 0.51 & 0.057 & 0.49 & 0.026 \\\\\\hline 800 & 0.52 & 0.050 & 0.49 & 0.022 \\\\\\hline 900 & 0.52 & 0.045 & 0.49 & 0.020 \\\\\\hline 1000 & 0.52 & 0.040 & 0.49 & 0.018\\\\\\hline \\end{tabular} \\end{center} \\end{table*}  \\begin{table}[!h] \\label{linear2} \\caption{Retention fraction fit results for a linear kick velocity distribution with very large negative slope, for a King model} \\begin{center} \\begin{tabular}{|c|c|c|} \\hline & \\multicolumn{2}{c|}{King}\\\\\\hline x & exponent&fraction\\\\\\hline 10 & 0.11 & 0.90\\\\\\hline 20 & 0.20 & 0.79\\\\\\hline 30 & 0.27 & 0.70\\\\\\hline 40 & 0.31 & 0.61\\\\\\hline 50 & 0.35 & 0.54\\\\\\hline 60 & 0.38 & 0.49\\\\\\hline 70 & 0.40 & 0.44\\\\\\hline 80 & 0.42 & 0.40\\\\\\hline 90 & 0.43 & 0.37\\\\\\hline 100 & 0.44 & 0.34\\\\\\hline \\end{tabular} \\end{center} \\end{table}    \\begin{table*}[!h] \\label{Gaussian} \\caption{Retention fraction fit results for a Gaussian kick velocity distribution.  The exponent values with asterisks denote cases where the retention fractions are quite small and the estimations suffer from small number statistics.} \\begin{center} \\begin{tabular}{|c|c|c|c|c|c|c|} \\hline & \\multicolumn{2}{c|}{King} & \\multicolumn{2}{c|}{P-center} & \\multicolumn{2}{c|}{P-spread} \\\\\\hline $\\sigma$ & exponent&fraction&exponent&fraction&exponent&fraction\\\\\\hline 20 & 0.68 & 0.260 & 1.15 & 0.06 & 1.29 & 0.0058 \\\\\\hline 40 & 1.06 & 0.051 & 1.43 & 0.005 & 1.46 & 0.0008 \\\\\\hline 60 & 1.18 & 0.017 & 1.52 & 0.001 & 1.57 & 0.0002 \\\\\\hline 80 & 1.23 & 0.007 & 1.60* & 0.000 & 1.66* & 0.0001 \\\\\\hline 100 & 1.26 & 0.003 & 1.69* & 0.000 & 1.77* & 0.0000 \\\\\\hline 120 & 1.28 & 0.002 & 1.81* & 0.000 & 1.89* & 0.0000 \\\\\\hline 140 & 1.30 & 0.001 & 1.94* & 0.000 & 2.01* & 0.0000 \\\\\\hline 160 & 1.30* & 0.001 & 2.09* & 0.000 & 2.13* & 0.0000 \\\\\\hline 180 & 1.32* & 0.001 & 2.25* & 0.000 & 2.26* & 0.0000 \\\\\\hline 200 & 1.33* & 0.000 & 2.42* & 0.000 & 2.38* & 0.0000 \\\\\\hline \\end{tabular} \\end{center} \\end{table*} }"
    },
    "0608/astro-ph0608539_arXiv.txt": {
        "abstract": "The astrophysical origin of the \\textit{r}-process nuclei is still unknown. Even the most promising scenario, the neutrino-driven winds from a nascent neutron star, encounters severe difficulties in obtaining requisite entropy and short dynamic timescale for the \\textit{r}-process. In this study, the effect of anisotropy in neutrino emission from a proto-neutron star surface is examined with semi-analytic neutrino-driven wind models. The increase of neutrino number density in the wind owing to the anisotropy is modeled schematically by enhancing the \\textit{effective} neutrino luminosity. It is shown that the neutrino heating rate from neutrino-antineutrino pair annihilation into electron-positron pairs can significantly increase owing to the anisotropy and play a dominant role for the heating of wind material. A factor of five increase in the effective neutrino luminosity results in $50 \\%$ higher entropy and a factor of ten shorter dynamic timescale owing to this enhanced neutrino heating. The nucleosynthesis calculations show that this change is enough for the robust \\textit{r}-process, producing the third abundance peak ($A = 195$) and beyond. Future multi-dimensional studies with accurate neutrino transport will be needed if such anisotropy relevant for the current scenario (more than a factor of a few) is realized during the wind phase ($\\sim 1-10\\, \\textrm{s}$). ",
        "introduction": "The astrophysical site of the rapid-neutron-capture nucleosynthesis (\\textit{r}-process), which accounts for about half of nuclei heavier than iron, has been a long-standing mystery. During the last decade, the neutrino-heated ejecta from a nascent neutron star \\citep[neutrino-driven winds,][]{Woos94} has been considered to be the most promising astrophysical site for the \\textit{r}-process. Previous studies show, however, severe problems in obtaining requisite high entropy and short dynamic timescale for the production of heavy \\textit{r}-process nuclei \\citep{Qian96, Otsu00, Sumi00, Wana01, Thom01}. The general relativistic effect for a very compact proto-neutron star \\citep[e.g., the mass of $2.0\\, M_\\odot$ with the radius of 10~km,][]{Otsu00, Wana01} or a magnetar-like magnetic field strength \\citep{Thom03, Suzu05} have been invoked to increase entropy and reduce the dynamic timescale of the winds. It is questionable, however, if such physical conditions can be the general requirements for the \\textit{r}-process nucleosynthesis.  \\citet{Qian96} have suggested that an additional energy input to the neutrino-driven wind at between 1.5 and 3 times the neutron star radius is efficient to increase entropy and reduce dynamic timescale of the wind material. In this \\textit{Letter}, it is shown that strong anisotropy in neutrino emission from the proto-neutron star, \\textit{if it exists}, acts as this extra energy source and helps the \\textit{r}-process. The neutrino-driven wind model with spherically symmetric, steady outflow approximation is used to obtain the wind trajectories (\\S~2). A sudden increase of neutrino number density in winds owing to anisotropic neutrino emission is modeled by enhancing the neutrino luminosity. Nucleosynthesis calculations with the obtained thermodynamic trajectories are performed to demonstrate this effect (\\S~3). Finally, a possible origin of this anisotropy in neutrino emission and some implications of this study are discussed (\\S~4). ",
        "conclusions": ""
    },
    "0608/astro-ph0608154.txt": {
        "abstract": "\\label{abs}  We present rotation periods for the solar-type primary stars in 13 close ($a \\la 5$ AU) single-lined spectroscopic binaries with known orbital periods ($P$) and eccentricities ($e$). All binaries are members of the open clusters M35 (NGC2168; $\\sim$ 150 Myr) and M34 (NGC1039; $\\sim$ 250 Myr). The binary orbital parameters and the rotation periods of the primary stars were determined from time-series spectroscopy and time-series photometry, respectively. Knowledge of the ages, orbital periods, and eccentricities of these binaries combined with the rotation periods and masses of their primary stars makes them particularly interesting systems for studying the rates of tidal circularization and synchronization. Our sample of 13 binaries includes six with orbital periods shortward of 13 days ($a \\la 0.12$ AU). The stars in these binaries orbit sufficiently close that their spins and orbits have evolved toward synchronization and circularization due to tidal interactions. We investigate the degree of tidal synchronization in each binary by comparing the angular rotation velocity of the primary stars ($\\Omega_{\\star}$) to the angular velocity expected if the primary star was synchronized ($e = 0$) or pseudo\u00adsynchronized ($e > 0$) with the orbital motion ($\\Omega_{ps}$). Of the six closest binaries two with circular orbits are not synchronized, one being subsynchronous and one being supersynchronous, and the primary stars in two binaries with eccentric orbits are rotating more slowly than pseudosynchronism. The remaining two binaries have reached the equilibrium state of both a circularized orbit and synchronized rotation. As a set, the six binaries present a challenging case study for tidal evolution theory, which in particular does not predict subsynchronous rotation in such close systems. ",
        "introduction": "\\label{intro}  Tidal and dissipative forces in close detached binary stars drive an exchange of angular momentum between the rotation of the stars and their orbital motion. The cumulative effects of such tidal interactions with time is referred to as tidal evolution. The characteristic signs of tidal evolution are: 1) Alignment of the stellar spin axes perpendicular to the orbital plane; 2) Synchronization of the rotation of the stars to the orbital motion; 3) Circularization of the orbits. In a population of coeval detached late-type binary stars tidal evolution can be observed among the closest binaries ($a \\la 0.2$ AU).  Observations of {\\it tidal circularization} have been the primary constraint on tidal theory over the past two decades \\citep[e.g.][]{mm05,mmd04,lst+02,mca+01,dmm92,mdl+92}. In particular, the distribution of orbital eccentricities with orbital periods (the $e-\\log(P)$ diagram) has provided clear evidence for tidal circularization in homogeneous and coeval populations of late-type binaries, and enabled a robust measure of the degree of circularization integrated over the lifetime of the binary population as a function of orbital period. \\citet{mm05} define the {\\it tidal circularization period} of a binary population as the longest orbital period to which binaries with initial eccentricities of e=0.35 circularize at the age of the population. Importantly, current theories of tidal circularization cannot account for the distribution of tidal circularization periods with population age.  In comparison, the amount of observational data suitable for measuring the rate of {\\it tidal synchronization} in late-type binaries is sparse. This is in part because binary orbital elements are simpler to obtain than rotation periods for stars in binaries. Even so, observations of the synchronization of the stars in a binary system provide a second window on the tidal effects on that system, and one of major importance for several reasons. % First, tidal theory makes explicit predictions for the relative rates of tidal circularization and synchronization that can be straightforwardly tested \\citep{ws02,zb89}. For a typical binary the two rates differ and the predicted evolutionary paths from an asynchronous, eccentric binary to a synchronized and circular binary take the stars in and out of synchronism during the evolutions of the stars, their orbital separation, and their orbital eccentricity (see Figure 1 in \\citet{zb89} and Figures 1 and 3 in \\citet{ws02}). % Second, the range in binary separations over which tidal synchronization can significantly affect the stellar angular momentum evolution provides an important constraint on the impact of binarity on stellar angular momentum evolution in late-type stars. % Third, synchronization of the observable surface layers is closely linked to the internal angular momentum transport in the star. Thus, the rate of synchronization of the surface layers can shed light on the coupling between those layers and the stellar interior. % Fourth, observations of magnetic field tracers (X-rays, chromospheric emission, etc.) during rapid stellar evolution in tidally synchronized binaries may enable determination of the evolution rates of stellar dynamos.  Finally, significant tidal synchronization and circularization is expected in the many star-planet systems where the planets orbits very close to their late-type host stars (i.e., ``hot Jupiter'' systems). Tidal interactions in star-planet systems may therefore play an important role in determining the observed distributions of mass, orbital period, and eccentricity of extra-solar planets \\citep[e.g.][]{ol04}.  To date, most published studies of synchronization have focused on early-type binaries \\citep[e.g.][]{ab04,alg02,gmm84a,gmm84c,levato74}. This emphasis is due, in part, to telescope and instrument capabilities which in the past have favored bright and rapidly rotating stars. Also, the use of archived data on eclipsing binaries has introduced a bias toward higher mass stars.  However, a few recent studies of tidal synchronization, e.g \\citet{gmm84d}, \\citet{cgc95} and \\citet{cc97}, include binaries with late-type main-sequence primary stars. These studies represent important contributions to the study of tidal synchronization. \\citeauthor{gmm84d} studied 43 detached double-lined eclipsing and non-eclipsing binaries with at least one late-type stellar component. They found that the observed degree of synchronism was in agreement with the theoretical predictions of \\citet{zahn77}. In half of the binaries in their study the primary or secondary star, or both, have evolved off the main sequence and stellar rotation was derived from either line-broadening ($v\\sin(i)$) or from periodic brightness variations. For the non-eclipsing binaries in their study, the stellar rotation velocities measured from line broadening are less suitable for constraining models of tidal synchronization, because of the ambiguities introduced by the unknown inclination of the rotation axis and non-rotational line-broadening due to the secondary spectrum.  \\citeauthor{cgc95} and \\citeauthor{cc97} studied tidal synchronization using eclipsing binary data from \\citet{a91}, of which 10 systems have late-type stellar components. The two studies compared the observed binary parameters against the model predictions of \\citet{tassoul87,tassoul88} and \\citet{zahn89}, respectively, and found general agreement between the observed level of synchronization and the predicted timescales for synchronization. Astrophysical parameters required to determine the theoretical time-scale for which synchronization is achieved were obtained by comparing theoretical models \\citep{claret95,cg95} directly to each star of the Andersen sample. Rotation velocities ($v\\sin(i)$) for all stars were derived from line-broadening.  Knowledge of binary ages is critical in order to measure the rate and evolution of tidal synchronization.  Furthermore, because of the sensitive dependence of tidal effects to stellar radius, knowledge of the stellar evolutionary state and history is essential for testing tidal evolution. For example, tidal theory predicts that the pre main-sequence (PMS) and the early main-sequence ($t \\la 500$ Myr) are the most active phases of tidal evolution. The PMS phase because of the large radii and deep convective envelopes, and the early main-sequence phase because the stars are spinning super-synchronously after contracting onto the zero-age main sequence. Similarly, post-main-sequence evolution of one or both components of a binary greatly increases the rate of tidal evolution.  Thus an optimal binary sample for the study of tidal evolution would comprise a coeval population of late-type binaries with accurate information about age, evolutionary stage, orbital parameters, and rotational angular velocities of the stars. Arguably such binary samples with ages $t \\la 500$ Myr are particularly interesting. Given these needs, young open clusters are superb laboratories for the study of tidal evolution.  So motivated, we have undertaken parallel spectroscopic and photometric surveys of the open clusters M35 ($\\alpha_{2000} = 6^{h}~9^{m}$, $\\delta_{2000} = 24\\degr~20\\arcmin$) and M34 ($\\alpha_{2000} = 2^{h}~42^{m}$, $\\delta_{2000} = 42\\degr~46\\arcmin$) to derive orbital periods and eccentricities as well as stellar rotation periods for late-type binary stars. The goal of this study is to measure the degree of tidal synchronization of the primary stars in the closest binaries ($a \\la 0.2$ AU) and to compare the results to the predictions from tidal theory.  M35 \\citep[$150$ Myr;][]{vss+02,deliyannis06} and M34 \\citep[$250$ Myr;][]{steinhauer06} provide populations of close late-type binaries with ages during the most active phase of tidal evolution, making them attractive targets for observational testing of models of tidal synchronization. M35, in particular, provides a rich population of close binaries which has allowed determination of a well-defined tidal circularization period at $10.2^{+1.0}_{-1.5}$ days \\citep{mm05}. Indeed, eight out of the nine M35 binaries with periods less than $\\sim$ 10 days have been circularized to eccentricities less than 0.05. In M34, our ongoing spectroscopic survey has led to the discovery of five circular binaries ($e < 0.1$) with periods of less than 5.5 days. A tidal circularization period has not yet been determined for M34.  We begin by briefly introducing the current theories of tidal evolution in Section~\\ref{theory}. Section~\\ref{obs} outlines our observational program and Section~\\ref{results} describe our observational results. In Section~\\ref{effects} we address potential complications of measuring the rotation periods of stars in close binary systems. In Section~\\ref{loop} we evaluate the degree of tidal synchronization and circularization for the closest binaries in M35 and M34 and introduce the $\\log(\\Omega_{\\star}/\\Omega_{ps})-\\log(P)$ diagram which presents the dependence on stellar separation. We compare the observed tidal evolution of individual binaries to the predictions of current tidal theory in Section~\\ref{discussion}. Section~\\ref{conclusions} summarizes and presents our conclusions.  % ---------------------------------------------------------------------------- ",
        "conclusions": ""
    },
    "0608/astro-ph0608363_arXiv.txt": {
        "abstract": "We present a physical framework that can account for most of the observed spectral properties of the prompt gamma-ray burst emission. This includes the variety of spectral shapes, shape evolutions, and spectral correlations between flux and spectral peak,  within bursts described by Borgonovo \\& Ryde, and among bursts described by Amati/Ghirlanda. In our proposed model the spectral peak is given by the photospheric emission from a relativistic outflow for which the horizon length is much smaller than the radial width. The observed duration of the thermal flash will be given by the radial light-crossing time. This then gives that the typical emission site is at $\\sim 10^{11}$cm from the center, with a Lorentz factor of $\\sim 300$. This emission is accompanied by non-thermal emission from dissipation locations outside the photosphere. The relative strength of these two components depend on injection effects at the central engine leading to varying relative location of the saturation and photospheric radii. The total emission can then reproduce the observed variety. The spectral correlations are found by assuming that the amount of energy dissipated depends non-linearly on the averaged particle density. Beside the spectral correlations this also gives a description of how the relative strength of the thermal component varies with temperature within a burst. ",
        "introduction": "The emission below the spectral peak in gamma-ray bursts (GRBs), which typically lies at around $200 - 300$ keV, is indicative of the radiative process in charge of the spectrum formation. In this sense a truly crucial observation is that the power-law slope of the sub-peak emission in a fair fraction of observed, instantaneous spectra are harder than what is expected from optically-thin synchrotron emission \\citep{preece00}. Indeed, the {\\it Compton Gamma Ray Observatory} BATSE instrument [25--2000 keV] recorded spectra which were as hard as a thermal emission throughout the burst \\citep{ryde04} as well as during the initial interval of bursts \\citep{ghirlanda, ryde04, kaneko05, bosnjak}. Furthermore, {\\it Beppo-SAX} [2-700 keV] detected several bursts which exhibited very hard low-energy spectra (see e.g. \\citet{frontera00}).  Moreover, observations made by {\\it Ginga} [$\\sim$ 2-400 keV] \\citep{strohmayer} indicated a substantial number of bursts with breaks below 10 keV, i.e., below the observable range of BATSE. The authors suggested that bursts might have two different spectral breaks, one in the BATSE range and one below this, close to 5 keV. Similar results are reported by {\\it Swift} [0.3-500 keV, with {\\it KONUS}] for GRB050820A (Osborne et al. 2005) whose spectrum indicates that there can be an additional spectral break around 10 keV. These evidence indicate that the high-energy spectra of GRBs are more complicated than that of a single broken power-law. Similarly, the need for additional high-energy spectral components beyond 1 MeV in several bursts (e.g. \\citet{hurley, atkins, gonz}) reinforces the same conclusions.  Another crucial point was raised by Ryde (2005) who argued that a strong thermal component could very well be a common feature in the $\\gamma$-ray spectra of GRBs. Most of the strong, pulsed bursts analyzed could be described by a two-component model consisting of a thermal component, modelled by a Planck function, and a non-thermal component, modelled by a power-law function, the sum of which creates the observed spectral shapes and evolutions. The typical peak in a GRB spectrum is thus argued to be caused by a thermal component. The behaviors of the two separate components on their own are, in general, similar from burst to burst, even though the combined behavior leads to the large variety in spectral evolutions that is observed. The typical behavior of the temperature over a single pulse episode is, with a few exceptions, that it decreases monotonically as a broken power-law in time. The non-thermal emission component, which is parameterized by a power-law photon-index $s \\equiv d\\,\\log N_{\\rm E}/d\\,\\log E$, where $N_E$ is the photon flux, exhibits a typical decrease in $s$ from approximately $-1.5$ to approximately $-2$. Other behaviors do occasionally exist (see Ryde et al. in prep.) These results also hold for bursts which are consistent with optically-thin synchrotron emission, in the sense that spectral fits with a broken power-law function gives a photon-index of the low-energy power-law which is smaller (softer) than $\\alpha = -2/3$.  Such thermal--non-thermal hybrid emission is indeed common in astrophysical sources and is a natural consequence of heating of a background plasma and acceleration of particles. This is, for instance, typical for solar flares, accretion disks and cluster of galaxies (see, e.g., Petrosian \\& Liu 2004). In the case of GRBs, the thermal component is thought to stem from the fireball photosphere while the non-thermal component stems from accelerated electrons beyond the photospheric radius, in the optically-thin region.  We will denote this model as the photosphere model since the main feature of it, its energy flux peak, and spectral evolution is given by the thermal component. This certainly emanates from the photosphere, where the optical depth is unity, even though the finally observed emission can be altered due to scattering and diffusion effects.  If we are indeed observing the photospheric, thermal emission, as suggested by the analyses above, it simplifies the physical interpretation since we are detecting the fireball and its evolution directly. This is in contrast to the internal shocks which give us indirect information through the dissipation processes, randomly distributed through out the fireball wind, which by necessity complicates the interpretation. However, the origin of the thermal component and its observed evolution is not obvious. Strong photospheric components were predicted in the spectra of GRBs, in kinetic models by \\citet{MR00, MRR, DM, ryde04}, and in Poynting-flux dominated models by  e.g. \\citet{DS}. More recently, discussions on various scenarios and theoretical aspects have been given by, e.g., \\citet{RM05, thomson, asaf, TMR06}. In this paper we will assume that the thermal component is predominantly from close to the photosphere and that the evolution we detect is mainly due to change in the property of the photosphere. First, we discuss the observational and statistical significance of the thermal fits in \\S \\ref{sec:sign}. We then discuss some basic properties of the fireball evolution that we can learn from the observed $\\gamma$-ray behavior in \\S \\ref{sec:theory} and derive several of the correlations that have been observed for bursts. We discuss our results in \\S \\ref{sec:discussion} and comment on alternative models. ",
        "conclusions": "\\label{sec:discussion}  Even though a major part of the total energy in a GRBs is emitted during the prompt phase,  consensus has not yet been reached regarding the emission sites and mechanisms. For a model to successfully describe this phase the following observational facts has to be appropriately addressed:  \\noindent (i) Existence of hard spectral slopes below the peak energy, i.e., large values of the parameter $\\alpha$, giving rise to the line-of-death problem.  \\noindent (ii)  Beside producing both hard and soft sub-peak spectra, the model should be able to reproduce more complicated spectral shapes such as the ones in Figs. \\ref{fig:decomp} and \\ref{fig:decomp2}. In addition, the shape of the spectrum around the peak that should not be too broad. The values of the model parameter that are inferred by the observations should have a distribution that has a reasonable explanation.  \\noindent (iii) The {\\it evolution} of the spectral shape. This could be represented by parameters of non-thermal models, such as the Band-function $\\alpha$, or parameters of the photosphere model, such as $kT$ and/or $s$. In the former case the $\\alpha$ parameter translates into physical parameters depending on the chosen  model.  \\noindent (iv) The distribution of peak energies; why the peaks mainly fall into the range they do.  \\noindent (v) The correlation between peak energy and flux during a burst, i.e. the hardness-intensity correlation, HIC, which is characteristic for every burst \\citep{RS00, BR01} .  \\noindent (vi) The Amati/Ghirlanda \\citep{amati, ghirlanda} relations for an ensemble of bursts should be predictable on reasonable grounds by the model.  \\noindent (vii) Burst and pulse durations. The derived radii at which the emission process takes place as well as the derived Lorentz factors should not lead to the compactness problem.  \\noindent (viii) The microphysics implied: does the cooling occur in situ with the acceleration or do they occur in separated regions? What underlying particle distributions is needed to describe the observations, given the acceleration and dissipation processes?  Here we have presented a physical framework that addresses most of these points. In particular, it can explain pulses that are thermal through out their duration. The analysis made above in \\S \\ref{sec:sign} and in Ryde (2004, 2005) indicate that the temperature has a canonical behavior with the temperature decaying as a broken power-law, independent of the apparent behavior of the burst. In the model presented here this simply reflects the temperature distribution within the shell. These results are for the class of long bursts (see e.g. \\citet{balazs}) and the corresponding behavior for short and intermediate duration bursts \\citep{Horv06} will be discussed elsewhere.  For the bursts discussed in this paper, the peak energy was observed to be within the BATSE energy range. However, the case could very well be that $kT$ is actually outside of that range, say at energies lower than 10 keV. In such a case the spectrum within the BATSE window would mainly reflect the non-thermal emission. The thermal component would reveal itself by an upturn in the spectrum at low energies, similarly to the case in the lower right-hand panel in Figure \\ref{fig:960530:2}. We have also assumed that the thermal emission is not greatly affected by a scattering photosphere, to the extent that it pertains its Planckian shape as indicated by the observations. However, strong scattering can be important and has been discussed elsewhere \\citep{RM05, asaf}.  In this connection, it is worth mentioning the work by Vetere et al. (2005) who studied rapidly varying burst which are accompanied by a slowly varying component. They suggest that the latter component is thermal and associated with the photosphere while the former, variable component is due to external shocks, giving non-thermal spectra. Such a picture fits in nicely with the model presented above and deserves further investigation.  What alternative models does there exist? As shown above a purely optically-thin synchrotron (OTS) emission (isotropic pitch angles) is rejected purely by the spectral shapes. In the canonical diffusive shock model the typical acceleration times are much shorter than the typical synchrotron cooling times. However, the acceleration and the cooling do not necessarily need to occur in the same place. This will lead to electrons acquiring a fast cooling distribution, since the cooling time is shorter that the dynamical time scale \\citep{g00}. This results in a spectral softening, with an expected $\\alpha = -1.5$ instead of $-2/3$. OTS then becomes even less plausible. Another concern with the OTS is that the peak energy will depend on several of the model parameters; $\\Gamma$, $\\gamma_{\\rm min}$, and $B_\\perp$ making it hard to get the preferred energy breaks as observed. \\citet{LP00} discussed the possibility that the hard spectral slopes are due to synchrotron self-absorption, i.e. the electrons absorb the photons themselves, which would lead to harder spectral slopes in the photon fluxes; $E^1$ or $E^{3/2}$ \\citep{RL}. They noted that to get a high optical-depth to self-absorption in the BATSE energy range somewhat extreme parameters have to be used, in particular, one needs to invoke a very high magnetic field. At a dissipation distance of $10^{13}$ cm, particle density of $10^8$ cm$^{-3}$, and with  $\\Gamma = 1000$ one would need approximately $10^8$ G. Note that the generation of magnetic fields in the outflow, in particular such strong fields, is not well understood (see e.g. \\citet{Trier}). Furthermore, the Compton $y$-parameter will in such a case be of order unity which thus reduces the efficiency in the observed range since emission will be up-scattered to higher energies (Kumar 1999).  Another efficient radiative mechanism is inverse Compton scattering: the gamma-ray spectrum could instead be an inverse Compton image of the synchrotron spectrum, now located in the optical since the energy increase goes as $\\gamma^2$. Such a synchrotron self-Compton model was explored by \\citet{PM00}. The low-energy slope will be the same as for the synchrotron seed spectrum, $\\alpha < -2/3$. But if the latter is self-absorbed, a situation which now requires much lower magnetic field strengths, the low-energy spectrum could be as hard as $\\alpha = 0$ \\citep{J68}. However, as \\citet{BaringB} pointed out, the spectrum around its peak is notably broader than what is observed. Alternatively, the inverse Compton emission could be seeded by soft photons with a narrow energy distribution, i.e. a quasi--mono-energetic distribution, producing a less broad spectrum and with a limiting value again at $\\alpha = 0$. Compared to synchrotron emission such an emission mechanism also alleviates the requirements on the magnetic field strength.  Yet another alternative is the small pitch-angle synchrotron-emission (SPA) \\citep{E73, EP73} and similarly jitter radiation \\citep{Med00, Med06}, which produce hard spectral slopes as well, $\\alpha < 0$. Here the $\\gamma$-dependant distribution of pitch angles is needed as well as a scenario to set up a distribution for which the pitch angles are not much greater than $\\gamma^{-1}$. This could be done by rapid cooling and inefficient diffusion transverse to the magnetic field \\citep{LP00}.  The latter two models are thus the main contenders for describing the prompt phase with a purely non-thermal model. Diffusive shock acceleration (Fermi processes) is assumed to accelerate electrons from a shocked thermal population into a non-thermal distribution, which cool by these mechanisms and emit the observed radiation. \\citet{BaringB} noted, however, that both these emission processes need an almost purely non-thermal electron distribution to be able to fit the observed spectra. This is difficult to reconcile with shock models, which often have a strong contribution of a thermal population. We note here that the non-thermal component in the photosphere model, presented in this paper, does not have its break energy within the BATSE window. Therefore such fitting constraints are not found and a thermal component of the electron population can very well be present in that case.  It will be necessary to have a broader energy coverage to be able to reach firm conclusions on the actual spectral shape in the X- and  $\\gamma$-rays, and thereby conclusively identifying the emission process. The {\\it GLAST} satellite is expected to provide better clues with its much broader energy range compared to presently available observations."
    },
    "0608/astro-ph0608680_arXiv.txt": {
        "abstract": "We present an $R=45,000$ Keck spectrum of the microlensed Galactic bulge G-dwarf OGLE-2006-BLG-265, which has high ($\\sim 60$) signal-to-noise ratio despite its short (15 min) exposure time because the source was magnified by $A\\sim 135$. While it is very metal-rich ([Fe/H]=0.56), the higher temperature of this star compared with the luminous red giants usually measured in the bulge gives its spectrum many unblended atomic lines. We measure the abundances of 17 elements, including the first abundances for S and Cu in a bulge star. The [$\\alpha$/Fe] ratios are subsolar, while the odd-Z elements are slightly supersolar, trends that are also seen in the more metal-rich stars in the Bulge and the local Galactic disk. Because the star is a dwarf, the [O/Fe], [Na/Fe], and [Al/Fe] ratios cannot be attributed to internal mixing, as is sometimes claimed for giants. Similar high-resolution spectra could be obtained for about a dozen bulge dwarf stars per year by means of well-designed target-of-opportunity observations. ",
        "introduction": "}  While high-resolution spectroscopy is the gold standard of stellar analysis, the daunting photon requirements of this technique generally limit its application to the brightest stars within a given stellar population, i.e., the giants.  The one major exception is stars from the local neighborhood, where the dwarfs can partially compensate for their dimness by their greater number and hence (statistically) their greater proximity.  For distant systems, like globular clusters and the Galactic bulge, high-resolution spectra of dwarfs would require exposure times that are at least 100 to 1000 times longer than those of giants, which in most cases is prohibitive.  Nevertheless, there are compelling reasons to obtain dwarf spectra in these systems.  Element abundances in dwarfs and giants are not necessarily the same even when the progenitors of these stars formed from gas with identical compositions.  Mixing on the giant branch can destroy some elements and dredge up others. Measuring the surface-composition differences between dwarfs and giants in the same population is an important test of theories of stellar evolution. Moreover, even when the dwarf and giant element abundances are identical, the {\\it derived} values may differ if the atmospheric models used to estimate the abundances are in error.  Hence, comparing dwarf and giant spectra is also a powerful way to test atmospheric models. Driven by these considerations, several groups (e.g., \\citealt{boesgaard:98}), have, by prodigious efforts, obtained high-resolution spectra of subgiant and main-sequence turnoff stars in globular clusters.  High-magnification gravitational microlensing events provide a unique opportunity to acquire such spectra. Microlensing events are rare, affecting only about $\\tau\\sim 10^{-6}$ Galactic bulge stars at any given time. The fraction of all events with maximum magnifications $A_\\max\\gg 1$ is only $A_\\max^{-1}$.  Nevertheless, by monitoring more than a hundred million bulge stars, the Optical Gravitational Lens Experiment (OGLE-III)\\footnote{http://www.astrouw.edu.pl/$\\sim$~ogle/ogle3/ews/ews.html} Early Warning System (EWS, \\citealt{ews}) has been finding about 600 microlensing events in each of the last four years, of which about a dozen per year have $A_\\max > 100$.  The Microlensing Observations for Astrophysics\\footnote{http://www.massey.ac.nz/$\\sim$iabond/alert/alert.html} collaboration has recently inaugurated the second phase of their experiment and expects to soon find a comparable number of high-magnification events.  High-magnification microlensing events remain within half their peak magnification for a time interval \\citep{einstein36}, \\begin{equation} t_{1/2} = {\\sqrt{12}\\,t_\\e\\over A_\\max} \\sim 1\\,{\\rm day}\\,{t_\\e\\over 30\\,{\\rm days}} \\biggl({A_\\max\\over 100}\\biggr)^{-1} \\label{eqn:thalf} \\end{equation} where $t_\\e$ is the Einstein timescale of the event. Hence, for typical Einstein timescales ($\\sim 30$ days), high-magnification events will be observable by any given telescope at magnifications $A> 50$.  At substantially higher magnifications, the number of events accessible to an observatory falls as $A^{-2}$, one factor of $A^{-1}$ because the events are intrinsically rarer and a second such factor because the chance that the event will peak near the observatory drops.  Nevertheless, there is at least one example of an $A_\\max=3000$ event, OGLE-2004-BLG-343, that peaked over the Very Large Telescopes in Chile and was basically recognized as such several hours before peak \\citep{ob04343}. Therefore, at least in principle, it would be possible to obtain of order a dozen spectra of dwarf stars in the Galactic bulge per year, each magnified by a factor $A\\sim 100$, with a few at much higher magnifications.  \\citet{benetti:95} took the first spectrum of a microlensing event towards the bulge, albeit at low dispersion, but sufficient to determine the spectral type. \\citet{minniti:98} reported the possible detection of Li in a bulge dwarf. \\citet{cavallo:03} presented a preliminary abundance analysis for six stars with high-resolution spectra including three dwarfs magnified by factors between 2.3 and 30. Here we report on the abundances of 17 elements in a microlensed bulge dwarf star based on high-resolution spectra, including the first measurements of S and Cu.    { ",
        "conclusions": ""
    },
    "0608/astro-ph0608155_arXiv.txt": {
        "abstract": "We have searched for microsecond bursts of emission from millisecond pulsars in the globular cluster M28 using the Parkes radio telescope. We detected a total of 27 giant pulses from the known emitter \\psrc. At wavelengths around 20\\,cm the giant pulses are scatter-broadened to widths of around 2\\,$\\mu$s and follow power-law statistics.  The pulses occur in two narrow phase-windows which correlate in phase with X-ray emission and trail the peaks of the integrated radio pulse-components.  Notably, the integrated radio emission at these phase windows has a steeper spectral index than other emission.  The giant pulses exhibit a high degree of polarization, with many being 100\\% elliptically polarized.  Their position angles appear random. Although the integrated emission of \\psrc\\ is relatively stable for the frequencies and bandwidths observed, the intensities of individual giant pulses vary considerably across our bands.  Two pulses were detected at both 2700 and 3500\\,MHz.  The narrower of the two pulses is 20\\,ns wide at 3500\\,MHz.  At 2700\\,MHz this pulse has an inferred brightness temperature at maximum of $5 \\times 10^{37}$\\,K.  Our observations suggest the giant pulses of \\psrc\\ are generated in the same part of the magnetosphere as X-ray emission through a different emission process to that of ordinary pulses. ",
        "introduction": "\\label{sec:introduction}  The Crab pulsar was discovered through the observation of a series of strong bursts of radio emission \\citep{sr68}.  The sampling interval of this search was about 18 times the 33\\,ms pulse period, so if all of its pulses had been equal to the average pulse it would have been rendered invisible.  Instead, the pulses followed power-law statistics, and so some pulses with energies many times stronger than the average pulse-energy were detected.  These giant pulses (GPs) are composed of $\\lesssim 2$\\,ns ultra-bright shots \\citep{hkwe03} that can occur in emission envelopes that are as long as $\\sim 100$\\,$\\mu$s \\citep{sbh+99,kni06}.  Propagative effects have a frequency-dependent nature.  As the spacing and number of shots are frequency independent, the observed multiplicity of components is due to the emission mechanism itself \\citep{sbh+99}.  The GPs are most readily found in the main-pulse and inter-pulse radio components, which correlate in phase with the optical, X-ray, and \\gree\\ pulse peaks \\citep{lcu+95}.  Other radio pulse-components do not have associated high-energy emission, but some GPs are seen at their phases \\citep{jsk+05}.  The individual nano-shots of GPs from the Crab pulsar have a high degree of circular polarization and come in both orientations \\citep{hkwe03}.  The distribution in degree of circular polarization of unresolved GPs can be accurately modelled by assuming superpositions of many ($\\sim 100$) nano-pulses that are 100\\% circularly polarized \\citep{psk+06}.  The millisecond pulsars B1937+21, J1824$-$2452A, J1823$-$3021A, J0218+4232, and B1957+20 have also been shown to emit short-timescale pulses with inferred brightness-temperatures ($T_{B}$) that are very high \\citep{cstt96,rj01,kbmo05,kbm+06,kni06}.  Although many of these ultra-bright pulses are weaker in pulse-energy terms than the average pulse (see, e.g. Soglasnov \\etal 2004)\\nocite{spb+04}, they can be defined as GPs through a number of other criteria.  These are that they are always unresolved on microsecond timescales, occur in narrow phase-windows that align with X-ray emission, and that their energies follow power-law statistics.  Unlike the GPs of the Crab pulsar, the GPs from the millisecond pulsars do not occur at all phases where integrated emission is seen, and do not contribute markedly to the integrated emission profile.  For PSR~B1937+21 the GPs occur at the extreme trailing edges of the two integrated emission components. \\citet{spb+04} found that the GPs from PSR~B1937+21 usually consist of single $\\sim 16$\\,ns shots, but also presented a pulse with structure over hundreds of nanoseconds.  \\citet{pskk04} found that at 600\\,MHz the GPs from PSR~B1937+21 are $< 10$\\% linearly polarized and highly circularly polarized.  \\psrc\\ is an isolated millisecond pulsar with a period of $P=3.054$\\,ms.  It is located 5.6\\,kpc from the Sun in the globular cluster M28 \\citep{lbm+87,har96a}.  Acceleration in the potential of its host cluster means the period derivative ($\\dot{P}$) contribution from spin down is not directly observable, and consequently other derived parameters are not known with certainty.  However, indirect evidence exists that the large $\\dot{P}$ observed is indicative of the true value --- \\psrc\\ glitches \\citep{cb04}, emits GPs, emits strong X-ray pulses \\citep{rjm+98}, and has a large second period derivative ($\\ddot{P}$) \\citep{cbl+96}.  If the observed $\\dot{P}$ is similar to the intrinsic value, then \\psrc\\ is around 30\\,Myr old and it and PSR~J1823$-$3021A are the youngest millisecond pulsars known.  This assumption also means that PSRs J1824$-$2452A and B1937+21 have inferred magnetic fields at their light cylinders ($B_{\\rm LC} \\propto P^{-2.5}\\dot{P}^{0.5}$) and spin-down luminosities ($\\dot{E} \\propto P^{-3}\\dot{P}$) that are the highest of all millisecond pulsars.  \\citet{rj01} (hereafter RJ01) discovered GPs at a center frequency ($\\nu$) of 1520\\,MHz from \\psrc\\ with energies as low as 30 times the mean pulse-energy, \\eav.  Use of a hardware filter-bank on a high dispersion measure source ($\\sim 120$\\dmunit) limited their effective time resolution to 200\\,$\\mu$s and they were unable to resolve the pulses.  Six of the eight brightest pulse detections fell within a single phase-window, which slightly trailed one of the three pulse components.  To date, it has been unclear how narrow the GPs from \\psrc\\ are, and whether they have substructure like GPs from the Crab pulsar.  In this paper we use a baseband recorder to achieve much finer time resolution than RJ01 in an attempt to uncover and study a larger population of GPs emanating from millisecond pulsars in M28.  We resolve the GPs from \\psrc, better quantify their phase and energy distributions, and perform a polarimetric analysis. ",
        "conclusions": "We have performed baseband searches for microsecond bursts of emission from sources in the globular cluster M28.  We detected 25 pulses at 20\\,cm and 2 pulses at both 2700 and 3500\\,MHz, all from \\psrc.  The emission rate of these giant pulses relative to the emission of ordinary pulses is approximately half that of PSR~B1937+21.  The pulses occur in two narrow regions of phase that coincide with the X-ray peaks, but occur on the trailing edges of the radio components. At 2700\\,MHz individual pulses were detected at phases where the corresponding component of the integrated emission was too weak to detect.  The integrated radio emission in the main emission regions has a very steep spectrum of up to $\\alpha_{\\rm spec} \\sim -4.4$.  The X-rays are emitted over a wider phase range than the giant pulses, so the two emission phenomena likely originate from similar regions in the magnetosphere but may not originate from the same physical mechanism.  The giant pulses have polarimetric properties that are completely different to those of ordinary emission.  The integrated emission has no discernable circular polarization, but the giant pulses are up to 60\\% circularly polarized.  Their linear polarization seems to have no preferred orientation, and a variety of position angles are observed. This could mean that the magnetic field in the emission regions is variable and overwhelms the dipole field of the pulsar.  Some pulses are 100\\% elliptically polarized; the average polarization fraction being 79\\%.  The two pulses observed at 3500\\,MHz have different widths and this may be interpreted as being due to resolution of the envelopes of nano-pulses.  Alternatively, the emission volumes may differ.  One of the pulses is resolved to be 20\\,ns wide.  The strongest pulse seen at 2700\\,MHz had an inferred brightness temperature at maximum of $5\\times 10^{37}$\\,K.  Although individual pulses are seen over 800\\,MHz of bandwidth, within 64\\,MHz bands at 20\\,cm the intrinsic pulse-intensity fluctuates significantly.  This indicates that the giant pulses could consist of nano-pulses that are individually narrow band."
    },
    "0608/astro-ph0608319_arXiv.txt": {
        "abstract": "{} {To determine whether ``vortex creep'' in neutron stars, the slow motion of neutron vortices with respect to pinning sites in the core or inner crust, is consistent with observations of long-period precession.}  {Using the concept of vortex drag, I discuss the precession dynamics of a star with imperfectly-pinned ({\\sl i.e.}, ``creeping'') vortices. } {The precession frequency is far too high to be consistent with observations, indicating that the standard picture of the outer core (superfluid neutrons in co-existence with type II, superconducting protons) should be reconsidered. There is a slow precession mode, but it is highly over-damped and cannot complete even a single cycle. Moreover, the vortices of the inner crust must be able to move with little dissipation with respect to the solid.}  {} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608633_arXiv.txt": {
        "abstract": "We examine the clustering properties of \\hiNS-selected galaxies through an analysis of the \\hi Parkes All-Sky Survey Catalogue (\\hicatNS) two-point correlation function.   Various sub-samples are extracted from this catalogue to study the overall clustering of \\hiNS-rich galaxies and its dependence on luminosity, \\hi gas mass and rotational velocity.  These samples cover the entire southern sky $\\delta < 0^\\circ$, containing up to 4,174 galaxies over the radial velocity range $300-12,700$ \\kmsNS.  A scale length of $r_0 = 3.45 \\pm 0.25\\, h^{-1}$Mpc and slope of $\\gamma=1.47\\pm0.08$ is obtained for the \\hiNS-rich galaxy real-space correlation function, making gas-rich galaxies among the most weakly clustered objects known.  \\hiNS-selected galaxies also exhibit weaker clustering than optically selected galaxies of comparable luminosities.  Good agreement is found between our results and those of synthetic \\hiNS-rich galaxy catalogues generated from the Millennium Run CDM simulation. Bisecting \\hicat using different parameter cuts, clustering is found to depend most strongly on rotational velocity and luminosity, while the dependency on \\hi mass is marginal.  Splitting the sample around $v_{\\rm rot} = 108$ \\kmsNS, a scale length of $r_0 =  2.86\\pm 0.46 \\,h^{-1}$ Mpc is found for galaxies with low rotational velocities compared to $r_0 =  3.96\\pm0.33\\,h^{-1}$ Mpc for the high rotational velocity sample. ",
        "introduction": "\\label{sec:intro}  The statistical analysis of galaxy clustering provides key information on the cosmological parameters of the universe and the formation and evolution of galaxies. A simple way of parametrizing galaxy clustering is though the two-point correlation function in its various redshift-space, projected and real-space forms \\citep[][]{groth1977,davis1982}.  With the advent of large-scale optical spectroscopic surveys such as the 2dF Galaxy Redshift Survey \\citep[2dFGRS, ][]{colless2001} and the Sloan Digital Sky Survey \\citep[SDSS, ][]{york2000}, the properties of the galaxy distribution are now able to be studied on cosmologically representative scales.  The large sample sizes have also enabled clustering properties to be examined as a detailed function of parameters such as optical luminosity, morphology, star formation activity and color. From these studies it has been found that the strongest clustering is exhibited by galaxies with high luminosities \\citep{norberg2001,norberg2002}, red spectral energy distributions \\citep{norberg2002,zehavi2002,zehavi2005}, and relatively passive star formation \\citep{madgwick2003}.  In this work, we focus on galaxies that are selected not on their stellar content, but on the amount of cold gas they contain.  Optical selection concentrates on the current stellar properties of galaxies, whereas \\hi selection identifies galaxies on their potential to form stars. These galaxies represent a population of more slowly evolving galaxies, still having a large fuel reservoir available for conversion into stars.   The \\hi Parkes All-Sky Survey (\\hipassNS) Catalogue \\citep[\\hicatNS;][]{meyer2004} provides such a sample over the entire southern sky.  As a blind \\hi survey, \\hipass is also not biased by extinction, providing a unique view of regions such as the Zone of Avoidance that are difficult to observe in the optical.  This paper provides a detailed analysis of the \\hicat two-point correlation function, building on earlier work by \\citet{meyer2004b} and \\citet{ryan2006} which also examine the clustering properites of \\hipass galaxies.  Here, we present a more detailed analysis, with a strong emphasis on the clustering dependencies on various galaxy parameters.  The basic properties of \\hicat are described in Section~\\ref{sec:cf_data}, with a discussion of the two-point correlation function technique following in Section~\\ref{sec:corr_func_desc}.  Sections~\\ref{sec:clustering}~and~\\ref{sec:dependencies} present the main results of this work, first covering the redshift-space, projected and real-space correlation functions, followed by an examination of the dependency of galaxy clustering on luminosity, \\hi mass, and rotational velocity.  These results are discussed in Section~\\ref{sec:discussion} with a summary given in Section~\\ref{sec:conclusions}.  A Hubble constant of \\ho = 100 \\hubble is used throughout to compare results with existing published work. ",
        "conclusions": "\\label{sec:conclusions}  Existing studies of galaxy clustering find strong dependencies on a number of parameters, highlighting an underlying trend for clustering to be strongest for more luminous galaxies \\citep{norberg2001,norberg2002}, earlier morphological types \\citep[e.g.][]{loveday1995}, and galaxies with old stellar populations \\citep{norberg2002,zehavi2002,madgwick2003}.  The low clustering measured for \\hiNS-rich galaxies is consistent with these trends, \\hiNS-rich galaxies having preferentially spiral morphologies, active star formation and blue colors.  However, the scale length of $r_0 = 3.45 \\pm 0.25\\,h^{-1}$ Mpc and slope $\\gamma=1.47 \\pm 0.08$ place \\hiNS-rich galaxies among the most weakly clustered objects known and at the extreme weak end of the observed clustering distribution.  Compared to results from the 2dFGRS, \\hiNS-rich galaxies are also more weakly clustered than optically selected galaxies of similar luminosities.  Dividing the \\hicat sample by \\hi mass around a threshold of \\mhi = $10^{9.25}\\,h^{-2}M_{\\odot}$, only a very marginal dependence of galaxy clustering strength on \\hi mass is observed.   The scale length for the low \\hi mass sample is found to be $r_0 = 3.26\\pm0.23\\,h^{-1}$Mpc and for the high mass sample $r_0 = 3.65\\pm0.30\\,h^{-1}$Mpc.   Alternatively dividing the sample on the basis of rotational velocity, a stronger dependence is seen.  The clustering scale length for galaxies $v_{\\rm rot} < 108$ \\kms is $r_0 = 2.86 \\pm 0.46\\,h^{-1}$ Mpc compared to $r_0 = 3.96 \\pm 0.33\\,h^{-1}$ Mpc for the high rotational velocity sample.   This is similar to the luminosity trend, where $r_0 = 2.90 \\pm 0.33$ for galaxies $Bj > -19.5$ and $r_0 = 3.89 \\pm 0.30$ for $Bj \\le -19.5$.  Our results are consistent with galaxy clustering being fundamentally a function of halo mass, which is well traced by stellar luminosity but poorly traced by HI gas mass.  In this scenario, \\hiNS-rich galaxies preferentially occupy lower mass halos compared to the general galaxy population, accounting for their low clustering strength.  Both environmental processes and initial conditions may lead to this effect."
    },
    "0608/astro-ph0608296_arXiv.txt": {
        "abstract": "Saturn's magnetic field is remarkably axisymmetric. Its dipole axis is inclined by less than $0.2^\\circ$ with respect to its rotation axis. Early evidence for nonaxisymmetry came from the periodicity of Saturn's kilometric radio bursts (SKR). Subsequently, percent variations of the SKR period were found to occur on timescales of years. A recent breakthrough has been the direct detection of a nonaxisymmetric component of the field that rotates with a period close to that of the SKR.  Because this component's magnitude varies only weakly with distance from Saturn, it must be supported by currents external to the planet. These currents flow along field lines that connect plasma in the equatorial region of the outer magnetosphere (the plasma disk) to the high latitude ionosphere. The plasma originates from mass lost by the planet's rings and moons.  It is tightly coupled to the magnetic field and its motion can be ascribed to the large scale interchange of flux tubes.  Heavily loaded tubes drift outward and are replaced by lightly loaded ones which drift inward. {\\it This process of rotationally driven convection is responsible for breaking the axisymmetry of Saturn's external magnetic field.} The convection pattern rotates rigidly with the angular velocity of the plasma at its source. Its rotation provides the clock that controls the periods of both the SKR bursts and the nonaxisymmetric magnetic perturbations. We distinguish two types of currents. Those that flow in the azimuthal direction in both the ionosphere and the plasma disk provide a radial force that enables the plasma to corotate with the planet's spin.  They are responsible for stretching magnetic field lines and thus exposing the polar ionosphere to the incoming solar wind. Other currents flow in the latitudinal direction in the ionosphere and the radial direction in the plasma disk. These act to transfer angular momentum from the planet's spin to the outflowing plasma. They slow down the rotation rate of the ionosphere below that of the underlying atmosphere and are the reason the clock referred to above runs slow. Moreover, these currents are the source of nonaxisymmetric magnetic perturbations whose strength varies inversely with radial distance in the planet's equatorial plane. Quantitative agreement with the magnitude of these perturbations requires a mass loss rate of order $10^4{\\, \\rm g\\, s^{-1}}$, similar to that believed to come from Saturn's E-ring. ",
        "introduction": "\\label{sec:intro}  Saturn's atmosphere exhibits strong ($\\lesssim 400 \\m \\s^{-1}$) and stable (over decadal time scales) zonal winds.  This precludes assigning a unique rotation period to its outer layers. Although the planet is fluid throughout, its deep interior must be in near solid body rotation \\citep{liu06}. Thus the internal rotation rate might be revealed by observing the nonaxisymmetric components of the planet's magnetic field.\\footnote{This technique applied to Jupiter has found a rotation period constant to within seconds over 50 years.}  However, application of this technique to Saturn has been hampered by the extreme axisymmetry of the planet's magnetic field.  \\emph{Voyager} observations of Saturn Kilometric Radiation (SKR) bursts coming from the planet's auroral regions showed a periodicity of 10h 39 min 24 $\\pm$ 7 s \\citep{des81}. The periodicity was suspected to arise from a small non-axisymmetry of the planet's internally generated magnetic field. Detections of small in situ magnetic anomalies from \\emph{Voyager} and \\emph{Pioneer 11} magnetometers were also reported \\citep{gia04}, with periods consistent with that of the SKR. \\emph{Ulysses} observations of the variability of the SKR period, by of order 1 \\% on timescales of 1 year from 1994 to 1997 \\citep{gal00}, challenged its interpretation as the rotation period of the planet's deep interior.\\footnote{A change of 1 \\% in spin period of the planet's core over such a short timescale is energetically impossible.} \\emph {Cassini} confirmed this variability, measuring an SKR period of 10h 45min 45 $\\pm$45s on approach to Saturn in 2004 \\citep{gur05}.  Most recently, magnetometer data obtained by \\emph{Cassini} showed a small ($ \\sim$ few nT) signal with period 10h 47min 6 $\\pm$40 s \\citep{gia06} that was stable during 14 months of observation. \\citet{gia06} suggested that this period might be that of Saturn's interior spin. However, the decline of the perturbation amplitude with radius is too slow to be due to a current source within Saturn, so this interpretation cannot be correct. Instead, we propose that the non-axisymmetric component of Saturn's external magnetic field is generated by rotationally driven convection in the planet's magnetosphere. The convection transports plasma from the inner magnetosphere to the magnetopause where it joins the solar wind. The plan of the paper is as follows. In \\S \\ref{sec:electro} we provide a simplified version of equations governing rotationally driven magnetospheric convection. We apply these equations in \\S \\ref{sec:mag} to estimate the nonaxisymmetric magnetic perturbations it produces. \\S \\ref{sec:clock} is devoted to a discussion of the clock that controls the perturbations' rotation rate. A short summary is given in \\S \\ref{sec:sum}.  \\subsection{Nominal parameters} \\label{subsec:parameters}  We adopt cgs units for length, mass, time, and Gaussian units for electrodynamical quanitites.  In order to focus our discussion on Saturn, we provide numerical estimates along with some of the major equations. The parameters used in these evaluations are displayed in Table \\ref{tab:nom}. \\begin{table} \\centering \\begin{tabular}{clc} \\hline Quantity& Value adopted&  Notes\\\\ \\hline $R$&$6 \\times 10^9$ cm&\\\\ $\\Omega$&$1.6 \\times 10^{-4} \\rm{~s}^{-1}$&\\\\ $B_P$&0.2 G&\\\\ $\\Sigma_p$&$10^{13} \\; \\Sigma_{p,13} \\,\\rm{~cm~s}^{-1}$& [1]\\\\ \\hline $a_i$&$ 4R \\; a_{i,4R}$&[2]\\\\ $\\Delta \\phi_i$&$\\pi \\; \\Delta\\phi_{i,\\pi}$&\\\\ $\\dot{M}$&$10^4 \\; \\dot{M}_4 \\, \\rm{~g~s}^{-1}$&[3]\\\\ \\hline $\\mu$&$3 \\times 10^{-24}$ g&\\\\ $\\gamma$&1.4&\\\\ $T$&$140 \\; T_{140} \\,$ K&[4]\\\\ $\\nu$&$10^8 \\; \\nu_8 \\rm{~cm}^2\\rm{~s}^{-1}$&[5]\\\\ $\\rho^I$&$3 \\times 10^{-13} \\; \\rho^I_{-12.5} \\rm{~g~cm}^{-3}$&[6]\\\\ $g$&$1 \\times 10^3 \\rm{~cm~s}^{-2}$&\\\\ \\hline \\end{tabular} \\caption{The numerical quantities adopted in order of magnitude calculations throughout the paper.  The meanings of the symbols are given in the text where first used. Key to numbered notes: [1] scaled between estimates for low and high latitude ionosphere \\citep{atreya84}; [2] scaled to orbit of E-ring; [3] scaled to estimates for plasma production from Saturn's E-ring \\citep{leisner06}; [4] scaled to atmospheric temperature at base of ionosphere \\citep{atreya84}; [5] scaled to eddy diffusion coefficient at homopause \\citep{atreya84}; [6] scaled between estimates for low latitude ionosphere and auroral ionosphere \\citep{atreya84}.} \\label{tab:nom} \\end{table} ",
        "conclusions": "\\label{sec:sum}  Rotationally driven convection of magnetospheric plasma breaks the axisymmetry of Saturn's external magnetic field. Field aligned currents transfer angular momentum from the planet to a tongue of outflowing plasma. This transfer slows the rate of rotation of the ionosphere relative to that of the underlying atmosphere. The currents are the source for the non-axisymmetric components of the field. The common rotation rates of these components and Saturn's kilometric radio (SKR) bursts is that of the plasma near the orbit of Enceladus, and by extension the rotation rate in the ionosphere to which this plasma is coupled. This rate tells us nothing about the rotation rate of Saturn's deep interior. Of that we remain ignorant.  Magnetic perturbations with magnitudes similar to those observed by \\emph{Cassini} are produced for ${\\dot M}\\approx 10^4 \\rm{~g~s}^{-1}$, a value similar to estimates for the rate of production of plasma from Saturn's E-ring.  Enhancement of the SKR occurs in a narrow range of longitudes where the tip of the outgoing plasma stream connects to the auroral ionosphere via field lines that are bowed outwards by currents that supply the plasma's centripetal acceleration."
    },
    "0608/astro-ph0608069_arXiv.txt": {
        "abstract": "We test the hypothesis that metal-poor globular clusters form within disk galaxies at redshifts $z>3$.  Numerical simulations demonstrate that giant gas clouds, which are cold and dense enough to produce massive star clusters, assemble naturally in hierarchical models of galaxy formation at high redshift.  Do model clusters evolve into observed globular clusters or are they disrupted before present as a result of the dynamical evolution?  To address this question, we calculate the orbits of model clusters in the time-variable gravitational potential of a Milky Way-sized galaxy, using the outputs of a cosmological $N$-body simulation.  We find that at present the orbits are isotropic in the inner 50 kpc of the Galaxy and preferentially radial at larger distances.  All clusters located outside 10 kpc from the center formed in satellite galaxies, some of which are now tidally disrupted and some of which survive as dwarf galaxies.  The spatial distribution of model clusters is spheroidal and the fit to the density profile has a power-law slope $\\gamma \\approx 2.7$, somewhat shallower than but consistent with observations of metal-poor clusters in the Galaxy.  The combination of two-body relaxation, tidal shocks, and stellar evolution drives the evolution of the cluster mass function from an initial power law to a peaked distribution, in agreement with observations.  However, not all initial conditions and not all evolution scenarios are consistent with the observed mass function of the Galactic globular clusters.  The successful models require the average cluster density, $M/R_h^3$, to be constant initially for clusters of all mass and to remain constant with time.  Synchronous formation of all clusters at a single epoch ($z=4$) and continuous formation over a span of 1.6 Gyr (between $z=9$ and $z=3$) are both consistent with the data.  For both formation scenarios, we provide online catalogs of the main physical properties of model clusters. ",
        "introduction": "\\label{sec:intro}  Observations in the past decade with the Hubble Space Telescope and major ground-based telescopes have revolutionized the field of star cluster research.  The classical division into old, massive globular clusters and young, diffuse open clusters has been challenged by the discoveries of young massive star clusters in nearby merging, interacting, and even normal star forming galaxies \\citep[e.g.,][]{holtzman_etal92, whitmore_etal93, whitmore_schweizer95, oconnell_etal95, zepf_etal99, larsen02}.  These young clusters have the range of masses and sizes similar to Galactic globulars and are believed to evolve eventually into globular clusters \\citep[e.g.,][]{ho_filippenko96, mengel_etal02, degrijs_etal04}. However, the mass functions of old and young clusters are very different.  While the mass distribution of globular clusters has a well-defined peak at $\\approx 2\\times 10^5\\, \\Msun$, the mass function of young clusters is consistent with a single power law between $10^4$ and $10^7\\, \\Msun$ \\citep{zhang_fall99, degrijs_etal03, anders_etal04}. {\\it Did old globular clusters form by a similar physical process to the young clusters but gradually evolve into the observed distribution?}  This question is vital to understanding early galaxy formation and is still a matter of debate \\citep{fall_zhang01,vesperini_etal03}. Globular clusters contain relic information about star formation processes in galaxies at high redshift and have been often used to constrain cosmological theory.  Uncovering the formation of primeval galaxies is a major driving force behind the development of future major ground and space observatories.  In contrast to the formidable difficulties of detecting the birthplaces of globular clusters at high redshift, local and interacting galaxies (M33, M51, M82, the Antennae) offer plenty of detail of the formation of young star clusters in giant molecular clouds \\citep{wilson_etal03, engargiola_etal03, keto_etal05}.  Using ultrahigh-resolution cosmological simulation, \\citet{kravtsov_gnedin05} identified similar molecular clouds within their simulated high-redshift galaxies and proposed that they harbor massive star clusters.  In this paper, we investigate whether the dynamically-evolved distribution of these model clusters can be reconciled with the current properties of globular clusters.  \\citet{kravtsov_gnedin05} find that giant clouds assemble during gas-rich mergers of progenitor galaxies, when the available gas forms a thin, cold, self-gravitating disk.  The disk develops strong spiral arms, which further fragment into separate molecular clouds located along the arms as beads on a string (see their Fig. 1).  As a result, massive clusters form in relatively massive galaxies, with the total mass $M_{\\rm halo} > 10^{9}\\ \\Msun$, beginning at redshift $z \\approx 10$.  The mass and density of the molecular clouds increase with cosmic time, but the rate of galaxy mergers declines steadily.  Therefore, the cluster formation efficiency peaks at a certain extended epoch, around $z \\approx 4$, when the Universe is only 1.5 Gyr old.  The mass function of model clusters is consistent with a power law \\begin{equation} dN/dM = N_0 \\, M^{-\\alpha}, \\label{eq:dndm} \\end{equation} where $\\alpha = 2.0 \\pm 0.1$, similar to the observations of nearby young star clusters.  According to the prescription for metal enrichment by supernovae in the \\citet{kravtsov_gnedin05} simulation, model clusters have iron abundances $[Fe/H] \\le -1$ at $z > 3$.  They correspond to the metal-poor (dominant) sub-population of the Galactic globular clusters, which we will use for the comparison.  We adopt this model to set up the initial positions, velocities, and masses for our globular clusters.  We then calculate cluster orbits using the results of a separate collisionless $N$-body simulation described in \\S~\\ref{sec:orbits_cl}.  This is necessary because the original gasdynamics simulation was stopped at $z \\approx 3$, due to limited computational resources.  By using the $N$-body simulation of a similar galactic system, but complete to $z=0$, we are able to follow the full dynamical evolution of globular clusters until the present epoch.  We use the evolving properties of all progenitor halos, from the outputs with a time resolution of $\\sim 10^8$ yr, to derive the gravitational potential in the whole computational volume at all epochs.  We convert a fraction of the dark matter mass into the analytical flattened disks, in order to model the effect of baryon cooling and star formation on the galactic potential.  We calculate the orbits of globular clusters in this potential from the time when their host galaxies accrete onto the main (most massive) galaxy. Using these orbits, we calculate the dynamical evolution of model clusters, including the effects of stellar mass loss, two-body relaxation, tidal truncation, and tidal shocks.  Since the efficiency of each of these processes depends on both cluster mass, $M$, and half-mass radius, $R_h$, we consider several evolutionary dependencies $R_h(M)$.  We also consider two possible formation scenarios, one with all clusters forming in a short interval of time around redshift $z=4$, and the other with a continuous formation of clusters between $z=9$ and $z=3$.  Note that we study only the collisionless evolution of star clusters, using the analytical formalism developed in previous studies and described in \\S~\\ref{sec:destruction}.  We do not include the recently suggested ``infant mortality'' effect, which may dissolve a large number of unbound clusters within $\\sim 100$ Myr of their formation independent of cluster mass \\citep[e.g.,][]{fall_etal05, mengel_etal05}.  Such universal reduction in numbers is therefore absorbed in our normalization of the initial GCMF.  The models calculated in this paper apply only to the initially bound clusters. ",
        "conclusions": "\\label{sec:conclusions}  We have verified the hypothesis that the properties of globular clusters formed within high-redshift disk galaxies are consistent with the observed properties of Galactic metal-poor clusters.  After calculating the orbits of model clusters in the time-variable gravitational potential of a Milky Way-sized galaxy, we find that at present the orbits are isotropic in the inner 50 kpc and preferentially radial at larger distances.  All clusters located outside 10 kpc from the center formed in satellite galaxies, some of which are now tidally disrupted.  The spatial distribution of model clusters is spheroidal, with a spherical density profile that can be fit by a power law with the slope $\\gamma \\approx 2.7$, consistent with the observations of Galactic metal-poor clusters.  Two-body relaxation dominates over other disruption processes and drives the evolution of the cluster mass function from an initial power law to the observed peaked distribution.  The evolution of the mass function in our best models is in close agreement with the models of \\citet{fall_zhang01} at late times and is somewhat slower at early times.  We find no significant variation of the mass function with radius, which is largely due to the rate of two-body relaxation being independent of cluster location.  Knowledge of realistic orbits is important for accurate calculations of the globular cluster evolution, although we find that tidal shocks play only a moderate role in cluster disruption.  Therefore, any inaccuracies in orbit calculations are not expected to significantly affect the evolution of the mass function.  Interestingly, we find that not all initial conditions and not all evolution scenarios are consistent with the observed mass function. If the half-mass radius remains constant throughout the evolution instead of decreasing with mass, the rate of two-body relaxation is suppressed and the final mass distribution is skewed towards too many low-mass clusters.  On the other hand, if the half-mass radius decreases too quickly, proportionally to the mass, then two-body relaxation is too efficient and there are no low-mass clusters left. The successful models (Sb-ii and Cb-ii) require the average cluster density, $M/R_h^3$, to be constant initially for clusters of all mass, as in the model of \\citet{kravtsov_gnedin05}, and to remain constant with time.  We have investigated two formation scenarios.  Synchronous formation of all clusters at a single epoch ($z=4$) and continuous formation over a span of 1.6 Gyr (between $z=9$ and $z=3$) both appear consistent with the observed mass function, spatial and kinematic distributions of the Galactic globular clusters.  However, clusters that formed at later epochs have more extended spatial distribution than the clusters that formed at early epochs.  This is expected for objects that trace the distribution of high-density peaks in hierarchical galaxy formation.  If additional clusters form after $z=3$, they are expected to have a shallower density profile than the clusters we study in this paper.  The profile of our clusters is consistent with, but is already somewhat shallower than the profile of the Galactic metal-poor clusters.  Therefore, we do not expect a significant fraction of metal-poor clusters to have formed after $z=3$.  Metal-rich clusters, associated with the Galactic disk, are not included in our present formation model and will be studied in future work.  In online Tables~\\ref{table:catalog1} and \\ref{table:catalog2} we provide catalogs of the physical properties of model clusters that survive to the present time, for the best models Sb-ii and Cb-ii, respectively.  These catalogs can be used to compare with other models of the dynamical evolution, to study selection effects in extragalactic systems, and to model the kinematics of globular cluster systems."
    },
    "0608/astro-ph0608543_arXiv.txt": {
        "abstract": "{We investigate the deposition of energy and momentum due to the annihilation of neutrinos ($\\nu$) and antineutrinos ($\\bar\\nu$) in the vicinity of steady, axisymmetric accretion tori around stellar-mass black holes (BHs). This process is widely considered as an energy source for driving ultrarelativistic outflows with the potential to produce gamma-ray bursts.} {We analyze the influence of general relativistic (GR) effects in combination with different neutrinosphere properties on the \\nnb-annihilation efficiency and spatial distribution of the energy deposition rate.} {Assuming axial symmetry, we numerically compute the annihilation rate 4-vector. For this purpose we construct the local neutrino distribution by ray-tracing neutrino trajectories in a Kerr space-time using null geodesics. We vary the value of the dimensionless specific angular momentum $a$ of the central BH, which provides the gravitational field in our models. We also study different shapes of the neutrinospheres, spheres, thin disks, and thick accretion tori, whose structure ranges from idealized tori to equilibrium non-selfgravitating matter distributions. Furthermore, we compute Newtonian models where the influence of the gravitational field on the annihilation process is neglected.} {Compared to Newtonian calculations, GR effects increase the total annihilation rate measured by an observer at infinity by a factor of two when the neutrinosphere is a thin disk, but the increase is only $\\approx25\\%$ for toroidal and spherical neutrinospheres. Comparing cases with similar luminosities, thin disk models yield the highest energy deposition rates by \\nnb-annihilation, and spherical neutrinospheres the lowest ones, independent of whether GR effects are included. Increasing $a$ from 0 to 1 enhances the energy deposition rate measured by an observer at infinity by roughly a factor of 2 due to the change of the inner radius of the neutrinosphere. General relativity and rotation cause important differences in the spatial distribution of the energy deposition rate by \\nnb-annihilation.} {}  \\keywords     {gamma-rays: bursts -- neutrinos -- accretion, accretion disks -- relativity -- black hole physics -- stars: neutron} ",
        "introduction": "%   It is widely believed that systems powering gamma-ray bursts (GRB) could be newborn, stellar-mass black holes (BHs) accreting matter at hyper-critical rates (up to several solar masses per second) from a surrounding accretion disk with a mass of some hundredth of a solar mass up to possibly a solar mass (see, \\eg \\citealp{2005.Piran}). These central engines may form in a `collapsar' event where the core of a massive, rotating Wolf-Rayet star collapses to a BH and the accretion of the stellar envelope may eventually lead to a GRB-supernova event with relativistic mass ejection along the rotation axis \\citep{1993.Woosley,1999.MacFadyen,2000.Aloy}. Accreting BHs may also be the remnants of mergers of two compact objects in close binaries \\citep{1989.Eichler,1993.Mochkovitch}. In the first scenario the system is embedded in the envelope of the progenitor star, and the accretion disk is fed by stellar matter yielding very long ($\\sim10-1000\\,$s) accretion time scales comparable to the collapse time scale of the progenitor star. In the second scenario viscous transport in the accretion torus sets the secular time scale of the system ($\\sim0.01-1\\,$s) formed during the merger.  The conditions in the vicinity of steady-state, hyperaccreting BHs have been analytically studied by \\cite{1993.Jaroszynski,1996.Jaroszynski,1999.Popham,2001.Narayan,2002.Matteo,2002.Kohri, 2006.Chen}, who determined the efficiency of energy loss by neutrino emission and the efficiency of energy conversion by neutrino-antineutrino (\\nnb) pair annihilation into electrons and positrons. Three-dimensional hydrodynamic simulations have explored the {\\em time-dependent} accretion in BH-torus systems, which are the remnants of neutron star--neutron star (NS+NS) and NS+BH mergers \\citep{1999.Ruffert,2004.Setiawan,2004.Lee,2005.Lee.A,2005.Lee.B}. These investigations considered a rather compact torus (typical size: $\\sim 10 - 20$ Schwarzschild radii) containing between a few hundredth of a solar mass and some $0.1\\,\\Msun$. Such torus masses result from NS+NS and NS+BH merger simulations \\citep{1999.Janka,1999.Ruffert,2002.Janka,2003.Rosswog,2003.Shibata, 2005.Shibata,2005.Oechslin}. The torus is partly opaque to neutrinos because of its high density. Typically, neutrino luminosities in excess of $10^{53}\\,$erg$\\,$s$^{-1}$ are produced. Under these conditions, the reactions $\\nu + \\bar\\nu \\rightarrow {\\rm e}^+ + {\\rm e}^- \\rightarrow \\gamma + \\gamma$ give rise to energy deposition in the close vicinity of the BH at rates ranging from several $10^{49}\\,$erg$\\,$s$^{-1}$ up to more than $10^{51}\\,$erg$\\,$s$^{-1}$ \\citep{1999.Ruffert, 1999.Janka,2004.Setiawan,2005.Setiawan}. The resulting e$^+$e$^-$-pair plasma-photon fireball may power an ultrarelativistic outflow of baryons with typical Lorentz factors of $10^2 - 10^3$, provided the baryon loading (\\ie the baryon rest mass compared to the internal energy in the outflow) remains sufficiently low (see, \\eg \\citealp{2005.Aloy} and references therein).  Several studies involving different levels of sophistication have been performed, in order to determine the amount of energy which can be released by \\nnb-annihilation near a (rotating) stellar-mass BH.  \\cite{1993.Jaroszynski,1996.Jaroszynski} investigated stationary configurations consisting of massive, dense and non-selfgravitating tori with different angular momentum distributions and different specific entropies, orbiting stellar-mass Kerr BHs. In particular, he considered the neutrino emission of isentropic tori in the external field of Kerr BHs as an approximation of what results from merger events or failed supernovae. Using the neutrino opacities of \\cite{1986.Burrows} and assuming that the neutrinos have an equilibrium (Fermi-Dirac) distribution given by the temperature and chemical potential at the neutrinosphere, \\cite{1996.Jaroszynski} determined the neutrino radiation field at a given point by following backwards null geodesics in Kerr spacetime until a point is reached in the torus where the neutrino optical depth is $\\tau_{ \\nu} \\approx 1$. He found that the energy deposition rate due to \\nnb-annihilation increases with increasing entropy and with increasing spin of the Kerr BH. He also claimed that BH-torus configurations resulting from NS+NS merger events do not provide sufficient energy to likely be sources of GRBs.  Extending earlier Newtonian calculations of \\cite{1986.Cooperstein} and \\cite{1987.Goodman}, \\cite{1999.Salmonson} analytically determined the proper energy deposition rate per unit proper time by \\nnb-annihilation near the surface of a spherical NS (the neutrinosphere was assumed to coincide with the NS surface) including general relativistic (GR) effects. They concluded that the inclusion of the GR effects of ray bending and redshift (neutrino trajectories were computed in the Schwarzschild metric) enhances the proper deposition rate per unit of proper time compared to the Newtonian values by up to a factor of 4 for a neutrinosphere located at 2.5 Schwarzschild radii relevant for a proto-NS. An enhancement factor of 30 is possible, if the radius of the neutrinosphere shrinks to 1.5 Schwarzschild radii, which happens during the collapse of a NS to a BH.  \\cite{2000.Asano} studied the influence of GR effects on the \\nnb-annihilation rate assuming two different geometries of the \\nnb-emitting regions: a spherically symmetric emission region, and a thin disk emitter surrounding a Schwarzschild BH. The spherical geometry was already studied by \\cite{1999.Salmonson}, but \\cite{2000.Asano} improved on this work by approximately taking into account that some fraction of the deposited energy might not escape from the gravitational potential well, and thus cannot contribute to power GRBs. For disk-shaped neutrinospheres they determined the energy deposition rate per unit world time for an observer located at infinity by computing the annihilation rate near the symmetry axis. The latter restriction allows for a semi-analytic treatment by making use of the axial symmetry of the neutrino source. Contrary to \\cite{1999.Salmonson}, they found that GR effects do not substantially change the energy deposition rate, neither for the spherically symmetric case nor for the disk case. This discrepancy arises, because the energy deposition rates were calculated with different constraints in both investigations. \\cite{1999.Salmonson} considered the proper energy deposition rate per unit proper time, which is enhanced by the GR effects of gravitational blueshift and bending of trajectories, for a fixed neutrino luminosity at infinity, while \\cite{2000.Asano} computed the (local) energy deposition rate per unit world time by assuming a given value of the effective temperature of the neutrinosphere and thus of the neutrinosphere luminosity. The inverse redshift factor relating the local luminosity to the luminosity at infinity explains the enhancement of the rate in case of \\cite{1999.Salmonson}. However, as argued by \\cite{2000.Asano}, the neutrinosphere luminosity, which is restricted or provided by models of the GRB engine, is the appropriate quantity to use when determining the influence of GR effects on the \\nnb-annihilation process.  Improving on their earlier work, \\cite{2001.Asano} considered a non-selfgravitating, geometrically infinitely thin accretion disk surrounding a Kerr BH. The disk was either assumed to be isothermal or to have a prescribed temperature gradient. They found that if the energy deposition (along the rotation axis) is mainly due to neutrinos from the central part of the disk near the horizon, redshift effects are dominant, and the energy deposition rate is consequently reduced compared to the case, in which GR effects were neglected during the computation. Instead, if neutrinos that are emitted at larger radii dominate the energy deposition rate, GR bending effects become important. This causes a GR enhancement of the energy deposition rate by a factor of about 2, irrespective of the spin of the Kerr BH.  \\cite{2003.Miller} considered the full three-dimensional problem of numerically computing the \\nnb-annihilation around a Kerr BH / thin accretion disk system using the full geodesic equations. They determined the energy-momentum deposition rate 4-vector per unit 4-volume in a given local observer's orthonormal frame. From this covariant quantity they derived the (coordinate-dependent, i.e. nonconvariant) energy-momentum deposition rate per proper time in Boyer-Lindquist coordinates for their observer, which is an adequate approximation to the 4-momentum per proper time at the observer's location as long as the observer is not too close to the horizon of the BH. \\cite{2003.Miller} imaged the accretion disk for a specified, but arbitrary off-axis observer. They did this by integrating geodesics along the Boyer-Lindquist energy-momentum vectors until they either hit the BH, the disk, or reach $r=50M$. They confirmed the findings of \\cite{2001.Asano} as to an only moderate GR enhancement of the energy deposition rate near the rotation axis, but they also showed that the dominant contribution to the energy deposition rate comes from near the surface of the disk, where the rate is a factor 10--20 times larger than on the rotation axis. This enhancement of the rate is independent of the spin of the BH. \\cite{2003.Miller} performed their analysis for five Kerr BHs with different angular momenta, and for each case they binned the Boyer-Lindquist time component of the energy-momentum deposition rate 4-vector reaching $r=50M$ into 20 polar angle bins. The resulting energy-momentum deposition rate per unit solid angle (as a function of polar angle) peaks along the surface of a cone centered on the rotation axis. The cone has a half-opening angle of $\\pi/4$, because the spatial components of the energy-momentum deposition rate 4-vector are tilted by about 45 degrees towards the rotation axis near the surface of the disk. The total (integrated over polar angle) energy-momentum deposition rate approximately varies linearly with the spin of the BH, the energy deposition rate of a maximally rotating Kerr BH being about twice that of a non-rotating Schwarzschild BH. This is in rough agreement with the results of \\cite{1993.Jaroszynski}.  The previous studies discussed above were concerned with a restricted set of specific, idealized cases: spherical neutrinospheres (\\citealp{1999.Salmonson,2000.Asano}), geometrically infinitely thin disks and hence neutrinospheres (\\citealp{2000.Asano,2001.Asano,2003.Miller}), and specifically designed sequences of torus models without the possibility to identify the separate influence of different components and properties of the neutrino emitting system (\\citealp{ 1993.Jaroszynski,1996.Jaroszynski}). In the following we present a more comprehensive and systematic study of energy-momentum deposition by \\nnb-annihilation. We analyze how the neutrino distribution and the resulting \\nnb-annihilation are influenced by GR effects, the geometry and the properties of the neutrinosphere, and the mass and spin of the central BH. To this end, we examine the most likely configurations occurring in compact astrophysical systems: spherical neutrinospheres present in NSs, and thin disks and thick tori surrounding stellar-mass BHs. We explicitly remark that the absolute values of the neutrino luminosity and of the energy-momentum deposition rate by \\nnb-annihilation do not play an important role for the discussion in this paper. Due to the simplicity of the neutrino source models considered here, our numbers for these quantities should not be interpreted quantitatively. They sensitively depend on the location and temperature of the neutrinospheres, which must be expected to be different in detailed hydrodynamic simulations that include some treatment of the neutrino transport. However, for the purpose of comparing the different effects, on which this paper focuses, only the relative changes of the neutrino luminosity and \\nnb-annihilation matter.  The paper is organized as follows: In Sect.\\,\\ref{sec:theory} we discuss the theoretical fundamentals for calculating the \\nnb-annihilation rate in a given Kerr spacetime, and how we construct equilibrium accretion tori surrounding Kerr BHs. In Sect.\\,\\ref{sec:numerics} we give a description of the numerical implementation of the ray-tracing algorithm used to calculate the annihilation rate. The results of our parameter study are discussed in Sect.\\,\\ref{sec:results}, and the conclusions of our work are presented in Sect.\\,\\ref{sec:conclusions}. Finally, technical details concerning the calculation of the annihilation rate are given in App.\\, \\ref{sec:appendix_maths}, and convergence tests performed on our ray-tracing algorithm are presented in App.\\,\\ref{sec:appendix_convergence}. ",
        "conclusions": "% \\label{sec:conclusions}   The main goal of this paper is to study generic properties of the process of energy deposition by neutrino-antineutrino (\\nnb) annihilation in the vicinity of systems consisting of a central stellar-mass black hole (BH) and an accretion disk or torus surrounding it. Assuming that the thermodynamic conditions in the accretion flow are such that a copious flux of \\nnb-pairs can be produced, we have performed a systematic parameter study of the influence of the mass and dimensionless angular momentum parameter of the BH, of the shape and thermal properties of the neutrinosphere, and of the importance of different general relativistic (GR) effects on the amount of energy released by \\nnb-annihilation.  On the one hand we considered idealized models having an isothermal neutrinosphere of prescribed temperature and geometry, and on the other hand non-selfgravitating, axisymmetric equilibrium tori bound to a central BH of given properties. In the latter models the neutrinospheres of neutrinos and antineutrinos do not coincide, because they are computed for opacities of $\\nu_{\\rm e}$ and $\\bar{\\nu}_{\\rm e}$ that differ due to absorptions on free neutrons or free protons, respectively (similar to the work of \\citealp{1993.Jaroszynski,1996.Jaroszynski}). Using both sets of models we numerically calculated the annihilation rate energy-momentum 4-vector. For this purpose, we constructed the local neutrino distribution by ray-tracing neutrino trajectories in a Kerr space-time using GR geodesics. Our study is a generalization of the work of \\cite{2003.Miller} to axisymmetric neutrinospheres of different shapes, which are not necessarily spatially coincident for $\\nu_{\\rm e}$ and $\\bar{\\nu}_{\\rm e}$.  We considered three different neutrinosphere geometries in our set of idealized models covering cases that may be encountered in astrophysical systems: infinitely thin disks, tori, and spheres. Infinitely thin disks are mathematical idealizations. However, they are investigated as the limiting case of very oblate tori. For this reason they have been widely used in the literature. By comparing with corresponding Newtonian models, where the influence of the gravitational field on the neutrino propagation is neglected, we have found that for an observer at infinity GR effects enhance the energy deposition rate by a factor of 2 in the non-rotating, thin disk case in agreement with \\cite{2001.Asano}. In the other two cases the influence of GR effects enhances the annihilation rate only by $\\approx 25\\%$.  Independent of whether GR effects are included, the energy deposition rate that may lead to the acceleration of relativistic outflow is largest for infinitely thin disks. For more prolate tori the energy deposition rate drops in comparison with a disk that has the same neutrino luminosity, because a sizeable fraction of the energy deposited by \\nnb-annihilation is swallowed by the BH. Spherical neutrinospheres are the least favorable geometry for a large total annihilation rate. Accretion disks encountered in gamma-ray burst (GRB) scenarios are likely to be geometrically thick, in which case the geometry of their neutrinospheres is torus-like. For such a situation neglecting GR effects in the evaluation of \\nnb-annihilation is a good approximation.  We also analyzed the influence of the dimensionless rotation parameter $a$ of the BH, and we confirm the findings of \\cite{2003.Miller} that modifying $a$ without changing the innermost disk radius (\\ie leaving the thin disk geometry unchanged) has no effect on the amount of released \\nnb-annihilation energy. The independence of the results of the BH rotation holds both for idealized models with toroidal neutrinospheres and for equilibrium accretion torus models. However, in a consistent accretion disk model, the innermost radius of the disk (close to the innermost stable circular orbit) will shrink as $a\\rightarrow 1$. We have found that a smaller innermost disk radius leads to a substantial increase of the energy deposition rate, even when the local Newtonian (\\ie GR effects disregarded) neutrino luminosity is the same. This can be understood by the higher neutrino density in the vicinity of a more compact disk and was also found by \\cite{2003.Miller}.  Depending on the mass of the accretion torus, models containing a maximally rotating BH ($a=1$) can release roughly a factor of 2 more energy (as measured by a distant observer) by \\nnb-annihilation than models involving non-rotating BHs. For a small accretion disk mass ($\\sim 0.05 \\Msun$) the energy release in case of a rotating ($a=0.8$) BH is more than 2.5 times larger than in case of a non-rotating BH. The difference is smaller for larger torus masses. Considering that low-mass accretion tori rotating around a BH with intermediate values of the dimensionless angular momentum parameter ($a \\sim 0.6-0.8$) may be typical products of mergers of compact objects (\\eg \\citealp{1997.Ruffert,2004.Setiawan}), we have demonstrated that all other GR effects are of moderate importance when calculating the amount of energy released in such systems by \\nnb-annihilation. However, we point out that the spatial distribution of the released energy exhibits very important differences between Newtonian and GR models: Close to the rotation axis and close to the bounding of the `up' region, GR effects can enhance the local energy deposition rate by a factor of $\\sim 10$. Thus, a relativistic jet driven by the energy release in a region close to the stagnation surface (which will form in the vicinity of the boundary of the `up' region) should receive a much larger energy input due to GR effects. The question whether this difference in the spatial distribution of the energy deposition is more favorable for producing ultrarelativistic jets can not be addressed by the present work, but requires time-dependent hydrodynamic jet simulations including a detailed description of \\nnb-annihilation (see below).  In our set of idealized models, a thin disk (sphere) is the most (least) favorable geometry to convert neutrino energy into ${\\rm e}^+{\\rm e}^-$-pairs by \\nnb-annihilation. Our idealized models yield efficiencies $q_{\\nu\\bar\\nu}^{{\\rm up},\\infty}$ of only a few tenths of a per cent. Larger efficiencies, of the order of several percent, have been obtained for our more realistic equilibrium models. These findings are in the ball-park of the results obtained by \\cite{1996.Jaroszynski} (for models with specific entropies $\\gsim 8$ and \\cite{2004.Setiawan}; for models with the largest $\\alpha$-viscosity).  \\cite{2003.Miller} have pointed out that the main contribution to the energy-momentum deposition rate comes from large polar angles (particularly from regions above the neutrinospheres) and not from regions near the symmetry axis. This behavior, however, results from their restriction to infinitely thin disk models, which are mathematically idealized cases. According to our studies, only such idealized models show this strong enhancement of the annihilation rate towards the equatorial plane, while this effect is absent in the more realistic equilibrium torus models. Hence, this finding of \\cite{2003.Miller} appears to be of little relevance for the amount of energy that is available to produce a GRB. In any realistic situation the accretion disk/torus will be inflated vertically, i.e. the baryon density will be highest in the equatorial plane and decrease in perpendicular direction, the density scale height of the baryon distribution being larger for large disk masses and high disk temperatures.  As demonstrated in \\cite{2005.Aloy}, who considered similar equilibrium tori as initial models for their hydrodynamic simulations of relativistic jets from BH-torus systems, time-dependent numerical simulations are needed in order to determine the feedback of \\nnb-energy deposition on the torus structure, which, in its turn, can yield a number of highly non-linear relativistic hydrodynamic effects on the outflowing jet \\citep{2006.Aloy}. This also decides about which fraction of the energy deposited by \\nnb-annihilation is finally useful for accelerating an ultrarelativistic fireball powering a GRB event. This fraction can be different from the annihilation efficiency $q_{\\nu\\bar\\nu}^{{\\rm up},\\infty}$ defined in the present work. Our equilibrium models produce \\nnb-annihilation rates $\\dot{E}_{\\nu\\bar\\nu}^{{\\rm up},\\infty} \\gsim 10^{52}$\\ergsec\\, for torus masses $m_{\\rm tor} \\gsim 0.1\\Msun$. These figures agree within a factor of 2--4 with the results of dynamical simulations (\\citealp{2005.Setiawan}) and they are about one order of magnitude higher than the ones quoted in \\cite{1996.Jaroszynski} for his models with specific entropies per baryon in the range of 8 to 10. The difference compared to \\citeauthor{1996.Jaroszynski}'s estimates is probably due to the differences in the temperature distribution on the neutrinosphere caused by our somewhat different way of constructing the accretion tori. The neutrino luminosity is very sensitive to the neutrinospheric temperature, and a $\\sim 30 \\%$ larger value of this temperature can account for a factor of more than 10 larger total \\nnb-annihilation rate. Considering a certain accretion disk mass and its corresponding neutrino luminosity and annihilation efficiency means an upper bound of the physically likely situation, in which the neutrino luminosity decreases with time as the mass of the disk decreases. However, we remark that the time evolution of the \\nnb-annihilation rate and the annihilation efficiency can be non-monotonic and does not necessarily decrease with time as demonstrated by \\cite{2005.Setiawan}.  When attempting to link our present results to GRB observations, we must therefore be aware of the restrictions stated in the previous paragraph and in particular of the fact that any reliable estimate of the fraction of the \\nnb-annihilation energy that drives an ultrarelativistic fireball and the determination of the collimation of such outflow requires hydrodynamic simulations (see also \\citealp{2006.Janka}). So we conclude with caution that the energy release by \\nnb-annihilation in some of our models (especially those having the most massive tori or large values of $a$) could be sufficient even to fuel the most distant and most powerful short GRB discovered so far (GRB060121, $E_{\\gamma} \\approx 3 \\cdot 10^{51}\\,$erg; \\citealp{2006.Postigo}) without the need of invoking either ultra-intense magnetic fields ($B>10^{16}\\,$G) or a different progenitor class.  Finally, we point out that \\citeauthor{2005.Aloy} employed a simple fit with time-independent geometry to describe the energy deposition by \\nnb-annihilation above the poles of a stellar-mass BH in their relativistic hydrodynamic jet simulations. Preferably, future numerical simulations should include a time-dependent treatment of the energy deposition by \\nnb-annihilation using the refined methods employed in this work."
    },
    "0608/astro-ph0608386_arXiv.txt": {
        "abstract": "The unexpected faintness of distant Type Ia Supernovae (SNe) has been used to argue for an accelerated expansion of the universe, with the understanding that these are thermonuclear disruptions of accreting white dwarfs.  However, the high velocity and polarized features observed in SNe Ia, and their inverse relation to luminosity, particularly for polarization, are consistent with an extreme version of the axisymmetry now seen in SN 1987A, which could be the result of double-degenerate merger-induced core-collapse.  This could be the correct paradigm for many Type I and II SNe, where Ia's are both thermonuclear {\\it and} core-collapse objects, which leave weakly magnetized, rapidly spinning ($\\sim$2 ms) pulsars.  In this paradigm, a Ia/c is produced from the merger of two degenerate cores of common envelope Wolf-Rayet stars, or of two CO white dwarfs.  Thus the same explosive mechanism that underlay 10--15 M$_\\odot$ in SN 1987A, underlies only 0--1.5 M$_\\odot$ in SNe Ia/c.  Its now visible polar blowout features produce the observed high velocity and polarized spectral features in Ia's, and its equatorial bulge is much brighter in Ia's, due to the greater fraction of $^{56}$Ni contained within it.  Such merger SNe become classified as Ia's when viewed from the merger equator, and Ic's when viewed from the poles (given sufficient matter in excess of the 1.4 M$_{\\odot}$ lost to core-collapse), where a hypernova signature and a gamma-ray burst (GRB) will be observed for lines of sight close to the merger axis.  This complication may mean that cosmology determined strictly from Ia's alone is flawed, because the local sample may be selectively biased.  Finally, all GRBs start as the short-duration, hard-spectrum variety (sGRBs), after which some are modified to the long duration/soft spectrum variety ($\\ell$GRBs) by interaction with the merger common envelope and/or previous polar ejection.  Thus the initial photon spectrum of nearly all GRBs is known. ",
        "introduction": "Type Ia supernovae (SNe Ia) have been used by at least two groups, and all without any explicit foreknowledge of their progenitors,\\footnote{In a Las Campanas 2.5-m run on SN 1987A during 1995 Feb.~my late colleague, Dr.~Jerry Kristian, mentioned to me that he considered the Ia cosmology effort ``a perversion of the science,'' because no one knew what SNe Ia really were.} to argue that the expansion of the universe is accelerating, and hence for the existence of ``dark energy,'' or a cosmological constant, $\\Lambda$ (see, e.g., Riess et al.~1998; Perlmutter et al.~1999). This has the appearance of convenience, as it helps several other lines of inquiry, including the scale size of the fluctuations of the surface of last scattering of the cosmic microwave background (CMB), and measurements of the clustering mass on large scales (see, e.g. Eisenstein et al.~2005), converge to a consistent set of parameters, generally $\\Omega_m \\sim 0.3$ and $\\Omega_\\Lambda \\sim 0.7$. However, at present SNe Ia represent the only firm, direct evidence for the existence of dark energy (see, e.g., Conley et al.~2006).  If dark energy has been convenient for cosmologists, it is certainly not so for the Standard Model, which allows no such subtle effect by 120 orders of magnitude (Weinberg 1989). Still, a lot of other recent efforts have gone into shoring up the case for cosmic acceleration (see, e.g., Clocchiatti et al.~2006), and there is no doubt that all are seeing the same effect. In addition, many have pointed out that the effect is still present even without using the width-luminosity (WL) relation, determined for local SNe Ia a decade earlier by Phillips and others (see, e.g., Phillips et al.~1999, and references therein), to correct the Ia luminosities. However, recent observations have shown that SNe Ia have high velocity and polarized line features (HVFs \\& PLFs) at the few \\% level \\citep{Le05,Maz05a}, as well as a few tenths of a \\% mean continuum polarization. In addition, core-collapse (CC) in Ia's may be {\\it necessary} to explain many of the ms pulsars (MSPs) in the globular clusters (GCs -- see $\\S$\\ref{sec:Ia/Ic?}), and the abundance of Zn \\citep{Ko06}.  In this letter I argue that many SNe Ia are caused by double-degenerate merger-induced CC (DD) which are classified as Ia/c's when viewed from the merger equator/poles, the latter occurring only when sufficient matter exists in excess of that lost to the CC to screen the Ia thermonuclear (TN) products.  This complicates the use of SNe Ia in cosmology.  I also argue that nearly all gamma-ray bursts (GRBs) are due to DD, with all starting as the short duration and hard spectrum variety (sGRBs), and only later are some modified to the long duration/soft spectrum variety ($\\ell$GRBs), by interaction with the common envelope (CE) and previous polar ejection. ",
        "conclusions": ""
    },
    "0608/astro-ph0608179_arXiv.txt": {
        "abstract": "We present preliminary results from a survey of CO emission from members of a volume-limited sample of non-cluster elliptical galaxies. Our intent is to compare the gas properties of these ellipticals to a sample of lenticulars selected using similar criteria. The data, although still sparse, suggest that the cool gas in ellipticals shows the same puzzling upper mass cutoff found in the lenticular galaxies. We find, however, significantly lower detection rates and possibly much lower H$_2$/HI mass ratios in the ellipticals. The detection rate is higher among the lower-mass galaxies, as has been found previously.  This seems puzzling given that the deeper potential wells of the larger galaxies ought to make gas retention easier, but perhaps that effect is overwhelmed by feedback from the central supermassive black hole. As we have observed $\\sim 40$ percent of our full sample, the conclusions are necessarily tentative at this time. ",
        "introduction": "There is as yet no clear, coherent and complete explanation for the origins of elliptical galaxies and their gas contents.  Since the early numerical simulations of Toomre \\& Toomre (1972) it has been evident that mergers between spirals can produce objects whose structure resembles that of an elliptical galaxy (e.g. Barnes \\& Hernquist 1992; Naab \\& Burkert 2003).  A current view is that bright ellipticals come from early, major mergers between disks (c.f. Naab et al. 1999; Naab \\& Burkert 2003).  Minor mergers - those involving mass ratios of roughly 3:1 and larger - probably formed the fainter galaxies, those with disk components and large rotational speeds.  In all cases some dissipation appears to be needed (e.g. Barnes \\& Hernquist 1996; Shioya \\& Bekki 1998; Cretton et al. 2001; Robertson et al. 2006; Naab, Jessiet \\& Burkert 2006), especially for making fainter galaxies, but how much is not clear.  To further complicate the situation, it is becoming clear that mergers between gas-poor objects have contributed importantly to the growth of many luminous ellipticals (van Dokkum 2005, Naab et al. 2006; Bell et al. 2006; Boylan-Kolchin et al. 2006) - the so-called dry or red merger hypothesis.  On the other hand, the classical monolithic formation idea has refused to die; a modified version suggested by Kormendy (1989) and Kobayashi (2004) is still attractive for explaining radial gradients in absorption line strengths (Ogando et al 2005).  There is not yet even clear agreement on the observed properties of ellipticals, especially the properties of their interstellar media.  For example, Lees et al. (1991) and Knapp \\& Rupen (1996) found molecular gas in about half of the IRAS-selected galaxies they observed, while Georgakakis et al. (2001) reported CO in just 25 percent (HI in almost 50 percent). Wiklind et al. (1995) found CO emission from 55 percent. The Georgakakis et al. work, however, was based on a sample of interacting galaxies at different stages of evolution (the stage was determined by the optical colours according to the prescription given in Georgakakis et al. 2000), while the other three studies used an IRAS criterion -- the 100 micron flux $S_{100}$ had to be $>1$ Jy. Wiklind et al. (1995) carefully assessed the morphology of each galaxy in their sample, and when they removed the merger candidates the detection rate (of presumably normal ellipticals) dropped to about 40 percent.  What is clear, and has been since Faber \\& Gallagher (1976), is that gas from evolved stars {\\it must} be returned to the interstellar medium of a galaxy, where after a Hubble time it will (if not recycled back into stars or ejected in galactic winds) account for up to $\\sim$10 percent of the visible mass (i.e. Ciotti et al. 1991, Kennicutt et al. 1994, Brighenti \\& Mathews 1997). It is also clear that new gas can be acquired by merging with like-sized galaxies, absorbing dwarf companions, or accreting pristine material from the nearby intergalactic medium.  Simple physical arguments show that the mixing of returned gas within a galaxy dominated by random motions will heat that gas to temperatures of $10^6-10^7$ K, which explains why ellipticals with $L_{\\rm B} \\ge 3\\times 10^{10}$ L$_{\\odot}$ have hot X-ray emitting halos (O'Sullivan, Forbes \\& Ponman 2001).  Continuing energy input from aging stars (supernovae, novae and winds from very massive stars) keeps at least a portion of that gas hot, and may even blow it out of the galaxy.  The details of energy transfer to the ISM, though, are complex and the general outcome correspondingly uncertain (c.f. Mathews 1990).  Simulations show that significant amounts of gas can cool near the centers (e.g. Brighenti \\& Mathews 1997; Pellegrini \\& Ciotti 1998).  Efforts to follow the evolution of the hot ISM in elliptical galaxies have led to the strong prediction that only massive galaxies will develop central reservoirs of cool gas.  It is an ongoing puzzle, therefore, that observers have reported higher HI and CO detection rates among less luminous galaxies (Lake \\& Schommer 1984, Lees et al. 1991 for HI and CO, respectively).  Lees et al. point to a possible bias against detecting broad, faint emission lines in massive objects.  Another possibility is that ISM reheating by a central AGN is more effective in massive galaxies; for example, di Matteo et al. (2005) found in simulations that when the mass of the central black hole exceeds $\\sim 10^7$ M$_{\\odot}$ that the outflows generated can turn off star formation.  A variety of mechanisms, such as ram pressure stripping, galaxy-galaxy interactions, and perhaps thermal interactions with the hot gas or sputtering and grain destruction, affect the ISM of galaxies moving within a cluster; field galaxies should be free of such complications, and are therefore the obvious starting points for understanding internal processes governing ISM evolution.  Published studies, though, have been biased not only towards cluster membership but also in favor of objects likely to contain a detectable ISM (e.g. far-infrared detection, optical dust features). Following the philosophy of our recent study of S0 galaxies (Welch \\& Sage 2003), we have defined a volume-limited sample of non-cluster elliptical galaxies with the aim to determine systematically their cool gas contents. We used integration times designed to reveal (at the 5$\\sigma$ level) just one percent of the mass expected to be returned by stellar evolution (Ciotti et al. 1991), so that even non-detections give us physically meaningful constraints.  We also take advantage of increased bandwidth compared to earlier surveys, which facilitates detection of rotationally broadened lines.  This paper reports our first results: CO observations of 18, of which we have detected emission (in either the 2-1 or 1-0 lines) from six. ",
        "conclusions": ""
    },
    "0608/astro-ph0608453_arXiv.txt": {
        "abstract": "New millimeter-wave CO and HCN observations of the host galaxies of infrared-excess Palomar Green quasi-stellar objects (PG QSOs) previously detected in CO are presented. These observations are designed to assess the validity of using the infrared luminosity to estimate star formation rates of luminous AGN by determining the relative significance  of dust-heating by young, massive stars and active galactic nuclei (AGN) in QSO hosts and {\\it IRAS} galaxies with warm, AGN-like infrared colors.  The analysis of these data is based, in part, on evidence that HCN traces high density ($>10^4$ cm$^{-3}$) molecular gas, and that the starburst-to-HCN luminosity ratio, $L_{\\rm starburst} / L'_{\\rm HCN}$, of {\\it IRAS}-detected galaxies is constant.  The new CO data provide a confirmation of prior  claims that PG QSO hosts have high infrared-to-CO luminosity ratios, $L_{\\rm IR} / L'_{\\rm CO}$, relative to {\\it IRAS} galaxies of comparable $L_{\\rm IR}$. Such high $L_{\\rm IR} / L'_{\\rm CO}$ ratios may be due to significant heating of dust by the QSO, or to an increased star formation efficiency in QSO hosts relative to the bulk of the luminous {\\it IRAS} galaxy population. The HCN data show a similar trend, with the PG QSO host IZw1 and most of the warm {\\it IRAS} galaxies having high $L_{\\rm IR} / L'_{\\rm HCN} (>1600$) relative to the cool {\\it IRAS} galaxy population for which the median $ \\left< {L_{\\rm IR} / L'_{\\rm HCN}} \\right> _{\\rm cool} \\sim 890^{+440}_{-470}$.  If the assumption is made that the infrared emission from cool {\\it IRAS} galaxies is reprocessed light from embedded star-forming regions, then high values of $L_{\\rm IR}/L'_{\\rm HCN}$ are likely the result of dust heating by the AGN. Further, if the median ratio of $L'_{\\rm HCN} / L'_{\\rm CO} \\sim 0.06$ observed for Seyfert galaxies and IZw1 is applied to the PG QSOs not detected in HCN, then the derived $L_{\\rm IR} / L'_{\\rm HCN}$ correspond to a stellar contribution to the production of $L_{\\rm IR}$ of $\\sim 7-39$\\%, and star formation rates $\\sim2-37$ M$_\\odot$ yr$^{-1}$ are derived for the QSO hosts.  The corresponding values for the warm galaxies are $\\sim10-100$\\% and $\\sim 3-220$ M$_\\odot$ yr$^{-1}$. Alternatively, if the far-infrared is adopted as the star formation component of the total infrared in cool galaxies, i.e., $\\left< L_{\\rm FIR} / L'_{\\rm HCN} \\right> _{\\rm cool} \\sim L_{\\rm starburst} / L'_{\\rm HCN}$, the stellar contributions in QSO hosts and warm galaxies to their $L_{\\rm FIR}$ are up to 35\\% and 10\\% higher, respectively, than the percentages derived for $L_{\\rm IR}$. This raises the possibility that the $L_{\\rm FIR}$ in several of the PG QSO hosts, including IZw1, could be due entirely to dust heated by young, massive stars. Finally, there is no evidence that the global HCN emission is enhanced relative to CO in galaxies hosting luminous AGN. ",
        "introduction": "Over the last two decades, considerable progress has been made in investigating the possible evolutionary connection between ultraluminous infrared galaxy mergers (ULIGs: i.e., $L_{\\rm IR} [8-1000{\\rm \\mu m}] \\geq 10^{12}$ L$_\\odot$) and quasi-stellar objects (QSOs) first proposed by Sanders et al.  (1988a,b). These studies have transformed our view of both the environments in which QSOs are triggered and fed, and the connection between star formation and active galactic nuclei (AGN) activity.  Approximately 70\\% of the optically-selected Palomar-Green (PG) QSO sample have {\\it IRAS}\\footnote{I.e., the Infrared Astronomical Satellite.} detections consistent with thermal emission from dust (Sanders et al. 1989). In addition, approximately a dozen low-redshift QSOs have been detected in CO($1\\to0$) to date (Sanders et al. 1988c; Barvainis et al. 1989; Alloin et al. 1992; Evans et al. 2001; Scoville et al. 2003), indicating substantial quantities of molecular gas out of which to fuel star formation and possibly AGN activity.  The detection of CO($1\\to0$) emission does not, in itself, imply that active star formation is occuring. Further, in galaxies known to harbor AGN, it is unclear what fraction of the infrared emission is reprocessed star light or AGN light, and thus using $L_{\\rm IR}$ to estimate star formation rates may not be justified. The present paper addresses these issues by reporting the results of an HCN($1\\to0$) survey of QSO previously detected in CO. HCN has a higher dipole moment and larger Einstein A coefficient than CO, and thus traces much denser gas ($\\gtrsim 10^4$ cm$^{-3}$). In the Galaxy, HCN emission is associated with molecular cloud cores where active star formation is occuring, and it has been detected in a large number of infrared luminous galaxies (Gao \\& Solomon 2004a). Gao \\& Solomon (2004b) have shown that, for spiral and infrared luminous galaxies spanning several orders in $L_{\\rm IR}$,  the infrared-to-HCN luminosity ratio, $L_{\\rm IR} / L'_{\\rm HCN}$, is relatively constant and has smaller deviations than $L_{\\rm IR}/L'_{\\rm CO}$ measured over the same $L_{\\rm IR}$ range.  Gao \\& Solomon conclude that HCN is a much better tracer of gas that ultimately gives rise to the young stellar population in these galaxies, and that the constancy of $L_{\\rm IR} / L'_{\\rm HCN}$ is also proof that the $L_{\\rm IR}$ is due entirely to dust-heating by young, massive stars. With this as a working assumption, the data presented here will be used to make a comparison between the $L_{\\rm IR} / L'_{\\rm HCN}$ of the QSO host sample and the {\\it IRAS} galaxy sample. Similarities between the $L_{\\rm IR} / L'_{\\rm HCN}$ ratios of the bulk of the {\\it IRAS} galaxy population and that of QSO  and {\\it IRAS} galaxies with warm, AGN-like infrared colors are expected if the $L_{\\rm IR}$ of QSO hosts and warm galaxies is due to dust-heating by young, massive stars.  This paper is divided into 6 sections. Section 2 is a summary of sample selection.  The observing procedure and data reduction are presented in \\S 3, followed by the results in \\S 4.  In \\S 5, the infrared, CO, and HCN data of the PG QSO hosts are compared with that of both cool and warm {\\it IRAS} galaxies to assess the relative contribution of star formation to the infrared emission of luminous AGN hosts.  Section 6 is a summary of the paper. Parts of this discussion have been presented elsewhere (Evans 2003, 2005a, 2005b).  Throughout this paper, we adopt an $H_0 = 75$ km s$^{-1}$ Mpc$^{-1}$ and $q_0 = 0.5$ cosmology for consistency with Evans et al. (2001, 2005c). ",
        "conclusions": "New CO and HCN observations of a sample of $z<0.17$ optically-selected PG QSO hosts with infrared excesses are presented. These data have been compared with similar data of {\\it IRAS}-detected galaxies to assess the validity of using the infrared luminosity to estimate star formation rates in luminous AGN hosts. The following conclusions are reached:  1. The new CO measurements confirm the results presented in Evans et al. (2001) -- that QSO host as a class may have low values of $L'_{\\rm CO}$ for their $L_{\\rm IR}$ relative to {\\it IRAS} galaxies, or relatively high values of $L_{\\rm IR} / L'_{\\rm CO}$. Such high $L_{\\rm IR} / L'_{\\rm CO}$ ratios may be evidence that AGN contribute significantly to the heating of dust in their host galaxies, or that they have high star formation efficiencies.  2. A comparison of the QSO hosts and {\\it IRAS} galaxies surveyed in HCN(1$\\to$0) to date show the upper distribution of $L_{\\rm IR} / L'_{\\rm HCN}$ ($>1600$) to be mostly populated by warm {\\it IRAS} galaxies and AGN hosts.  This distribution overlaps at the low end with cool {\\it IRAS} galaxies for which $\\left< {L_{\\rm IR} / L'_{\\rm HCN}} \\right> _{\\rm cool} \\sim 890^{+440}_{-470}$ L$_\\odot$ (K km s$^{-1}$ pc$^2$)$^{-1}$.  There are also 2 (out of 8) warm galaxies with $L_{\\rm IR} / L'_{\\rm HCN}$ comparable to the median $L_{\\rm IR} / L'_{\\rm HCN}$ of cool galaxies.  An enhanced $L_{\\rm IR} / L'_{\\rm HCN}$ relative to cool {\\it IRAS} galaxies thus appears to be an indication that an AGN is contributing significantly to heating the dust. However, the possibility that the dust in most of the QSO hosts and warm galaxies are heated primarily by young, massive stars cannot be entirely ruled out.  If the median ratio of $L'_{\\rm HCN} / L'_{\\rm CO} \\sim 0.06$ observed for Seyfert galaxies and the PG QSO host IZw1 is applied to the PG QSOs not detected in HCN, then the derived $L_{\\rm IR} / L'_{\\rm HCN}$ correspond to a stellar contribution to the production of $L_{\\rm IR}$ of $\\sim 7-35$\\% and star formation rates $\\sim2-37$ M$_\\odot$ yr$^{-1}$ are derived for the QSO hosts.  The corresponding values for the warm galaxies are $\\sim10-100$\\% and $\\sim 3-220$ M$_\\odot$ yr$^{-1}$.  3. The average $\\left< L_{\\rm FIR} / L_{\\rm IR} \\right>$ of the cool, warm, and PG QSO galaxies considered in this study are $0.77\\pm0.05$, $0.64\\pm0.11$, and $0.44\\pm0.10$, respectively.  Thus, the disparity between the QSO (and high-redshift AGN) hosts and the cool galaxies is diminished by comparing the $L_{\\rm FIR}$ of QSO hosts and cool galaxies. If the far-infrared is adopted as the star formation component of the total infrared in cool galaxies, i.e., $\\left< L_{\\rm FIR} / L'_{\\rm HCN} \\right> _{\\rm cool} \\sim L_{\\rm starburst} / L'_{\\rm HCN}$, then five of the nine QSO hosts, including IZw1, potentially have all of their $L_{\\rm FIR}$ generated by dust heated by massive stars. The stellar contributions to $L_{\\rm FIR}$ in QSO hosts are up to 35\\% higher than those derived for $L_{\\rm IR}$, but the warm galaxies have contributions only 10\\% higher than those derived for $L_{\\rm IR}$.  4. The PG QSO host IZw1 and the {\\it IRAS} galaxies with Seyfert optical classifications have $L'_{\\rm HCN} / L'_{\\rm CO}$ distributions comparable to HII-region like galaxies.  It is thus unlikely that the global HCN emission in galaxies harboring luminous AGN is enhanced relative to CO by radiation from the AGN.  A HCN survey of a more statistically complete sample of  {\\it IRAS} galaxies and PG QSO hosts is required to assess the robustness of many of these results. The apparent faintness of HCN emission in PG QSO hosts, and in many warm galaxies, necessitates the use of a larger millimeter-wave telescope. Such observations may be possible with the soon-to-be commissioned 50m-diameter Large Millimeter Telescope (LMT), or the 100m-diameter Green Bank Telescope (GBT). Frequencies as high as 90 GHz should be accessible with the GBT in a few years."
    },
    "0608/astro-ph0608209_arXiv.txt": {
        "abstract": "We investigate the possibility of constraining the ionization state of the Intergalactic Medium (IGM) close to the end of reionization ($z\\approx 6$) by measuring the size of the \\HII regions in high-$z$ quasars spectra. We perform a combination of multiphase smoothed particle hydrodynamics (SPH) and three-dimensional (3D) radiative transfer (RT) simulations to reliably predict the properties of typical high-$z$ quasar \\HII regions, embedded in a partly neutral IGM ($x_\\nHI=0.1$). In this work we assume a fixed configuration for the quasar lifetime and luminosity, \\ie $t_Q=10^7$ yr and $\\dot{N_\\gamma}=5.2 \\times 10^{56}$s$^{-1}$. From the analysis of mock spectra along lines of sight through the simulated QSO environment we find that the \\HII region size derived from quasar spectra is on average 30 per cent smaller than the physical one. Additional maximum likelihood analysis shows that this offset induces an overestimate of the neutral hydrogen fraction, $x_\\nHI$, by a factor $\\approx 3$. By applying the same statistical method to a sample of observed QSOs our study favors a mostly ionized ($x_\\nHI < 0.06$) universe at $z=6.1$. ",
        "introduction": "Studies of the Gunn-Peterson (Gunn \\& Peterson 1965, GP) test have now firmly established that at $z \\gtrsim 6$ the mean volume (mass) weighted neutral hydrogen fraction, $x_\\nHI$, is higher than $10^{-3}$ ($10^{-4}$) (e.g. Fan \\etal 2006a); a more precise determination is hampered by the strong sensitivity of \\lya photons resonant scattering to even tiny amounts of \\HI. Nevertheless, assessing the exact value of $x_\\nHI$ at $z\\sim6$  would represent a crucial information to discriminate between different reionization scenarios and identify the nature and distribution of ionizing sources.  The size of \\HII regions around high-$z$ luminous quasars prior to complete reionization is strongly dependent on $x_\\nHI$. Previous studies (\\eg Wyithe \\& Loeb 2004; Wyithe \\etal 2005) have tried to use this observable to improve upon the above GP constraints.  In a simple  (pre-overlap) reionization picture in which isolated quasar \\HII regions expand into an intergalactic medium (IGM) with mean neutral hydrogen fraction $x_\\nHI$, the (physical) size of the \\HII region is \\beq \\label{radius} R_d \\approx \\left(\\frac{3 \\dot{N}_\\gamma t_Q}{4 \\pi n_\\nH x_\\nHI}\\right)^{1/3} \\eeq where $\\dot{N}_\\gamma$, $t_Q$ are the ionizing photons emission rate and the quasar lifetime, respectively, and $n_\\nH$ is the hydrogen number density. This expression does not account for the recombinations occurring in the ionized gas, which can be neglected as typical quasar lifetimes ($t_Q=10^{6-8}$ yr) are much smaller than the IGM recombination time scales at these epochs. Also note that eq. \\ref{radius} does not hold during the relativistic phase of the \\HII region expansion, which lasts as long as $\\dot{N}_\\gamma t \\gtrsim (4\\pi/3)n_\\nH x_\\nHI(ct)^3$. For typical high-$z$ quasars luminosity this phase could last as much as $t_Q$. However (see \\eg White \\etal 2003), the size of the {\\it observed} \\HII region along the line of sight (LOS) to the quasar turns out to have exactly the same time evolution as deduced by assuming an infinite speed of light. In summary,  eq. \\ref{radius} can be safely used to describe the size of the quasar \\HII region observed in quasar absorption spectra and to put constraints on $x_\\nHI$.  In practice, this procedure is affected by a number of possible sources of error, which range from uncertainties in the quasar properties (lifetime, luminosity, highly biased environment) to intrinsic problems in the practical definition of the \\HII region size from the observed spectra. Despite all these problems, the available set of quasar spectra with a transmission region inside the GP trough have been exploited to constrain $x_\\nHI$. Wyithe \\& Loeb (2004) and Wyithe  \\etal (2005) performed a statistical comparison between observed and predicted radii, and derived a mean $x_\\nHI\\gtrsim 0.1$ at $z\\approx 6$. An independent study by Mesinger \\& Haiman (2004, hereafter MH04), not involving a direct measurement of the \\HII region size, found $x_\\nHI \\gtrsim 0.2$ by modelling observed properties of the \\lya and Ly$\\beta$ regions of the spectrum of the $z=6.28$ quasar SDSS J1030+0524.  These results seem to indicate a largely neutral IGM at those epochs, somewhat at odd with recent interpretations of the 3-yr WMAP data (Choudhury \\& Ferrara 2006, Gnedin \\& Fan 2006). Independent analysis of the same QSOs spectra by Yu \\& Lu (2005) and Fan \\etal (2006a) find much lower values for $x_\\nHI$. Furthermore a recent theoretical analysis by Bolton \\& Haenhelt (2006) argues that current observed spectra are consistent with a very broad range of $x_\\nHI$ values.  In this work we move a step further: by performing a detailed statistical analysis of mock quasar spectra extracted from the combined SPH and 3D RT simulations, we quantify the confidence level at which $x_\\nHI$ can be constrained by inverting eq.\\ref{radius}. ",
        "conclusions": "In this Letter we have discussed the robustness of constraints on the IGM \\HI fraction inferred from the extent of \\HII regions around high redshift quasars by means of their absorption spectra.  We have performed a combination of state-of-art multiphase SPH and 3D RT simulations to reliably predict the properties of a typical high-$z$ quasars \\HII region (\\eg extent, geometrical shape, inner opacity). We have found that RT effects do not induce strong deviations from spherical symmetry. The RT-induced dispersion in the \\HII region size along different LOS is in fact of the order of roughly 6 per cent of the mean radius which is likely smaller than the typical error induced on $R_d$ estimates by uncertainties in the quasar redshift determinations.  By deriving and analyzing mock spectra through the simulated quasar environment we have found that the \\HII region size deduced from quasar spectra, $R_f$, typically underestimates the physical one by 30 per cent.  The fact that the observed \\HII region sizes can substantially underestimate the size of the region impacted by the ionizing radiation of the quasar was already noted by Bolton \\& Haehnelt (2006), but our results gives the first quantitative estimate of the above underestimation, to which we refer to as apparent shrinking. It is worthwhile to notice here that the resolution of our simulations is not sufficient to properly resolve the IGM clumping which could result in an underestimate of RT effects. To this respect our results can be considered a lower limit for the amplitude of the apparent shrinking effect: resolving the high density clumps would result in fact in a higher inner opacity which is s expected to be the origin of a stronger apparent shrinking.  The amplitude of the apparent shrinking effect could be decreased by defining $R_f$ as the red side of the GP trough in the Ly$\\beta$ region ($R_f^{\\beta}$). A clear offset between the red side of the GP trough in \\lya and Ly$\\beta$ has been observed in the spectrum of the QSO SDSS J1030+0524 (White \\etal 2003, MH04. Recently, Bolton \\& Haehnelt (2006) have argued that the ratio $R_f^{\\beta}/R_f$ has a well-defined trend with $x_\\nHI$, which could be exploited to constrain such quantity.  However, the robustness of this result is uncertain given the small number of LOS used in the study and further investigation is needed.  In addition our analysis shows that, in the case of $x_\\nHI=0.1$, the apparent shrinking is almost completely due to resonant absorption of residual \\HI inside the ionized bubble. This contribution is highly fluctuating along different LOS, resulting in a large dispersion of the observed radii distribution, which is mainly due to the fact that $R_f$ depends on the inner \\lya opacity which is highly sensitive to $n_\\nHI$ fluctuations. This result implies that in order to properly account for these fluctuations, detailed RT calculations are required; the present work is unique in this respect, being the only study which accounts for the detailed effects of 3D RT.  The maximum likelihood analysis we have performed on a sample of 1000 mock spectra shows that the apparent shrinking effect induces an overestimate of the $x_\\nHI$ by a factor $\\approx 3$, if the IGM is only partially ionized ($x_\\nHI=0.1$). Moreover, by applying the above analysis to a sample of observed QSOs, we conclude that our study favors a mostly ionized universe at $z\\sim 6.1$ ($x_\\nHI\\lsim 0.06$), a somewhat different conclusion with respect to the previous results given in MH04 and  Wyithe \\& Loeb (2004) and Wyithe  \\etal (2005), which found a lower limit for $x_\\nHI$ of 0.2 and 0.1 respectively. Nevertheless, Bolton \\& Haenhelt (2006) posted as an outcome of their analysis that the sizes of the \\HII regions in the higher redshift quasar spectra are consistent with a significantly neutral surrounding IGM, as well as with a highly ionized IGM. Furthermore, as mentioned in the Introduction, independent studies by Yu \\& Lu (2005) and Fan \\etal (2006a) give evidence for much lower values of $x_\\nHI$. It is worth noting that our measurement agrees with the independent determination by Fan \\etal (2006) who found $x_\\nHI \\simeq 1.3\\times 10^{-3}$, based on an analysis of the HII region size evolution.  Uncertainties remain due to the fact that the range of acceptable $x_\\nHI$ values associated to a given sample of measured radii is quite large, \\eg our $R_f$ sample, having as most probable value $x_\\nHI=0.34$, still allows $x_\\nHI > 0.1,0.07$ to a 1-$\\sigma$,2-$\\sigma$ confidence level. This suggests that measurements of the \\HII size in quasar spectra can only provide rough constraints on $x_\\nHI$, as long as the knowledge of intrinsic properties of observed QSOs remains incomplete.  It is worth noting that quasar spectra could contain additional useful information on $x_\\nHI$. MH04 constrained $x_\\nHI$ from an analysis of the SDSS J1030+0524 absorption spectrum, which is completely independent from constraing $x_\\nHI$ using the radius extent. In this case the $x_\\nHI$ determination is based on the difference between the \\lya and Ly$\\beta$ optical depths and it is free from the uncertainties related to the apparent shrinking effect. It is interesting to notice that, differently from our results, MH04 find that the slightly smaller size of the \\HII region extent inferred from the \\lya and Ly$\\beta$ opacity is due to the presence of a damping wing absorption produced by the \\HI in the IGM unperturbed by the quasar. This difference can be due to the fact that MH04 compute the neutral hydrogen fraction at the equilibrium with the quasar ionizing flux which is diluted geometrically, without taking into account the further attenuation due to the inner opacity. As a consequence the inner opacity itself is underestimated. We found instead that the inner opacity, whose correct determination requires detailed 3D RT calculations, must be taken into account for a reliable determination of the apparent shrinking effect. A significant underestimate of the inner (resonant) opacity would require a higher $x_\\nHI$ outside the \\HII region in order to reproduce the observed difference between the apparent shrinking in the \\lya and Ly$\\beta$ flux. In any case, though, it is necessary to await for a larger sample of observed quasar and simulate a larger parameters space ($x_\\nHI$, $t_Q$, $\\dot{N}_\\gamma$) before being able to directly compare results coming from different approaches in the modelling as well as in the statistical analysis technique adopted."
    },
    "0608/astro-ph0608515_arXiv.txt": {
        "abstract": "We present the results of Monte Carlo simulation studies of the capability of the proposed NEMO telescope to detect TeV muon neutrinos from Galactic microquasars. In particular we determined the number of the detectable events from each known microquasar together with the expected atmospheric neutrino and muon background events. We also discuss the detector sensitivity to neutrino fluxes expected from microquasars, optimizing the event selection in order to reject the atmospheric background, and we show the number of events surviving the event selection. \\vspace{1pc} ",
        "introduction": "Microquasars are Galactic X-ray binary systems which exhibit relativistic jets, observed in the radio band \\cite{Chaty05}. Several authors propose microquasar jets as sites of acceleration of charged particles up to energies of about $10^{16}$ eV, and of high energy neutrino production. According to present models, neutrinos could be produced both in p$\\gamma$ \\cite{Levinson01,Distefano02} and pp interaction scenarions \\cite{Bednarek05,Aharonian05,Christiansen05,Romero05}.  The aim of this paper is to study the possibility to detect neutrinos from known microquasars with the proposed NEMO-km$^3$ telescope \\cite{nemo}. In particular, for each microquasar we simulated the number of expected events according to present models. We also simulated the expected background from atmospheric neutrinos and muons.  We applied an event selection in order to reject the atmospheric background. The event selection is a combination of different criteria, chosen optimizing the detector sensitivity to neutrino fluxes coming from each microquasar. Eventually we calculated the expected number of events surviving the selection referring to 1 year of data taking. ",
        "conclusions": "The possibility to detect TeV neutrinos from Galactic microquasars with the proposed NEMO-km$^3$ underwater \\v{C}erenkov neutrino telescope has been investigated.  A Monte Carlo was carried out to simulate the expected neutrino-induced muon fluxes produced by point-like sources and by atmospheric neutrinos. The expected atmospheric muon background was also simulated.  We computed the detector sensitivity for each microquasar, optimizing the event selection in order to reject the background. Finally, we applied the event selection and calculated the number of surviving events.  Our results show that, assuming reasonable scenarios for TeV neutrino production, the proposed NEMO telescope could identify microquasars in a few years of data taking, with a discovery potential for at least few cases above the 5$\\sigma$ level, or strongly constrain the neutrino production models and the source parameters."
    },
    "0608/astro-ph0608273_arXiv.txt": {
        "abstract": "The basic geometry of a stellar interferometer with two telescopes consists of a baseline vector and a direction to a star. Two derived vectors are the delay vector, and the projected baseline vector in the plane of the wavefronts of the stellar light. The manuscript deals with the trigonometry of projecting the baseline further outwards onto the celestial sphere. The position angle of the projected baseline is defined, measured in a plane tangential to the celestial sphere, tangent point at the position of the star. This angle represents two orthogonal directions on the sky, differential star positions which are aligned with or orthogonal to the gradient of the delay recorded in the $u-v$ plane. The North Celestial Pole is chosen as the reference direction of the projected baseline angle, adapted to the common definition of the ``parallactic'' angle. ",
        "introduction": "The paper describes a standard to define the plane that contains the baseline of a stellar interferometer and the star, and its intersection with the Celestial Sphere. This intersection is a great circle, and its section between the star and the point where the prolongated baseline meets the Celestial Sphere defines an oriented projected baseline. The incentive to use this projection onto the Celestial Sphere is that this ties the baseline to sky coordinates, which defines directions in equatorial coordinates which are independent of the observatory's location on Earth. This is (i) the introduction of a polar coordinate system's polar angle in the $u-v$ coordinate system associated with the visibility tables of the Optical Interferometry Exchange Format (OIFITS) \\citep{PaulsSPIE5491,PaulsPASP117}, detailing the orthogonal directions which are well and poorly resolved by the interferometer, and (ii) a definition of celestial longitudes in a coordinate system where the pivot point has been relocated from the North Celestial Pole (NCP) to the star.  In overview, Section \\ref{sec.vars} introduces the standard set of variables in the familiar polar coordinate systems. Section \\ref{sec.princ} defines and computes the baseline position angle---which can be done purely in equatorial or horizontal coordinates or in a hybrid way if the parallactic angle is used to bridge these. Some comments on relating distances away from the optical axis (star) in the tangent plane to the optical path delay conclude the manuscript in Section \\ref{sec.fov}. ",
        "conclusions": "To standardize the nomenclature, we propose to choose the North Celestial Pole as the ``reference'' direction (direction of position angle zero) and a handedness to define the sign of the position angles (mathematically positive if looking at the Celestial Sphere from the inside).   The set of position angles discussed here includes \\begin{itemize} \\item the direction of the zenith, the ``parallactic'' angle, \\item the direction toward the point where the baseline pinpoints the Celestial Sphere, the ``projected baseline angle,'' \\item position angles of ``secondary'' stars in differential astrometry. \\end{itemize} Computation of the angle proceeds via (\\ref{eq.tanpi1}) if the object coordinates are given in the local altitude-azimuth system, or via (\\ref{eq.tanpi}) if they are given in the equatorial system.  The mathematics involved is applicable to observatories at the Northern and the Southern hemisphere: positions $(A,a)$ or $(\\delta,h)$ are mapped onto a 2D zenithal coordinate system. The position angles play the role of the longitude (the North Celestial Pole the role of Greenwich or Aries). A distance between the object and the star (in the range from 0 to $\\pi$ along great circles through the Celestial Pole in an \\texttt{ARC} projection) might take the role of the polar distance---although an orthographic \\texttt{SIN} projection had been used for the calculations in this script. \\appendix"
    },
    "0608/astro-ph0608335_arXiv.txt": {
        "abstract": "We present here simultaneous INTEGRAL/RXTE observations of Sco X-1, and in particular a study of the hard X-ray emission of the source and its correlation with the position in the Z-track of the X-ray color-color diagram. We find that the hard X-ray (above about 30 keV) emission of Sco X-1 is dominated by a power-law component with a photon index of $\\sim 3$. The flux in the power-law component slightly decreases when the source moves in the color-color diagram in the sense of increasing inferred mass accretion rate from the horizontal branch to the normal branch/flaring branch vertex. It becomes not significantly detectable in the flaring branch, where its flux has decreased by about an order of magnitude. These results present close analogies to the behavior of GX 17+2, one of so-called Sco-like Z sources. Finally, the hard power law in the spectrum of Sco X-1 does not show any evidence of a high energy cutoff up to $100 - 200$ keV, strongly suggesting a non-thermal origin of this component. ",
        "introduction": "Hard X-ray emission in the brightest low-mass X-ray binaries (hereafter LMXBs), the so-called Z-sources, was occasionally detected in the past (see e.g.\\  Peterson \\& Jacobson 1966). These results received relatively little attention, mostly because the lack of a broad-band spectral coverage did not permit to establish whether an extra component was indeed required to fit the hard spectrum of these sources. Renewed interest in the hard X-ray emission properties of bright LMXBs was motivated by recent broad-band studies mainly performed with RXTE (2 -- 200 keV) and BeppoSAX (0.1 -- 200 keV). These have shown that most Z-sources display variable, hard power-law shaped components, dominating their spectra above $\\sim 30$~keV (see Di Salvo \\& Stella 2000 for a review).  The hard component detected in bright (otherwise soft) LMXBs can be fitted by a power law, with photon index in the range 1.9--3.3, contributing from 1\\% to 10\\% of the observed ($0.1-200$ keV) source luminosity. The presence of these components in Z sources seems sometimes to be related to the source state or its position in the X-ray color-color diagram (hereafter CD). The clearest example to date is in the BeppoSAX observation of GX~17+2, where the hard component (a power-law with photon index of $\\sim 2.7$) showed the strongest intensity in the horizontal branch (HB) of its CD (Di Salvo et al.\\ 2000). A factor of 20 decrease was observed when the source moved from the HB to the normal branch (NB), i.e.\\ from low to high (inferred) mass accretion rate. A hard tail was also detected in almost all the currently known Z sources (e.g. Di Salvo et al.\\ 2001; Iaria et al.\\ 2001; Di Salvo et al. 2002; Asai et al.\\ 1994). The fact that a similar hard component has been observed in several Z sources indicates that this is probably a common feature of these sources. However, the origin of this hard component is still poorly understood. While in most cases the hard component becomes weaker at higher accretion rates, HEXTE observations of Sco X--1 showed a hard power-law tail in 5 out of 16 observations, without any clear correlation with the position in the CD (D'Amico et al.\\ 2001). The thermal vs.\\ non-thermal nature of this component remains to be addressed, yielding important information on the production mechanism.  Sco X--1, the brightest persistent X-ray source in the sky, is also the brightest radio source among neutron star LMXBs, with a mean radio flux about 10 times higher than that of the other Z sources (e.g.\\ Fender \\& Hendry 2000). A hard X-ray power-law component has been observed in RXTE/HEXTE (20--200 keV) data of this source (D'Amico et al.\\ 2001). As already mentioned, contrary to the case of GX~17+2, in Sco X--1 the flux of this component was observed to vary without any clear correlation with the position in the CD. Interestingly, Strickman \\& Barret (2000) report that the hard X-ray emission present in OSSE data of Sco X--1 may be correlated with periods of radio flaring.  To study the hard X-ray emission in \\sco, the brightest of these sources, as well as its correlation with other source properties (such as radio emission and fast timing variability), we have performed a campaign of observations of Sco X--1 with INTEGRAL and RXTE. Part of these observations were also done simultaneously with radio VLBI observations (which will be discussed elsewhere). The INTEGRAL spectrum of \\sco\\ shows with high statistical significance the presence of a hard (power-law) component, without any clear exponential cutoff up to $\\sim 100-200$ keV. We also find clear evidence that the intensity of this component is correlated with the position of the source in the X-ray CD. ",
        "conclusions": "We report on a spectral analysis of a simultaneous $\\sim 300$~ks-long observation of \\sco\\ with INTEGRAL and RXTE. We show that the addition of a hard power-law component dominating the X-ray spectrum above $\\sim 30$ keV proves necessary for a good fit of the INTEGRAL ($20-200$ keV) spectrum. The power law is quite steep, with a photon index of about 3.1, contributing up to 12\\% of the observed $20 - 200$ keV luminosity, and does not show any evidence of a high energy cutoff up to $100-200$ keV. Similarly to what was observed in the BeppoSAX observation of the Z source GX~17+2, the presence of the hard component in \\sco\\ seems to be related to the position of the source in the Z track; in fact we observe a clear trend of the $20-200$ keV power-law flux to decrease when the inferred mass accretion rate increases (i.e.\\ from the HB to the NB/FB vertex). At the highest inferred mass accretion rate (i.e.\\ in the FB of the X-ray CD), the hard X-ray emission seems to disappear completely (the hard X-ray flux decreases by about one order of magnitude with respect to the flux measured in the HB/upper-NB). Note, however, that the behavior of the hard component in the FB may be variable (c.f.\\ D'Amico et. al.\\ 2001).  One of the most interesting results of this analysis is that the hard power law detected in \\sco\\ does not show any evidence of a high energy cutoff, with lower limits on the electron temperature that are in most cases above 100 keV (the lower limit on a thermal cutoff in the \\sco\\ averaged INTEGRAL spectrum, where the statistics is the highest, is about 200 keV). Such high temperatures are not expected in a system like \\sco, because of the strong Compton cooling due to the primary soft spectrum, where most of the energy is emitted. These results therefore indicate that, in analogy with the steep hard power-law emission detected in some soft states (e.g.\\ intermediate and very high) of systems containing a black hole candidates, which do not show any energy cutoff up to a few hundreds keV, the hard power law observed in Z sources may be of non-thermal origin.  The most probable origin of these components are therefore Comptonization in a hybrid thermal/non-thermal corona (where the Maxwellian velocity distribution of the electrons have a non-thermal high-velocity tail, e.g.\\ Poutanen \\& Coppi 1998; see Farinelli et al.\\ 2005 and D'A\\'\\i\\ et al.\\ 2006 for a successful application of this model to the case of GX~17+2 and \\sco, respectively). For instance, it is possible that a hard population of electrons is accelerated in internal shocks at the base of a jet. Otherwise the hard power law may originate in a mildly relativistic ($v/c \\ga 0.5$) bulk motion of matter close to the compact object. Therefore non-thermal Comptonization in outflows (or jets) may be the origin of these components, with flatter power laws corresponding to higher optical depths of the scattering medium and/or higher bulk electron velocities (e.g.\\ Psaltis 2001).  It has been proposed that non-thermal, high energy electrons, responsible for the hard tails observed in Z sources, might be accelerated in a jet (Di Salvo et al.\\ 2000). This is also confirmed by the observed correlation between the strength of the hard X-ray emission and the position of the source in the X-ray CD. In fact all Z sources are detected as variable radio sources with the highest radio flux associated with the HB  and the lowest with the FB (Migliari \\& Fender 2006, and references therein). In other words, both the radio emission from these objects, which is probably due to jets (Fender \\& Hendry 2000), and the hard X-ray emission are anticorrelated with the inferred mass accretion rate. In this respect, it will be very important to study the direct correlation between the presence and strength of the hard X-ray emission and the radio activity of the source. The radio (VLA) observations of \\sco, simultaneous to part of the INTEGRAL observations, is under investigation and will be discussed elsewhere."
    },
    "0608/astro-ph0608103_arXiv.txt": {
        "abstract": "We study the cosmological  constraints on active-sterile neutrino oscillations $\\nu_e \\leftrightarrow \\nu_s$ for the case when $\\nu_s$ is partially filled initially, i.e. $0 < \\delta N_s < 1$. We provide numerical analysis of the cosmological production of $^4\\!$He, in the presence of  $\\nu_e \\leftrightarrow \\nu_s$ oscillations, effective after neutrino decoupling, accounting for all known oscillations effects on cosmological nucleosynthesis. Cosmological constraints on oscillation parameters corresponding to  higher than $5\\%$ $^4\\!$He overproduction and different non-zero initial populations of the sterile state  $\\delta N_s<1$ are calculated. These generalized cosmological constraints corresponding to $\\delta N_s >0$ are relaxed in comparison with the $\\delta N_s=0$ case and the relaxation is proportional to $\\delta N_s$.    Keywords: BBN constraints, neutrino oscillations ",
        "introduction": "Atmospheric, solar and  terrestrial neutrino oscillations experiments provided evidence for neutrino oscillations. The atmospheric neutrino data analysis confirmed the $\\nu_\\mu \\leftrightarrow \\nu_\\tau$ channel as the dominant atmospheric oscillation solution. The results of the terrestrial experiment KamLAND and the results of SNO salt phase prefer flavor oscillation LMA solutions to the $\\nu_e \\leftrightarrow \\nu_s$ one for the solar neutrino problem. Active-sterile neutrino oscillations are discussed as a supplementary sub-dominant channel. Even after the confirmation of LMA solution as a dominant solution of the solar neutrino problem, the oscillation parameters are not precisely known - the completeness of LMA solution was questioned and the scope for some possible sub-dominant transitions was explored \\cite{al,kaki2,kaki1,chapu,hosmir,desan1,desan2,pulchau}. The available results of the neutrino oscillations from the solar experiments, including the later SNO ES, CC and NC data, slightly favor the existence of a small sterile oscillating sector. Thus, it is still interesting to explore the cosmological influence of sterile oscillating neutrinos.  Neutrino $\\nu_e \\leftrightarrow \\nu_s$ oscillations affect the expansion rate and the neutrino involved processes in the early Universe and in particular, influence Big Bang Nucleosynthesis (BBN) \\cite{dol2,ki5}. This allows to put  stringent constraints on neutrino $\\nu_e \\leftrightarrow \\nu_s$ oscillations  from BBN considerations. The whole LMA and partially LOW $\\nu_e \\leftrightarrow \\nu_s$ solar oscillation solutions and atmospheric $\\nu_e \\leftrightarrow \\nu_s$ solution were excluded by cosmological considerations \\cite{bado2,bado1,enq,chiki,ki1,kichi3,kichi2,dol1,di,cir,ciret,dolvil,dol5,kichi5}.  Detail reviews discussing neutrino role in cosmology and neutrino masses and mixings may be found in refs. \\cite{dol5,kichi5,stru,lesg}.  Cosmological constraints were obtained usually assuming zero sterile neutrino state population before neutrino oscillations epoch, $\\delta N_s=\\rho_{\\nu_s}/\\rho_{\\nu_e}^{eq}=0$, where $\\rho_{\\nu_e}^{eq}$ is the equilibrium electron neutrino energy density. However, sterile neutrinos production is predicted by different types of models, like GUT models, models with large extra dimensions, many-fold Universe models, mirror matter models, neutrino oscillations models, etc., and may be present at the onset of BBN epoch. Hence, the degree of population of $\\nu_s$ and its initial energy spectrum distribution depends  on the $\\nu_s$ production model, and in general $\\delta N_s \\ge 0$.   The general case of initially (before oscillations) non-zero sterile neutrino population was examined and its influence on the oscillation effects at BBN epoch was discussed in reference \\cite{ki2}. It was found that {\\it the kinetic effect of oscillations is sensitive to the initial population of $\\nu_s$}. BBN constraints on neutrino oscillations parameters in the general case of  sterile neutrino state, partially occupied initially, were discussed  in reference \\cite{ki3} for oscillations effective after electron neutrino decoupling. The  case of $3\\%$ $^4\\!$He overproduction was  considered. {\\it The presence of  non-empty  sterile neutrino $0<\\delta N_s<0.54$ before oscillations was shown to lead to strengthening of the cosmological constraints proportional to the  $\\delta N_s$ value}.  Recent studies of the systematic uncertainties of primordial $^4\\!$He abundance suggest as a more reliable value of  $^4\\!$He observational uncertainty  $\\delta Y_p/Y_p>5\\%$. In what follows we derive 5.2$\\%$ cosmological constraints, corresponding to $\\delta Y_p/Y_p>5\\%$ $^4\\!$He overproduction on $\\nu_e \\leftrightarrow \\nu_s$ oscillation parameters for the same type of non-equilibrium neutrino oscillations. We show that in contrast to $3\\%$ $^4\\!$He constraints, {\\it in the $\\delta Y_p/Y_p>5\\%$ case the non-zero initial sterile population leads to a relaxation of the constraints on oscillation parameters proportional to the value of $\\delta N_s$}, and finally to their alleviation  for  $\\delta N_s=1$.  In the next section, we discuss briefly the primordial $^4\\!$He production in the case of  non-zero  sterile population at the start of oscillations. In the last section we present the cosmological constraints corresponding to $^4\\!$He-overproduction $\\delta Y_p/Y_p>5\\%$ and different levels of initial population of the sterile state  and compare them to $\\delta N_s=0$ case. ",
        "conclusions": "We have studied BBN  constraints on neutrino $\\nu_e \\leftrightarrow \\nu_s$ oscillations for the specific case when the sterile neutrino is partially filled initially $0<\\delta N_s<1$. Non-zero  $\\delta N_s$ has two-fold effect on BBN with neutrino oscillations: a dynamical effect leading to overproduction of He-4  and a kinetic effect, leading to underproduction of He-4 in comparison with the case of $delta N_s=0$. So, depending on the interplay between these opposite effects, the cosmological constraints may be either relaxed or strengthened.  We have provided  detail numerical analysis of the BBN production of $^4\\!$He, $Y_p$, in the presence of $\\nu_e \\leftrightarrow \\nu_s$ neutrino oscillations, effective after electron neutrino decoupling. We have calculated  and discussed  cosmological constraints on oscillation parameters, corresponding to higher than $5\\%$ uncertainty of helium-4, for non-zero initial population of the sterile state $\\delta N_s<1$. It was found that the cosmological constraints on oscillation parameters relax with the increase of $\\delta N_s$. The cosmological constraints for the cases $\\delta N_s \\le 0.5$ are slightly changed in comparison to $\\delta N_s=0$ case, however, for bigger  $\\delta N_s$ the constraints are relaxed considerably and for  $\\delta N_s=1$ they are alleviated.  Resuming the results of the works discussing  $\\delta N_s$ effect on BBN: Cosmological constraints corresponding to higher than $5\\%$ uncertainty of helium-4,  relax with the increase of the  initial population of the sterile state, while the constraints corresponding to  lower than $5\\%$ uncertainty of helium-4 strengthen with $\\delta N_s$.  It is remarkable, that in case of BBN with non-equilibrium oscillations between electron and sterile neutrinos, it is possible to obtain cosmological constraints on oscillation parameters even in the case when the $^4\\!$He abundance is known with uncertainty greater than 5$\\%$ (Actually, it is possible to derive constraints on neutrino oscillation parameters for He-4 uncertainty up to 32$\\%$ in the resonant oscillations case, and up to 14$\\%$ uncertainty in the nonresonant oscillations case, as far as these are the maximal possible helium overproduction values \\cite{Ap}). The  cosmological constraints  persist  while   initially the sterile state is non equilibrium. When $\\delta N_s=1$ initially, the kinetic effect of oscillations is zero and it is not possible to obtain constraints on oscillation parameters in the discussed  BBN models with non-equilibrium electron-sterile neutrino oscillations, as well as in BBN models with equilibrium electron-sterile neutrino oscillations.   \\ack {The authors thank the referee for the useful comments and suggestions. D. Kirilova appreciates  the visiting position at ULB, Bruxelles and the Regular  Associateship  of  the Abdus Salam ICTP, Trieste. This work  was supported in part by Belgian Federal Government (Federal Public Planning Service Science Policy (PPS Science Policy)).}"
    },
    "0608/astro-ph0608429_arXiv.txt": {
        "abstract": "We present an exact three-dimensional wave solution to the shearing sheet equations of motion.  The existence of this solution argues against transient amplification as a route to turbulence in unmagnetized disks.  Moreover, because the solution covers an extensive dynamical range in wavenumber space, it is an excellent test of the dissipative properties of numerical codes. ",
        "introduction": "The combination of magnetic fields and differential rotation is a volatile one, leading to the magnetorotational instability, or MRI, under very general conditions (e.g. Balbus 2003).  Laboratory evidence for the MRI in liquid sodium has recently been claimed (Sisan et al. 2004), % and there can be little doubt that in magnetized disks the MRI will be present and a source of the enhanced angular momentum transport needed for accretion to proceed.  But not all astrophysical disks need be everywhere ionized at a level at which the field and the gas are well-coupled. Therefore the question of whether or not an unmagnetized, laminar gas in Keplerian rotation is vulnerable to turbulent breakdown continues to attract attention.  The nature of the local stability of Keplerian disks has been investigated by detailed numerical simulation (Balbus, Hawley, \\& Stone 1996; Hawley, Balbus, \\& Winters 1999).  There is no evidence of nonlinear instability.  Indeed, the notion that Keplerian disks are intrinsically turbulent would be most surprising and problematic to MHD numericists, who routinely check the validity of an MHD code by ensuring that the disk quickly returns to stability when the magnetic field is removed. Even when nonlinear perturbations are directly seeded, they do not lead to turbulence.  Numerical simulations have been criticized, however, on the grounds that they lack the necessary resolution to reveal nonlinear instabilities (e.g. Richard \\& Zahn 1999; Afshordi,  Mukhopadhyay, \\& Narayan 2005). The claim is that the effective Reynolds number of available codes is too low, and that the numerical schemes dissipate nascent nonlinear instabilities.  In some cases, a specific destabilizing mechanism is suggested. Afshordi et al. (2005) suggest, for example, that the transient growth exhibited by linear nonaxisymmetric perturbations, in the course of evolving from leading to trailing wavenumbers, might become sufficiently large to cross a threshold and initiate nonlinear turbulence.  The first systematic numerical studies of Keplerian stability are now a decade old, and numerical codes have developed and improved. Our physical understanding of disks has also deepened.  We are motivated to reexamine the problem of the stability of Keplerian disks.   In this paper, we present precise analytic solutions for three-dimensional incompressible perturbations in Keplerian disks, recovering the previously known two-dimensional solutions as special cases.  We show that in fact all such solutions are valid both linearly and nonlinearly because the nonlinear terms in the equations of motion vanish.  Thus, these solutions by themselves do not trigger nonlinear instabilities, regardless of the level of transient amplification attained.  Since the solutions exhibit both strongly leading and strongly trailing large wavenumbers, they cover a very large dynamical range.  They are, in fact, a good test of the dissipation properties of numerical codes.  Conversely, the numerical codes test the stability of these exact time-dependent solutions, which would otherwise be a difficult analytic calculation.  Recently, laboratory experiments have weighed in on the controversy. A high Reynolds number ($\\sim 10^6$) Couette flow experiment found that Keplerian-like profiles are nonlinearly stable (Ji 2006, private communication).  This finding is compatible with numerical simulations, but in conflict with an earlier experiment (Richard 2001) that claimed that such profiles were unstable.  The Ji experiment uses differentially rotating endcaps to break up Ekman boundaries (which otherwise cause unwanted interior circulation), and Doppler velocimetry techniques to verify that the rotation profiles matched their analytic counterparts.  The Reynolds stress in the Ji experiment has also been measured, and residual fluctuations in the radial and azimuthal velocities are uncorrelated: the tensor vanishes on average.  These results are compelling reasons to believe that local Keplerian flow is nonlinearly stable, and that the numerical simulations have not been prohibitively dissipative or otherwise qualitatively misrepresentative. The explicit solutions we present in this paper will provide numericists with a well-posed numerical test bed that will sharpen our understanding of local small scale behavior in Keplerian disks.  An outline of this paper is as follows.  In \\S 2 we provide an overview of the issues that have been debated.  Section 3 sets up the mathematical background, and the full analysis is given in \\S 4.  The solutions derived in \\S 4 are applied as a code test problem in \\S 5.  Finally, \\S 6 summarizes our work.  \\section {A Brief Review}  Axisymmetric waves propagating in the interior of an unmagnetized disk are stable if and only if the specific angular momentum $l$ increases with cylindrical radius $R$. This is the classical Rayleigh stability criterion, and it is generally satisfied in astrophysical disks. When stable, fluid displacements in the disk plane oscillate about their equilibrium circular orbits with a characteristic frequency $\\kappa$. The value of $\\kappa^2$ is directly proportional to the specific angular momentum gradient $l$.  In particular, \\beq\\label{eqn:epicycle} \\kappa^2 = \\frac{1}{R^3} \\frac{dl^2}{dR} = 4\\Omega^2 + \\frac{d\\Omega^2}{d\\ln R}, \\eeq where $R$ is the cylindrical radius and $\\Omega =l/R^2$ is the local disk angular velocity.  The formal transition to a negative value of $\\kappa^2$ corresponds to the onset of rotational instability.  It is to be noted that this result applies only to axisymmetric disturbances; there is at present no known rigorous nonaxisymmetric criterion.  (Hence the ongoing debate.) A radially discontinuous rotation profile, for example, would be unstable to non-axisymmetric disturbances by the Kelvin-Helmholtz instability, even if the rotation rate increased outward and was formally Rayleigh-stable.  For a Keplerian disk, $\\kappa^2=\\Omega^2$, which is a stable profile by the Rayleigh criterion.  Absent pressure forces, this is unconditionally stable: the behavior of Keplerian orbits is not an object of current debate.  Pressure forces must, therefore, be at the heart of any proposed instability mechanism.  This is already a potential difficulty for proponents of instability, since any system that is close to Keplerian must be characterized by very small radial background pressure gradients, and perturbative pressure gradients generally behave acoustically. Even {\\em global} instabilities in flows that would be {\\em locally} Rayleigh-stable revert to stable behavior as the disk approaches a Keplerian profile (see, e.g., Goldreich, Goodman, \\& Narayan 1986).  The claims to nonlinear instability discussed in the Introduction are of course meant to apply to systems which are linearly stable.  It is not unprecedented for linear waves to be subject to a finite amplitude instability (e.g., Infeld \\& Rowlands 1990), and even at arbitrarily small amplitudes, one-dimensional adiabatic sound waves show nonlinear steepening behavior.  (In fact, the shearing wave solutions to be discussed are, in their tightly-wrapped phase, subject to dissipative instabilities that may well prevent them from achieving even modest amplitudes.)  But proponents of nonlinear Keplerian instability are arguing for something far more novel than wave modification.  The claim is that well-behaved local linear disturbances that formally begin and end their existence as small amplitude waves, heavily dominated by restoring forces, in actuality become self-sustaining turbulence. If only the small nonlinear terms were properly resolved, it is argued, the simulations would reveal this behavior.  In this view, current disk codes would be not merely imprecise, but grossly misrepresentative of the most fundamental character of Keplerian flow.  Planar Couette and Poiseuille viscous flows are two classic shear profiles that exhibit a breakdown into turbulence.  Their behavior has motivated the notion that Keplerian rotation ought to do the same.  Planar Couette flow is linearly stable (in a marginal sense) but nonlinearly unstable, whereas Poiseuille flow is linearly unstable if the appropriately defined Reynolds number $Re$ exceeds a large critical value ($Re_{crit}=5772$, see e.g., Drazin \\& Reid [1981] for details).  Both flows are exquisitely sensitive to nonlinear disturbances.  Planar Poiseuille and Couette flows have boundary layers and linear dynamical responses lacking a dominant restoring force.  The presence of viscous boundary layers means that the Reynolds number is an important feature of the flow stability, and at the same time allows for {\\em linear} instability in Poiseuille flow at sufficiently high Reynolds number.  By way of contrast, local Keplerian flow has no viscous boundary layers, and its destabilizing shear is strongly stabilized by the Coriolis force, $2\\Omega >|d\\Omega/d\\ln R|$.  It is not surprising, therefore, that the response of disks is sensitive neither to the Reynolds number, nor to the amplitude of the initial disturbances.  The mechanism by which shear flow nonlinearly breaks down into turbulence seems to be fairly-well understood, at least in its essentials: nonlinear perturbations are in essence linear disturbances that feed off of a marginally stable, finite amplitude two-dimensional perturbation that is imposed on the shear flow (Zahn, Toomre, \\& Spiegel 1974; Bayly, Orszag, \\& Herbert 1988).  If no such marginally stable, finite amplitude state exists---and for Keplerian disks it does not---the mechanism breaks down. In fact, detailed numerical simulations show no tendency for Keplerian rotation to evolve toward self-sustained turbulence, even when driven vigorously well into the nonlinear regime by a sympathetic programmer. A revealing example is the case of thermally-driven convective turbulence. In the Boussinesq limit, this results in a very small but measurable angular momentum flux, with the wrong sign to sustain turbulence (Stone \\& Balbus 1996).\\footnote{A negative Reynolds stress, which cannot be due to numerical dissipation, poses difficulties for the notion that Keplerian disks are formally turbulent, but at a level too small to be detected (Lesur \\& Longaretti 2005).  Shear turbulence, whatever its magnitude, can be sustained only by a {\\em positive} Reynolds stress.} Moreover, even codes with very different small scale dissipation properties converge when presented with the same initial set of finite amplitude disturbances (Hawley et al. 1999).   Such convergence would hardly be expected if each code was anomalously dissipative in its own way.   The current work is supportive of numerical simulation and the recent Ji laboratory experiment.  The analytic perturbation solutions are in fact {\\em exact} solutions to the dynamical equations.   The nonlinear terms that are supposed to trigger turbulence instead vanish.  This is problematic for the idea that these disturbances seed turbulence.  More positively, it is an opportunity to test the accuracy of current codes. The solutions may be determined to any desired order of accuracy either by well-established WKB techniques, or by direct numerical integration of a simple ordinary differential equation.  The wave forms involve time-dependent wavenumbers of arbitrarily high magnitudes, both leading and trailing.  Despite the fact that the amplitude of the departure from Keplerian flow can be made arbitrarily high, the solution remains formally valid.  Secondary instabilities feeding off this finite amplitude wave cannot be excluded {\\em a priori,} but in fact the recent findings of Shen, Stone, \\& Gardiner (2006) suggest that such secondary instabilities make matters even more difficult.  Very tightly-wound leading wave solutions (which would evolve toward the largest amplitudes) break down in the early stages of their evolution by what seems to be a Kelvin-Helmholtz instability.  Far from becoming nonlinear, the wave dissipates before there is any appreciable growth. ",
        "conclusions": ""
    },
    "0608/astro-ph0608117_arXiv.txt": {
        "abstract": "{}{ Hydrogen and helium line spectra are crucial diagnostic features for the quantitative analysis of OB stars. Hybrid non-LTE line-formation calculations for these elements have not been discussed thoroughly so far, despite their wide use for analyses of metal line spectra. We compute synthetic spectra based on a hybrid non-LTE approach in order to test the ability of these models to reproduce high-resolution and high-S/N spectra of dwarf and giant stars and also to compare them with published grids of non-LTE (OSTAR2002) and LTE (Padova) models. } { Our approach solves the restricted non-LTE problem based on classical line-blanketed LTE model atmospheres. State-of-the-art model atoms and line-broadening theories are employed to model the H and \\ion{He}{i/ii} spectra over the entire optical~range~and~in~the~near-IR. } { A comparison with published line-blanketed non-LTE models validates the suitability of the LTE approximation for modelling the atmospheric structure of late O to early B-type dwarf and giant stars at metallicities down to (at least) 1/5\\,$\\times$\\,solar. Our hybrid non-LTE synthetic spectra simultaneously match almost all measurable hydrogen and helium lines observed in six test stars over a wide spectral range from the Balmer limit to the near-IR, except for only a few well-understood cases. A robust starting point for studies of the metal spectra is thus established. Our approach reproduces other published non-LTE calculations, but avoids inconsistencies in the modelling of the \\ion{He}{i} singlets. These have recently been discussed in the literature in the context of O-type stars and we show that they persist in the early B-types. Our approach improves on published pure LTE models -- widely applied for OB star analyses -- in many aspects: non-LTE strengthening and the use of improved line-broadening data result in significant differences in the line profiles and equivalent widths of the Balmer and helium lines. Where possible, systematic effects on the stellar parameter determination are quantified, e.g. gravities derived from the H$\\gamma$ wings may be overestimated by up to $\\sim$0.2\\,dex for dwarfs at our upper temperature boundary of 35\\,000\\,K in LTE.} {} ",
        "introduction": " ",
        "conclusions": "A hybrid non-LTE approach, such as is often employed in the literature for analyses of metal line spectra, has been \\emph{thoroughly} tested for hydrogen and helium line formation in OB stars for the first time, using recently improved model atoms. The synthetic spectra simultaneously match almost all measurable hydrogen and helium lines in the optical and (where available) also in the near-IR spectral range of six test stars covering a wide range of stellar parameters.  The comparison of state-of-the-art line-blanketed non-LTE and LTE models confirms that the atmospheric structure of OB dwarf and giant stars is described well under the assumption of LTE, but not their spectral energy distribution and also not their line spectra. For these stars in the range 20\\,000\\,K\\,$\\le$\\,$T_\\mathrm{eff}$\\,$\\le$\\,35\\,000\\,K and 3.0\\,$\\le$\\,$\\log g$\\,$\\le$\\,4.5 (far from the Eddington limit), our hybrid non-LTE approach is equivalent to full hydrostatic non-LTE computations.  It succeeds also in providing synthetic spectra that correctly reproduce the observed \\ion{He}{i} singlet lines, avoiding inconsistencies recently reported in the literature.  In contrast to this, pure LTE modelling based on the Padova grid (or equivalent computations with ATLAS9$+$SYNTHE) may give rise to considerable systematic errors in the atmospheric parameter determination ($T_\\mathrm{eff}$, $\\log g$) and to subsequent elemental abundance studies for the hotter stars in particular. The problems with this approach may be remedied at least near the lower temperature boundary by implementing proper line broadening data.  Finally, we suggest that the \\ion{He}{i} triplet be preferred for analyses over the singlet lines from the OSTAR2002 grid in order to avoid systematic uncertainties, in agreement with the findings of Najarro et al.~(\\cite{Najarroetal06}). Hydrogen and helium line profiles from the Padova grid may be useful for quantitative analyses of stars cooler than $\\sim$22\\,000\\,K. Preference should be given to the \\ion{He}{i} $\\lambda\\lambda$3867, 4121, 4437, 4713, 5016, and 5048\\,{\\AA} transitions, which are least affected by non-LTE effects."
    },
    "0608/astro-ph0608321_arXiv.txt": {
        "abstract": "We present new optical and X-ray observations to show that the X-ray source \\wga\\ previously associated with the SNR G309.2-0.6 can be unequivocally identified with the emission line star \\hdstar\\ located in the foreground open cluster NGC 5281. Images from \\chandra in the X-ray band as well as from \\magell in the narrow optical H$\\alpha$ band show a coincidence of the source positions within 0.5\\arcsec. The X-ray source appears extremely hot for an OB-star identified as of O9.7e type. \\xmm spectra show plasma temperatures of 1 keV and $>$ 8 keV with an X-ray luminosity of 6.2$\\pm0.1\\times10^{32}$ \\ergsec. The optical and X-ray properties are very reminiscent of the prototype emission line star $\\gamma$ Cas. We discuss the ramifications of this similarity with respect to very early type emission line stars as a new class of hard X-ray sources. ",
        "introduction": "X-ray emission from early type stars has many origins and much progress has been made, particularly in recent years, to understand some of the underlying mechanisms. Since the discovery with the \\eino~\\citep{harnden1979, seward1979} it has been established that O-stars and most early-type B-stars emit X-rays of some form~\\citep{pallavicini1981,chlebowski1989, berghoefer1994,cassinelli1994}. Observations with \\ros established a fairly tight relationship between the X-ray and bolometric luminosity with $L_X/L_{bol} \\sim 10^{-7}$~\\citep{berghoefer1996}. OB-stars that follow this relation are thought to generate X-rays from a distribution of shocks in a line driven wind\\citep{lucy1980,lucy1982,owocki1988}. We have a relatively fair understanding how the shocks lead to X-ray emission and for involved wind velocities one can derive X-ray emissivities from plasma temperatures of up to 25 MK. However, once it comes to processes that lead to these shocks, models differ quite substantially~\\citep{owocki1988, feldmeier1995, howk2000, feldmeier2002}.  Soon after the launch of \\chandra and \\xmm it became clear that the origins of X-rays in hot stars is much more complex and puzzling. Some early observations confirmed the picture where X-rays arise from shock fragments embedded within about 10 stellar radii of the stellar wind~\\citep{kahn2001,waldron2001,cassinelli2001,miller2002}, but others showed unusually high plasma temperatures and X-ray line properties which seemed to be more related to forms of magnetic activity ~\\citep{schulz2000,waldron2001,cohen2003,schulz2003,gagne2005,schulz2006}. Hard X-rays from hot stars have now been observed from various circumstellar environments. Enhanced X-ray emission from magnetically confined winds has been observed in young OB-stars in and near the Orion Trapezium~\\citep{schulz2003,gagne2005,schulz2006} and also seems the likely process for $\\tau$ Sco~\\citep{cohen2003, donati2006}. Candidates for X-ray enhancement through magnetic fields or colliding winds may be located in other young embedded clusters~\\citep{wolk2002, rho2004, townsleyetal2003}. Hard X-rays from colliding wind binaries have been observed in several exotic O-star binaries ~\\citep{sana2004,corcoran2005,pollock2005}. Binary interactions can also lead to powerful hard X-ray outbursts when two magnetically active stars interact as has been recently proposed for the massive spectroscopic binary in $\\theta^2$ Ori A~\\citep{schulz2006}.  A third category has recently been identified as a possible source of intrinsic hard X-rays: hot emission line stars. The prototype of such a source is $\\gamma$ Cas, which before \\chandra was known as a hard X-ray source but was considered likely to be a binary system of a B0.5e star orbited by a degenerate companion, probably a white dwarf~\\citep{kubo1998, owens1999}. \\citet{smith2004} recently reported on high resolution spectra obtained with \\chandra which indicated quite a number of peculiar X-ray properties. The spectra featured various emission line spectral components from hot collisional plasmas which range from 100 eV to 12 keV as well as cold fluorescence. In their analysis~\\citet{smith2004} conclude that the X-ray emission is likely to originate close to the Be star and its disk.  In this paper we report on X-ray observations of another early-type star with bright, hot X-ray emission, the O9e emission line star \\hdstar, located in the open cluster NGC~5281. The open cluster itself is fairly well-studied. Photometry and proper motion of NGC~5281  were investigated by \\citet{sanner2001}, who found an age for the cluster of 45$\\pm$10~Myr and a distance of 1580$\\pm$150~pc from fits to the color-magnitude diagram after statistical field star subtraction. They find a membership probability $P = 0.78$ for \\hdstar\\ based on the proper motion. \\citet{levenhagen2004} spectroscopically determined the fundamental parameters for \\hdstar , an effective temperature of $31910 \\pm 550$~K, an age of $4 \\pm 1$~Myr, and a luminosity  $log(L/L_{\\sun}) = 4.64 \\pm 0.1$ implying a distance modulus consistent with NGC 5281. Given updated determinations of O star effective temperatures which account for line-blanketing \\citep{martins2002}, this effective temperature categorizes \\hdstar\\ as an O9.7 star, but it is certainly on the borderline between O and B, with \\citet{levenhagen2006} recently interpreting their measurements as a B0Ve star. The optical spectra from ~\\citet{levenhagen2004} do not show any evidence for binarity.  The unequivocal identification of the X-ray source \\wga\\ \\citep{white1994a,white1994b} with \\hdstar\\ had previously been thwarted by the lack of sufficient spatial resolution in the X-ray band to eliminate source confusion in the center of NGC~5281 as well as the proximity on the sky of a young supernova remnant (SNR) \\gton\\ for which \\wga\\ could potentially have been a pulsar-wind-nebula (PWN). Herein we present new \\chandra X-ray and \\magell H$\\alpha$ observations with the spatial resolution to unambiguously establish the correspondence between \\wga\\ and \\hdstar\\ and discuss the deep X-ray spectrum of \\hdstar\\ taken with \\xmm that even at CCD spectral resolution reveals emission lines from highly ionized Fe indicative of a hot thermal plasma. ",
        "conclusions": "Using the excellent spatial resolution of \\chandra and optical images from \\magell we were able to identify the X-ray source \\wga\\, thought to be a possible PWN at the center of SNR 309.2-0.6, as the O9.7e massive star \\hdstar. The region itself is rather complex. The line of sight hosts an extended radio structure~\\citep{gaensler1998} at a distance of about 5 kpc which appears positionally coincident with a second possibly extended X-ray source from the \\ros All Sky Survey (RASS)~\\citep{voges1999} 1RXS J134651.1-624843. Follow-up \\asca observations detected strong emission-lines from the extended structure confirming its identification as a young ejecta-dominated SNR~\\citep{rakowski2001}. As for the bright point source \\wga , \\ros contours already suggested the possibility of an association with the foreground cluster NGC 5281 and its third brightest star \\hdstar\\ \\citep{gaensler1998}. However, its high X-ray brightness as well as a measured spectral index from the \\asca observation were equally consistent with a pulsar or synchrotron nebula interpretation \\citep{rakowski2001}. The low absorbing column density favored an origin in NGC 5281.   The better than 1\\arcsec\\ agreement between the \\chandra source position and the optical position of \\hdstar\\ dispels any doubts about the identification of \\wga\\ with \\hdstar. Both the optical and X-ray properties of \\hdstar\\ are reminiscent of $\\gamma$ Cas where a recent X-ray gratings study concluded that the Be-star itself is the likely source of X-ray emission \\citep{smith2004}. First, \\hdstar\\ has been classified as a very late O or very early B emission line star, similar in type to $\\gamma$ Cas (B0.5e). Second, the likeness between the X-ray properties of \\hdstar\\ and $\\gamma$ Cas is striking. Our spectral analysis produced a two component fit that features a 1 keV soft component and a dominant hard component of about 10 keV. Significant \\ion{Fe}{25} and \\ion{Fe}{26} line emission highlights the hot thermal origin of the spectrum in \\hdstar. Likewise $\\gamma$ Cas exhibits a very hot emission component up to $kT=12$~keV along with some levels of soft emission. In both stars the hard emission dominates, with the soft component contributing only 5\\% of the flux in \\hdstar\\ and 14\\% in $\\gamma$~Cas~\\citep{smith2004}, which to our knowledge is unique in stellar systems. Furthermore, the measured luminosities are quite similar, $L_{X}=6.2\\times10^{32}$ \\ergsec, $log(L_{X}/L_{bol})=-5.4$ for \\hdstar\\ compared to  $L_{X}=5.1\\times10^{32}$ \\ergsec \\citep{kubo1998}, $log(L_{X}/L_{bol})=-5.4$ for $\\gamma$ Cas using $L_{bol}$ from \\citet{hamann1992}. That the emission from \\hdstar\\ is analogous to $\\gamma$~Cas, seems by far the most likely possibility.  By contrast, the colliding wind scenario would have to involve another similarly massive companion with extremely strong winds, no evidence for which is seen optically in imaging or spectroscopy. Furthermore the spectral type of \\hdstar\\ is O9.7e or later. For such late type stars the energetics for colliding winds as the origin of the hard X-ray emission are even more difficult to justify. A recent study by~\\citet{schulz2006} involving the massive triple $\\theta^2$ Ori A, which contains an O-star of about similar type, showed that even for the existence a close massive companion, the wind densities as well as the involved kinetics do not favor a colliding wind scenario. \\hdstar\\ should be even less favorable as optical spectra do not indicate a massive companion.  The possibility that the origin of the high-energy emission from \\hdstar\\ is a magnetically confined wind seems a viable alternative. However, conditions that lead to a confinement scenario, as modeled for the case of $\\theta^1$ Ori C~\\citep{gagne2005}, seem a bit unrealistic here given the extremely high plasma temperatures of \\hdstar. Furthermore most magnetic activity of such kind has been observed in much younger OB stars~\\citep{schulz2003}. Quite recently it was concluded that the field in young $\\tau$ Sco, another young early type star that shows hard X-ray emission \\citep{cohen2003}, is likely of fossil origin rather than some intrinsic dynamo activity~\\citep{donati2006}. The answer thus has to be more complex.  The remaining possibility points to the quite enigmatic case of $\\gamma$ Cas. For the moment, though, with only a CCD-resolution spectrum we cannot go much beyond a discussion about the similarities between $\\gamma$ Cas and \\hdstar. However, the case may bear implications with respect to the nature of the X-rays in both sources. \\hdstar\\ is quite well identified and neither optical spectra~\\citep{levenhagen2004} nor our high spatial resolution optical imaging yet show any evidence for a companion. Even though the high resolution spectra of $\\gamma$ Cas provided additional evidence that the observed X-ray emission is likely not due to a possible white dwarf companion, that hypothesis could not be entirely ruled out because of some resemblance between the X-ray emission from some cataclysmic variables and the HETG spectra presented by~\\citet{smith2004} More recently~\\citet{smith2006} reported on another Be-star (B1e), HD 110432 (BZ Crucis), which shares similarly peculiar X-ray and optical characteristics with $\\gamma$ Cas. The case of \\hdstar\\ thus seems to be another candidate of ``$\\gamma$ Cas analogs'' with the strong notion that the hard X-ray emission emanates from a single OB emission line star."
    },
    "0608/astro-ph0608051_arXiv.txt": {
        "abstract": "We perform three-dimensional (3D) compressible MHD simulations over many dynamical times for an extended range of sonic and Alfv\\'{e}n Mach numbers and analyze the statistics of 3D density and 2D column density, which include probability distribution functions, spectra, skewness, kurtosis, She-L\\'{e}v\\^{e}que exponents, and genus. In order to establish the relation between the statistics of the observables, i.e. column densities, and the underlying 3D statistics of density, we analyze the effects of cloud boundaries. We define the parameter space for 3D measures to be recovered from column densities. In addition, we show that for subsonic turbulence the spectra of density fluctuations are consistent with $k^{-7/3}$ in the case of a strong magnetic field and $k^{-5/3}$ in the case of a weak magnetic field. For supersonic turbulence we confirm the earlier findings of the shallow spectra of density and Kolmogorov spectra of the logarithm of density. We find that the intermittencies of the density and velocity are very different. ",
        "introduction": "Magnetohydrodynamic (MHD) turbulence is a key element for understanding star formation \\citep[see][and references therein]{mckee02,maclow04}. In a new emerging paradigm, density perturbations produced by compressible MHD turbulence create molecular clouds and clumps within the clouds. This makes studies of density statistics very timely. In addition, recent advances in the understanding of the properties of compressible magnetic turbulence \\citep[see review by][and references therein]{cho05} call for observational testing. While velocity statistics are currently emerging as an essential tool for turbulence studies \\citep[see review by][and references therein]{lazarian04}, density statistics are more readily available from observations.  Variations of density have been extensively studied  for a long time \\cite[see][]{kaplan70,dickman85}. Electron density fluctuations have become widely used to gain insight into small-scale properties of interstellar turbulence \\cite[see][and reference therein]{armstrong95}. Fluctuations of column density, at the same time, became a useful source of information about turbulence at larger scales \\citep{falgarone05,padoan03}. Those studies of column densities are not limited to studies of spectra. Higher order structure functions have been employed recently \\cite[see][]{falgarone05,padoan03}. In fact, it was claimed in \\cite{padoan04}, that studies of higher order correlations could determine the Mach number of the interstellar flows. This motivates further studies of this type of statistics.  In addition, the topology of density in the interstellar medium (ISM) can provide an important insight into ISM physics. For instance, the topology of the ISM is expected to be different for models of the ISM in which most of the gas is in the hot phase with islands of other phases embedded within \\cite[see][]{mckee77} and in models where a substantial part of the volume is filled with cold and warm gas \\cite[see][]{cox74}. It was pointed out in \\cite{lazarian02}, that the topology of density can also be studied using genus statistics \\cite[see][]{gott86}. However, no detailed study has been performed as far as we know.  The issues above call for a comprehensive study of density statistics, which we undertake in this paper. We use an extended set of MHD numerical simulations, obtained for different sonic ${\\cal M}_s$ and Alfv\\'{e}nic ${\\cal M}_A$ Mach numbers, to study spectra, higher order statistics, and the genus of the density. We follow turbulence for many dynamical time steps (e.g. 25 for 128$^3$ simulations) to provide good sampling. In all cases we provide results of synthetic observations to compare the statistics of density with the statistics of the column density that are available through observations.  In this article we investigate the scaling properties of the structure functions of density and the logarithm of density for compressible MHD turbulence with different sonic and Alfv\\'{e}nic Mach numbers. In \\S\\ref{sec:models} we describe the numerical models of compressible MHD turbulence. In \\S\\ref{sec:pdfs} we discuss one-point probability distribution functions (PDFs). In \\S\\ref{sec:spectra} we study spectra of density and the logarithm of density and the anisotropy of density structures, and we show their dependence on Mach numbers. In \\S\\ref{sec:self-similarity} we perform tests of the property of self-similarity in the case of density structures. In \\S\\ref{sec:intermittency} we present the intermittency of density and the logarithm of density and compare the intermittency of density and column density. In \\S\\ref{sec:moments} we analyze the statistical moments, skewness and kurtosis, as possible measures of the ${\\cal M}_s$ of turbulence. In \\S\\ref{sec:considerations} we study the effects of cloud boundaries, resolution and time averaging. In \\S\\ref{sec:topology} we analyze density structure using genus statistics. We compare the genus lines to those of projected density. In \\S\\ref{sec:discussion} we discuss our results and their relation to the previous studies. We draw our conclusions in \\S\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Density in MHD turbulence}  We feel that it is important for studies of turbulent density to be understood in the context of broad MHD turbulence. There has long been an understanding that MHD turbulence is anisotropic \\cite[e.g.][]{shebalin83}. A very important insight into incompressible MHD turbulence by \\cite{goldreich95} (GS95) has been followed by progress in the understanding of compressible MHD turbulence \\cite[CL02; ][]{cho03,lithwick01,vestuto03}. In particular, simulations in CL02 showed that the Alfv\\'{e}nic cascade evolves on its own and the exchange of energy between Alfv\\'{e}n, slow, and fast modes is marginal.\\footnote{The expression proposed and tested in \\cite{cho03} shows that the coupling of Alfv\\'{e}nic and compressible modes is appreciable on the injection scale if the injection velocity is comparable to the {\\it total} Mach number of the turbulence, i.e. with $(V_A^2+c_s^2)^{1/2}$, where $V_A$ and $c_s$ are the Alfv\\'{e}n and sound velocities respectively. However, the coupling becomes marginal at smaller scales, as turbulence cascades and turbulent velocities get smaller.} This understanding differs from some earlier claims of a high degree of coupling of compressible and incompressible motions. Incidentally, the notion of high coupling was used to explain the fast damping of compressible MHD turbulence reported in \\cite{stone98} and in \\cite{maclow98}. In view of the current understanding of MHD turbulence we claim that the fast decay of MHD turbulence is due to the fast decay of the Alfv\\'{e}nic mode \\cite[see][]{cho03}, the effect that is present even in incompressible MHD \\cite[see][]{cho02b}.  In view of the aforementioned picture of MHD turbulence, density perturbations arise from the slow and fast modes. Alfv\\'{e}n modes, however, shear and cascade the slow modes (GS95) and therefore affect the density spectrum. Note that for supersonic turbulence the slow modes get modified and produce shocks [the spectrum $E(k)\\sim k^{-2}$]. The large-scale shocks produced compress density and produce a shallow spectrum, as we observed in the simulations above (see also BLC05).  In terms of the study of PDFs and spectra, our present work goes beyond the one in BLC05, as we provide a wider parameter space survey. However, we reach conclusions similar to those in BLC05 \\cite[see also][]{padoan04}. Interestingly enough, our results are also similar to those obtained in non-magnetized fluids by \\cite{kim05}. Indeed, we find that the spectrum of density gets shallower than the Kolmogorov one as the sonic Mach number increases. The reason for this is the formation of small high-density compressed regions that dominate the spectrum. If the logarithm of density is used, these regions do not dominate and the underlying spectrum influenced by anisotropic Alfv\\'{e}n mode shearing emerges. As for the PDFs, they are close to lognormal for sub-Alfv\\'{e}nic turbulence even if the turbulence is supersonic. For super-Alfv\\'{e}nic turbulence the PDFs can be approximated by a lognormal distribution only for low sonic Mach number.  The remarkable regularity of $\\log \\rho$ should have a physical explanation. A possible one is related to the multiplicative symmetry with respect to density in the equations for isothermal hydrodynamics \\cite[see][]{passot98}. This means that if a stochastic process disturbs the density, it results in perturbations of density being multiplied rather than added together. Consequently, the distribution of density for a Gaussian driving of turbulence tends to be lognormal, which is consistent with our PDF measurements. Magnetic forces should affect the multiplicative symmetry above. However, they do not affect the compressions of gas parallel to the magnetic field. Compressions in magnetically dominated fluid will be of the highest intensity and therefore most important. They are sheared by Alfv\\'{e}nic modes, as their own evolution will be slower than that imposed by Alfv\\'{e}nic cascade \\cite[see theoretical considerations in GS95,][]{cho03}. We note that shearing does not affect the PDFs, but it does affect the spectra and the anisotropy of the turbulence. This explains the close relationship between the properties of velocity in our simulations and those of density.   \\subsection{Properties of 3D density}  We confirm the existence for density of a rather mysterious property of turbulence called ``extended self-similarity''. While our spectra have a limited inertial range, the moments of density normalized by the third moment show nice power laws. Interestingly enough, we find that the extended self-similarity is not true if one uses the normalization over the third moment of velocity. While we still may wonder about the origin of this phenomenon, the extended self-similarity is an indispensable tool for studyint turbulence intermittency.  Our study of the intermittency of density and the logarithm of density reveals the dependence of the She-L\\'{e}v\\^{e}que exponents on both sonic and Alfv\\'{e}n Mach numbers. Moreover, we note a rather weak dependence on the intermittency when ${\\cal M}_s < 1$. For supersonic turbulence the dependency is much stronger. The intermittency changes with the strength of the external magnetic field may frequently be more important. We observe less intermittent structures when $B_\\mathrm{ext}$ is strong (see Figure \\ref{fig:dens_expons}).  The question of anisotropy of structures is in a natural way related to the presence of magnetic field. When no magnetic field is present, there is a lack of any force that could organize these structures. On the other hand, the introduction of the external magnetic field to the turbulent system induces the Alfv\\'{e}nic mode, which is able to imprint its structure on random densities. In MHD turbulence, one direction is chosen, the direction of $B_\\mathrm{ext}$, so the anisotropy should be observed. In our sub-Alfv\\'{e}nic simulations the anisotropy of density is difficult to distinguish due to large dispersion of the structure functions. The situation improves in the case of models with a weaker magnetic field. Here the results are in accordance with the GS95 model (i.e. $l_\\parallel \\sim l_\\perp^{2/3}$). The anisotropy of the logarithm of density and velocity structures show similar GS95 behavior in all models with observed dependence on the scale and independence of ${\\cal M}_s$. However, the structures of $\\log \\rho$ and velocity for sub-Alfv\\'{e}nic models seem to be less anisotropic at the small scales.  Some basic properties of the density distribution in turbulence have been already studied. However, most papers were focused mainly on the variance of distribution. We analyzed here also two higher order statistical moments: skewness and kurtosis. We observe a strong dependence of these moments on ${\\cal M}_s$. For instance, skewness grows with the sonic Mach number for mid- and supersonic turbulence. For highly subsonic turbulence we observe the growth of skewness but with a decreasing ${\\cal M}_s$. Kurtosis in this case is zero and starts to grow rapidly when the value of ${\\cal M}_s$ exceeds about 0.7. Combining these two measures can allow one to determine the value of ${\\cal M}_s$ in turbulence. We do not observe any significant dependency of these moments on ${\\cal M}_A$.  Genus statistics reveal the changes of the structure of density with sonic Mach number. For lower values of ${\\cal M}_s$ density has a normal distribution, signifying roughly the same number of separate rarefactions and condensations. For higher sonic Mach numbers, overdense structures start to dominate over the underdense ones. This changes the structure of density significantly. For low sonic Mach numbers, the structure of the density is spongelike. In the supersonic turbulence, the shocks produce separated high-density 1D and 2D structures embedded in the medium of much lower density. In this sense, the genus statistics confirm the topology with the intermittency.   \\subsection{What can be found from observations?}  Many surveys provide extensive information on column density statistics. Those include H{\\footnotesize I} surveys, e.g. the GALFA (\\url{http://www.naic.edu/alfa/galfa/}) survey of galactic H{\\footnotesize I} distribution. The main goal of another observational project, the Wisconsin H-Alpha Mapper (WHAM), is to provide a survey of H$\\alpha$ emission from the ISM over the entire northern sky \\citep{haffner03,hill05}. The observations provide the dispersion and emission measures, which can be interpreted as the column density and the integral of the square of the electron density along the entire LOS. The survey supported by the numerical studies helps in better understanding the density and distribution of diffuse plasma in a major component of our Galaxy, the warm ionized medium. Combining data from different surveys and using the techniques in this paper, it is possible to get valuable insight into processes taking place in different ISM phases.  If the observable column density distributions are concerned, our study confirms the earlier conclusion in \\cite{ostriker03} that the PDFs are affected by the LOS distribution of density, and therefore, the relation between the column density PDFs and the underlying 3D PDFs is rather involved. On the contrary, the relation between the 3D and 2D density spectra is very straightforward, as was shown earlier by many authors \\cite[see][]{stutzki99,lazarian95}. This makes the observable column density spectrum a good measure of the underlying 3D density spectrum if the value of the spectral index is concerned. However, the anisotropies of the density distribution that can be revealed for low sonic Mach number by density itself and for high Mach number by the distribution of $\\log\\rho$  are affected by the LOS integration that obtaining the column densities involves. The problem here is that the GS95 anisotropies of the magnetic field and velocities that induce anisotropies of the density are defined in the system of reference aligned with the {\\it local} magnetic field \\cite[see discussion in ][]{clv03}. LOS integration averages the local field directions, and therefore, only the anisotropies have to be calculated with respect to the mean magnetic field. This both decreases the degree of anisotropy and erases the scale dependence of the anisotropy. The residual anisotropy is the anisotropy of the turbulence at the injection scale. Therefore, for super-Alfv\\'{e}nic turbulence for which magnetic fields are dynamically unimportant at the injection scale, the anisotropy of the column density statistics vanishes altogether.  \\cite{padoan03} studied maps of $^{13}$CO intensity of molecular clouds in Taurus and Perseus in terms of structure functions. The integrated intensity maps can be translated into column densities by simple multiplication by a factor \\citep[see][]{ossenkopf02}. \\cite{padoan03} show that the scaling exponents of the structure functions, normalized to the third order, follow the velocity scaling of the supersonic turbulence. This suggests that the density scaling can be described by a similar hierarchical model, like the one proposed by \\cite{she94}. Our study shows that the intermittencies of densities and velocities are intrinsically different and one should exercise care in relating the relevant She-L\\'{e}v\\^{e}que exponents. Although the intermittencies can reflect the sonic Mach number, this dependence that we observe is weak for subsonic models and the changes are below the errors (see Figure \\ref{fig:col_dens_scal}). The low-resolution models showed the weak dependence of the intermittency of density on the sonic Mach number. However, the intermittency is very sensitive to the resolution, which was stressed in \\S\\ref{sec:resolution}. The growth of error bars with the order of exponent $p$ is due to the amplification of any irregularities in the structure function by higher order moments. Another difficulty is the change of the intermittency while we switch from density to column density, as seen by the comparison of the scaling exponents in Figures \\ref{fig:dens_expons} and \\ref{fig:dens_int_expon}. We note, that the intermittency of the column densities in general is lower than the intermittency of density. In such situations we can only speak about some rough predictions taken from the observational data.  In addition, we see that the effects of cloud boundaries change the intermittency itself depending on ${\\cal M}_s$. For instance, the highly subsonic and supersonic turbulence reveal more intermittent structure after including the boundary. The effect is especially strong for turbulence with high ${\\cal M}_s$ and grows with its value. The strength of the external magnetic field does not influence significantly the effects of the cloud boundary.  We also present a number of different measurements that can be helpful in the determination of the properties of density structures in real cosmic objects, like molecular clouds. We compared the statistical measures applied for the 3D density and corresponding 2D column densities. By such comparison, we were able to regain information about the spatial structure from the synthetic maps of the projected density. For instance, we have shown, that the PDFs of density and column density can be related in particular cases.  We find that the statistical moments skewness and kurtosis, which are easy to obtain, can be very useful. They can be employed to determine the sonic Mach number for supersonic turbulence. The subsonic regime of turbulence should be treated more carefully, because the cloud boundary effects might be very important in this case, as was shown in \\S\\ref{sec:boundary}. Boundary effects influence the PDFs and the scaling of density structure (see Figures \\ref{fig:dens_sph_pdfs} and \\ref{fig:dens_sph_expon}, respectively). In the case of PDFs, however, some information can still be regained, because the boundary of the cloud mostly influences the low-density tail of the distribution. Three-dimensional density spectra can be recovered successfully from observations of column densities.  Another promising technique for the topology of density and column density study is genus statistics. Although it is a completely different technique than one- or two-point statistics, we have shown that it can both help us with understanding the topology of density structures in turbulence and relate the results obtained for density and column density (see \\S\\ref{sec:topology}).  Some measures presented in the article, like statistical moments and spectra, should be easily obtained from the molecular maps. Such observation of quantities can be directly compared with our numerical results, allowing quick determination of the conditions in investigated maps. Other measurements, like scaling exponents and genus statistics require more effort to calculate, but can provide more information related to the density structure. Thus, the techniques of the study of turbulence using 2D maps of observed column densities are very useful.     In this paper we investigated the statistics of density structure in compressible MHD turbulence using numerical experiments. We analyzed PDFs, spectra, structure functions, and the intermittency of density and the logarithm of density for experiments with different sonic ${\\cal M}_s$ and Alfv\\'{e}n ${\\cal M}_A$ Mach numbers. We addressed the issue of whether the numerical results can be compared with observations. We found:  \\begin{itemize} \\item The amplitude of density fluctuations strongly depends on ${\\cal M}_s$ both in weakly and strongly magnetized turbulent plasmas. \\begin{enumerate} \\item[a)] PDFs of density fluctuations have a lognormal distribution. \\item[b)] The width of the PDF of $\\log \\rho$ grows with ${\\cal M}_s$. \\item[c)] The flattening observed in density spectra is due to the contribution of the highly dense small-scale structures generated in the supersonic turbulence. \\end{enumerate} \\item Fluctuations of the logarithm of density are much more regular than those of density. The logarithm of density exhibits the GS95 scalings and anisotropies. The PDFs are close to Gaussian. \\item Structure functions of density and the logarithm of density reveal self-similarity. The density field is not self-similar to the velocity one. \\item The intermittency of density structure depends on both ${\\cal M}_s$ and ${\\cal M}_A$. We report only weak dependence on ${\\cal M}_s$ for subsonic turbulence. In the case of supersonic models the dependence is much stronger. \\item Additional care about resolution is necessary for intermittency studies. The She-L\\'{e}v\\^{e}que exponents results can be affected by the resolution. \\item Similar to their 3D counterparts, the 2D measures of column densities, i.e. the variance, skewness and kurtosis, depend on ${\\cal M}_s$ and can be used to determine the value of ${\\cal M}_s$ from observations. \\item The cloud boundary effects are important and should be accounted for by the study. The most sensitive to the effects of cloud boundaries are the PDFs and intermittency studies. \\item The genus statistic is sensitive to ${\\cal M}_s$. The 2D genus analysis is a promising tool for studying the topology of the ISM. \\end{itemize}"
    },
    "0608/astro-ph0608098_arXiv.txt": {
        "abstract": "We have studied the OH masers in the star forming region, W3(OH), with data obtained from the Very Long Baseline Array (VLBA). The data provide an angular resolution of $\\sim$5 mas, and a velocity resolution of 106 m s$^{-1}$. A novel analysis procedure allows us to differentiate between broadband temporal intensity fluctuations introduced by instrumental gain variations plus interstellar diffractive scintillation, and intrinsic narrowband variations. Based on this 12.5 hours observation, we are sensitive to variations with time scales of minutes to hours. We find statistically significant intrinsic variations with time scales of $\\sim$15--20 minutes or slower, based on the {\\it velocity-resolved fluctuation spectra}. These variations are seen predominantly towards the line shoulders. The peak of the line profile shows little variation, suggesting that they perhaps exhibit saturated emission. The associated modulation index of the observed fluctuation varies from statistically insignificant values at the line center to about unity away from the line center. Based on light-travel-time considerations, the 20-minute time scale of intrinsic fluctuations translates to a spatial dimension of $\\sim$2--3 AU along the sight-lines. On the other hand, the transverse dimension of the sources, estimated from their observed angular sizes of about $\\sim$3 mas, is about 6 AU. We argue that these source sizes are intrinsic, and are not affected by interstellar scatter broadening. The implied peak brightness temperature of the 1612/1720 maser sources is about $\\sim2\\times 10^{13}$ K, and a factor of about five higher for the 1665 line. ",
        "introduction": "\\label{sec-intro}  Interstellar hydroxyl maser sources are found in the Galaxy and many external galaxies. In the Galaxy, these sources are associated with star forming regions (predominantly in the 1665 and 1667 MHz lines), IR and late type stars (mainly at 1612 MHz) and supernova remnants (exclusively at 1720 MHz).  Numerous VLBI observations (e.g., Gwinn et al 1988; Fish et al 2005; 2006) have indicated that the apparent angular sizes of OH masers increase with distance in the Galaxy. OH maser sources in the inner Galaxy show larger angular sizes compared to sources in the outer Galaxy. These facts strongly suggest that the observed angular size is strongly dependent on interstellar scatter broadening. Some years ago, Burke et al (1968) and Gwinn et al (1988) have proposed that the apparent broadening of masers shows a wavelength dependence $\\propto\\lambda^2$, based on OH maser sizes at 1665 MHz and $H_2O$ maser sizes at 23 GHz. These data suggest that interstellar scattering is  the dominant cause for the observed angular sizes of these sources.  Desai, Gwinn \\& Diamond (1994) used the VLBA to study the details of the interstellar broadening of the OH masers in the distant ($\\sim$14 kpc) HII  region W49N. Anisotropic broadening was observed with the orientation of the minor axis preferentially parallel to the Galactic plane, an additional confirmation of the influence of interstellar scattering on the scattered sizes. An unambiguous determination of the intrinsic sizes of OH maser sources would represent an important input in understanding the OH maser pump and emission mechanism.  In addition to the intrinsic sizes, an important property of the masers that provides information about the pumping mechanism is the variability of the maser intensity as a function of time. Variability on long time scales (weeks to months) has been discussed by numerous investigators (Schwartz, Harvey \\& Barrett 1974; Coles, Rumsey \\& Welch 1968; Zuckerman et al 1972; Rickard, Zuckerman \\& Palmer 1975; Gruber \\& de Jager 1976; Clegg \\& Cordes 1991). The suggestion was made that the observed variability may arise from  changes in the number density of molecules exhibiting maser phenomenon, or physical conditions in the maser column. In contrast to the long timescale variations, short timescale variations have also been detected (Zuckerman et al 1972; Rickard et al 1975). Salem \\& Middleton (1978) provide  a model consisting of a sudden onset of a pumping mechanism that could  cause rapid quasi-periodic fluctuations in the observed intensity. The predicted fluctuations on time scales of roughly a day would have a $\\sim$25\\% modulation index.  Evans et al (1972) have investigated eight well known OH masers with the goal of describing the statistics of the radiation. They sampled the output of the 140 feet NRAO telescope rapidly and find that the radiation is of Gaussian nature to within a level of about one per cent.  The major existing study to date of short term variations consists of Arecibo (beam size of 3 arcmin) and VLA (beam size of 15$\\times$5 arcsec) observations probing variability on time scales in the range 16 s to 2 hours (Clegg \\& Cordes 1991). Typical variations were detected at the 5-10 per cent level with some large variations at the 100 per cent (or more) level for the sources W75S and NGC 6334F. As Clegg \\& Cordes point out, an identification of any intrinsic variability in these sources would provide a unique opportunity to determine the source extent using light-travel-time arguments. They do find prominent variations in some  of the sources with time scales of $\\sim$20 minutes (e.g S269). However, it is very difficult to distinguish between variations that are intrinsic to the source, and those that might arise from interstellar diffractive scintillations. These authors do not decisively choose between an intrinsic or scattering origin for the fluctuations. However, they do stress that if the fluctuations are due to interstellar scintillation, implausible brightness temperatures would be required.  As pointed out by Clegg \\& Cordes, several aspects can affect",
        "conclusions": "\\label{sec-discussion}  In this work, we have conclusively demonstrated the presence of ``narrowband'' intrinsic variations in these W3(OH) maser sources. These variations seem to have a typical time scale of about 15-20 minutes or longer, indicating that the faster time scale may correspond to the light travel time of the maser pumping column.  A very important aspect, which is worth stating again, is that the observed intrinsic variations that are clearly separable from other contaminants are necessarily ``narrowband'' in nature. All the other variations such as the instrumental gain variations and the interstellar diffractive scintillation, although their time scales may be of the order of a few tens of minutes, are distinguishable from the intrinsic variabilities, mainly because of their broad correlation bandwidth. The expected differential fluctuation within the line profile due to interstellar scintillation is only of the order of $\\sim$0.5\\%, much smaller than the observed intrinsic variability, whose magnitude was as much as 100\\% (modulation index), in the ``tail\" portions of the spectral lines.  There is one source of systematic narrowband error that could have potentially influenced our conclusions. In our VLBI measurements, if the Earth's motion were not correctly compensated for, certain systematic but spurious temporal variations in the line profiles would be induced due to differential Doppler shift. However, we would then expect the variations seen on the two sides of the line profile peak to be anti-correlated. We have examined carefully the relevant cross-correlations, and do not find any statistically significant signature such an anti-correlation, clearly indicating that our data are free of such artefact.  An important aspect of these maser sources is that their typical measured angular diameter is $\\sim$3 mas. With the distance to the star forming region of 2 kpc, this corresponds to a transverse distance scale of $\\sim$ 6 AU. As we have seen from Figure \\ref{fig:vrspec}, the fluctuation frequencies seen in the spectra are $\\lsim 10^{-3}$ Hz (fluctuation time scale, $T_f\\sim 1000$ s or longer). If indeed this time scale reflects the dimension of the source based on light travel time arguments, then the implied longitudinal spatial scale of the maser column would be 2--3 AU. It is important to note that the apparent transverse spatial dimension of the source as measured from the angular size is comparable to its longitudinal dimension implied by the fluctuation timescale, even though there seems to be no {\\it a priori} basis for the comparison, let alone the agreement. However, if indeed these orthogonal dimensions of the source are expected to comparable, it would imply that the VLBA observations may have actually resolved the intrinsic source size of the OH masers in W3(OH). Then the possible contribution of angular broadening caused by interstellar scattering to the apparent size of the source is minimal. This also implies that the expected decorrelation bandwidth of interstellar scintillation may be much wider than 100 kHz, further reinforcing the validity of our assumptions. However, if the scatterer is closer to the source instead of the midway location which is implicit in equation 1, the decorrelation bandwidth would be correspondingly narrower. We assess whether that may be the case, based on the two important estimates we have at hand, namely the measured upper-limit for scatter broadening $\\theta_H$ and the observed time scales (say, $T_{\\rm dif}$) of the narrow-band variability, if due to ISM, since that is our concern. How far are we from the naive assumption of uniform scatterer (or a strong scatterer mid-way) will be reflected in both of these parameters, which depend differently on the relative location of the scatterer (see for example, eq. 3 and 8 of Deshpande \\& Ramachandran 1998). If $D_s$ \\& $D$ are the distances of the scatterer and of the observer, respectively, as measured from the source, it is easy estimate an upper limit to ($D/D_s$ -1) using relevant expressions, and assuming that the observed time-scales are 1000 s or longer, \\& a wavelength of 18 cm. We find that ($D/D_s$ -1) should be less than or equal to $20/(\\theta_H V_{ism})$, where the scatter broadening $\\theta_H$ is in mas, and the ISM velocity $V_{ism}$ is in km s$^{-1}$. A value of 3 mas for $\\theta_H$ implies $(D/D_s -1) \\le 6.6/V_{ism}$. Noting that $V_{ism}$ will have the same lever-arm factor as source velocity does, and considering typical values of $V_{ism}$, we conclude that $(D/D_s -1)$ is close to unity, if not smaller. Hence, the decorrelation bandwidth that we estimate naively is most likely an under-estimate, given the above and that $\\theta_H$ may already be an over-estimate of the scatter broadening.  As already mentioned in the introduction, variability of astrophysical masers on longer time scales --- weeks to months --- observed before by several investigators, may be due to changes in the number density of relevant molecules or physical conditions in the maser column. On the other hand, very short time scale rapid variations have also been seen. Salem \\& Middleton (1978) suggest a model in which a sudden onset of a pumping mechanism can cause rapid quasi-periodic fluctuations in the observed flux density, and predict fluctuations with time scales of a day or so, with $\\sim$25\\% modulation index. These fluctuations may either correspond to propagation of a radiative or a collisional disturbance. In the former case the disturbance travels with the speed of light, and in the latter with a typical speed of $\\sim$10 km s$^{-1}$ or so. For radiative propagation of disturbance, with the typical flux density of the observed lines, the brightness temperature comes to $\\sim 10^{10}$K. For a collisional disturbance, the source dimension is only $T_f\\times 10$ km s$^{-1}$ $\\approx 10^7$ meters (angular diameter of $\\sim$0.1 $\\mu$arcsec). This corresponds to a brightness temperature of $\\sim 10^{20}$K (see also Clegg \\& Cordes 1991). In the latter case, since the implied intrinsic angular diameter is only $\\sim$1$\\mu$arcsec, the observed angular diameter is predominantly due to interstellar scattering. For an object at 2.2 kpc (distance of W3(OH) complex), a scatter broadening of 2--3 arcsec seems to be very large at the 18-cm wavelength. VLBI observations of several nearby (distance less than 2--3 kpc) pulsars at 327 MHz have shown a broadening of 10--20 mas, with an exception of the Vela pulsar, PSR B0833--45 (Gwinn, Bartel \\& Cordes 1993). In the case of Vela, the excess angular broadening is due to the presence of strong scattering screen (Deshpande \\& Ramachandran 1998). Assuming a Kolmogorov density irregularity spectrum, with the wavelength dependence of $\\lambda^{2.2}$, the expected angular broadening is only $<0.5$ mas at 1665 MHz. Based on these considerations, we argue against any significant over-estimation of the source size of W3(OH) due to scatter broadening.  It has been suggested by Elitzur (1991) that short time scale intensity variations can be produced by variations in the maser level population at small length scales. These fluctuations produced in the unsaturated medium of the maser core give rise to the spectrum of flux density variability observed. In our case, with the length scale of 2-3 AU (derived based on the light travel time arguments) represents an average ``seed'' length for such a fluctuation. This suggestion may well be the reason for the short time scale fluctuations that we have observed.  As mentioned earlier, the longitudinal dimension 2-3 AU estimated from the observed time scales of intrinsic variability compares well with the transverse spatial dimension of 6 AU estimated assuming that the apparent angular size to be approximately the intrinsic size of the source. However, for estimating brightness temperature the source we will use only its transverse size. A typical observed flux density of some 10 Jy at the line peaks (e.g. for 1612/1720 lines), implies a peak brightness temperature of $\\sim2\\times 10^{13}$ K. A five times higher value, i.e. $\\sim 10^{14}$ K, is implied by the correspondingly brighter peaks of the 1665 lines.  The intrinsic intensity variations that we observe are particularly confined to the ``shoulders'' of the lines and away from the line peak (see Figure \\ref{fig:variance}). The peak of the line profiles does not exhibit any statistically significant variation, in all of the six data sets. This behavior is not at all surprising if the peak of these line emissions corresponds to ``saturated'' maser action, and the rest of the line profile exhibits unsaturated emission. In any case, the absence of any intrinsic variability at the line peak enables use of the corresponding channel intensities to calibrate out any ``broadband'' variations from the dynamic spectrum.  Another interesting aspect apparent from the results in Figure \\ref{fig:variance} is that the statistically significant narrowband variations that we observe at the raising and falling edges of the line profiles are not identical between the Zeeman pairs (RR and LL components). For instance, the difference between the modulation indices of the two Zeeman components of 1720 MHz lines is clearly seen to be several times the root mean square noise level. The physical reason for this behavior needs to be explored. Given the birefringent nature of the medium close to or at the origin of the radiation, as well as interstellar propagation medium, effects of refraction on the apparent visibility of the two Zeeman components remain to be explored.  To summarize, we list our main conclusions as follows. \\begin{itemize} \\item The combination of our cross-correlation procedure and the variance analysis provides an effective tool for estimation and elimination of the ``broadband'' contribution from instrumental effects and the interstellar diffractive scintillation, and thus for identifying variations that are intrinsic and ``narrowband'' in nature. \\item Significant intrinsic ``narrowband'' variability is observed over most regions of the line profiles except at and about the line peak, suggesting ``saturated'' maser emission at the peak. \\item The velocity-resolved fluctuation spectra reveal that the time scales of the significant intrinsic variability are 15-20 minutes or longer. \\item Based on light travel time argument, the intrinsic variability time scale implies a longitudinal spatial scale of the maser column to be about 2-3 AU. \\item The apparent angular sizes of the sources are unlikely to be significantly affected by interstellar scattering. The implied transverse size of the source is comparable to its longitudinal dimension. \\item The peak brightness temperatures for the maser sources range between  $\\sim2\\times 10^{13}$ and $\\sim 10^{14}$ K. \\end{itemize}"
    },
    "0608/astro-ph0608267_arXiv.txt": {
        "abstract": "{Using the  {\\it International Ultraviolet Explorer} data archive, we have examined the {\\it SWP} echellograms of 74 B0--B2.5e stars for statistically significant fluctuations in the He\\,II (``H$\\alpha$\") $\\lambda$1640 line profile. In this sample we found that the He\\,II line is occasionally variable in 10 stars over short to long timescales. The He\\,II-variable stars discovered are $\\lambda$\\,Eri, $\\omega$\\,Ori, $\\mu$\\,Cen, 6\\,Cep, HD\\,67536, $\\psi^1$\\,Ori, $\\eta$\\,Cen, $\\pi$\\,Aqr, 2\\,Vul, and 19\\,Mon.  The most frequent two types of variability are an extended blue wing absorption and a weakening of the line along the profile. Other types of variability are a weak emission in the red wing and occasionally a narrow emission feature. In the overwhelming number of cases, the C\\,IV resonance doublet exhibits a similar response; rarely, it can exhibit a variation in the opposite sense. Similar responses are also often seen in the Si\\,IV doublet, and occasionally even the Si\\,III $\\lambda$1206 line.  We interpret the weakenings of He\\,II and of high-velocity absorptions of C\\,IV to localized decreases in the photospheric temperature, although this may not be a unique interpretation. We discuss the variable blue wing absorptions and red wing emissions in terms of changes in the velocity law and mass flux carried by the wind. In the latter case, recent experimental models by Venero, Cidale, \\& Ringuelet require that during such events the wind must be heated by 35,000\\,K at some distance from the star. ",
        "introduction": "\\label{int}  Together with the hydrogen lines, the lines of helium are among the most important in the spectra of hot stars. In atmospheres of B stars and most O stars, the dominant ion stage of helium is He$^+$, and the strongest of the He\\,II features is the $\\lambda$1640 (``H$\\alpha$\") line. In O stars this line is formed substantially in the wind -- so much so in supergiants that the line generally develops a strong P\\,Cygni emission structure. For spectral types later than O8--09, the line decreases in strength, but it remains visible as a photospheric diagnostic for spectra as late as B2.5 (Peters 1990, Rountree \\& Sonneborn 1991).  The He\\,II $\\lambda$1640 line is actually a complex of seven permitted transitions arising from lower levels at 40.8 eV. Its effective centroid wavelength is 1040.42\\,\\AA.~ In the outer atmospheres of hot stars this line is formed by the photoionization of He$^{2+}$ by extreme UV ($<$$\\lambda$228) photons, followed by recombination. The density and temperature sensitivities insures that the line's formation occurs substantially in the base of the wind or within the photosphere for O and B stars near the main sequence, respectively.  The $\\lambda$1640 line is mildly sensitive to departures from LTE in the He$^{1+}$ atom. As a result, the strengths computed from non-LTE models tend to be slightly stronger than those from LTE models. Since these effects are relatively small, there appears to be no major difficulties fitting this line approximately with conventional blanketed non-LTE model atmospheres.  Auer \\& Mihalas (1972) suggested that the near coincidence of central wavelengths of the He\\,II and some hydrogen lines could enhance emission of He\\,II $\\lambda$4686 and $\\lambda$1640 through optical pumping (Bowen fluorescence).  However, using more recent atomic parameters, Herrero (1987) demonstrated that these effects are negligible.  Recently,  Venero, Cidale, \\& Ringulet (2000; ``VCR\") have considered the behavior of the $\\lambda$1640 line for model atmospheres with $T_{\\rm eff}$ = 25,000\\,K and strong, isotropic, and heated winds. These authors find that even for model atmospheres of early-type B stars a fast and/or heated wind can alter the underlying photospheric profile. For example, in these models $\\lambda$1640 undergoes a near maximum {\\it absorption} strength in winds having a temperature determined by radiative equilibrium, i.e., with $T_o$/$T_{\\rm eff}$ = 0.6--0.8. However, if the wind is heated to 10,000\\,K above the $T_{\\rm eff}$, then emissions will be produced in one or both of the line wings. Thus, isotropic, heated ($T_o$ $\\ge$ $T_{\\rm eff}$) winds produce a P\\,Cygni-type profile, that is, with a distinctly blueshifted absorption and slightly redshifted emission. For standard wind models for early-type Be stars (unheated winds, with \\.M $\\sim$ 1$\\times$10$^{-8}$  M$_{\\odot}$ yr$^{-1}$), emission should be absent or undetectable. Even in the spectrum of the O9\\,V star 10\\,Lac, with its mass loss rate of 1.7$\\times$10$^{-7}$ M$_{\\odot}$ yr$^{-1}$, the $\\lambda$1640 profile is unshifted and has no redshifted emission (see Figure\\,18a of Brandt et al. 1998). It can be added that unpublished thesis work by Cidale (1993) suggests that these same trends continue with emission in C\\,IV. As a postulated hot temperature rise in these models is moved outward, a mild red emission component in C\\,IV is enhanced while the absorption component remains almost the same. Some of our examples of observed variations below will illustrate this behavior.  The $\\lambda$1640 line has been interpreted by some authors to be stable in strength and thus to be a good measure of an O or early B star's effective temperature.  As detailed below, variability in this line has only seldom been reported in hot, chemically homogeneous, single stars. According to the VCR models, any variability of this line would require substantial changes in the star's effective temperature, mass loss rate, or restructuring of the wind stratification. In cool stars $\\lambda$1640 variations  arise from variable EUV irradiation in their chromospheres (Linsky et al. 1998). In this paper we will test several claims in the literature for $\\lambda$1640 variability in Be star spectra and extend the search for this variability to a larger sample of B stars. Positive results of this search will be placed in the general context of the predictions of the VCR's {\\it ad hoc} and nonstandard models of winds in hot stars. ",
        "conclusions": "This paper provides a mini-atlas of He\\,II $\\lambda$1640 variability for a group of 10 Be stars selected from a much larger sample of early-type Be, Bn, and B normal stars. We have identified several basic types of variability. Weak red emissions and line weakenings occur over timescales of a few hours or less. In terms of the variability timescales, we have noted that the pattern of strengthening blue wings occurs over long timescales. In our view this is most likely explained by changes in the wind velocity law (cause unknown). Second, line weakenings likewise occur preferentially over long timescales. Long-term weakenings occur in half of our ten $\\lambda$1640-variable stars. Third, weakenings over the whole line or only the red wing can occur even within a few hours. We also point out that rapid variability was found preferentially in the stars $\\eta$\\,Cen and $\\lambda$\\,Eri. We believe these events speak to intrinsic properties of these stars rather than to observational sampling.  We have suggested that the properties which change the surface and wind properties of these stars are mediated by magnetic instabilities. This is among the few ways of interpreting aperiodic variations of a single star on a timescale sometimes much less than its rotation period. We note that our examples of variable $\\lambda$1640 do not include known Bp stars. Magnetic fields in these stars are thought to be dipolar and, most importantly, stable over at least several years. Strong stable fields resist the influence of hydrodynamical instabilites that might alter a  wind's structure or its geometrical flow. Therefore, we conjecture that at least for some of these stars magnetic fields must be localized on the surface.  Velocity perturbations due to nonradial pulsations and differential surface rotation would then offer plausible ways to trigger magnetic instabilities in these multipolar configurations."
    },
    "0608/astro-ph0608022.txt": {
        "abstract": "{} {We derive accurate parameters related to the CMD, structure and dynamical state of M\\,52 and NGC\\,3960, whose fields are affected by differential reddening. Previous works estimated their ages in the ranges $35 - 135$\\,Myr and $0.5 - 1.0$\\,Gyr, respectively.} {\\jj, \\hh\\ and \\ks\\ 2MASS photometry with errors $<0.2$\\,mag is used to build CMDs, radial density profiles, luminosity and mass functions, and correct for differential reddening. Field-star decontamination is applied to uncover the cluster's intrinsic CMD morphology, and colour-magnitude filters are used to isolate stars with high probability of being cluster members.} {The differential-reddening corrected radial density profile of M\\,52 follows King's law with core and limiting radii of $\\rm\\rc=0.91\\pm0.14\\,pc$ and $\\rm\\rl=8.0\\pm1.0\\,pc$. NGC\\,3960 presents an excess of the stellar density over King's profile ($\\rm\\rc=0.62\\pm0.11\\,pc$ and $\\rm\\rl=6.0\\pm0.8\\,pc$) at the center. The tidal radii of M\\,52 and NGC\\,3960 are $\\rt=13.1\\pm2.2$\\,pc and $\\rt=10.7\\pm3.7$\\,pc. Cluster ages of M\\,52 and NGC\\,3960 derived with Padova isochrones are constrained to $60\\pm10$\\,Myr and $1.1\\pm0.1$\\,Gyr. In M\\,52 the core MF ($\\chi_{\\rm core}=0.89\\pm0.12$) is flatter than the halo's ($\\chi_{\\rm halo}=1.65\\pm0.12$). In NGC\\,3960 they are $\\chi_{\\rm core}=-0.74\\pm0.35$ and $\\chi_{\\rm halo}=1.26\\pm0.26$. The mass locked up in MS/evolved stars in M\\,52 is $\\sim1200$\\,\\ms, and the total mass (extrapolated to 0.08\\,\\ms) is $\\sim3800$\\,\\ms. The total mass in NGC\\,3960 is $\\sim1300$\\,\\ms.} {Compared to open clusters in different dynamical states studied with similar methods, the core and overall parameters of M\\,52 are consistent with an open cluster more massive than 1\\,000\\,\\ms\\ and $\\sim60$\\,Myr old, with some mass segregation in the inner region. The core of NGC\\,3960 is in an advanced dynamical state with strong mass segregation in the core/halo region, while the somewhat flat overall MF ($\\chi\\approx1.07$) suggests low-mass star evaporation. The excess stellar density in the core may suggest post-core collapse. The dynamical evolution of NGC\\,3960 may have been accelerated by the tidal Galactic field, since it lies $\\approx0.5$\\,kpc inside the Solar circle.} ",
        "introduction": "\\label{intro}  Open clusters (OCs) can be used as test-beds of molecular cloud fragmentation, star formation, and stellar and dynamical evolution models. They are excellent probes of the Galactic disc structure (Janes \\& Phelps \\cite{JP94}; Friel \\cite{Friel95}; Bonatto et al. \\cite{BKBS06}). This follows from the relative simplicity in age and distance estimates, at least for bright clusters. Accurate OC parameters such as core and limiting radii, age, mass, density, and relaxation times are essential to constrain theoretical models. However, because most of the OCs lie close to the disk, and thus are more affected both by reddening and field-star contamination (Bonatto et al. \\cite{BKBS06} and references therein), this kind of analysis becomes restricted to the more populous OCs and/or those located a few kpc from the Sun.  According to the WEBDA\\footnote{\\em http://obswww.unige.ch/webda} OC database (Mermilliod \\cite{Merm1996}) the current number of OCs with known parameters (such as coordinates, age, distance from the Sun, and reddening) amounts to $\\sim700$. The catalogue of Dias et al. (\\cite{Dias2002}) includes 1756 optically visible OCs and candidates. More recently Bonatto et al. (\\cite{BKBS06}) simulated observational completeness in different directions throughout the Galaxy and estimated that $\\sim50$\\% of Trumpler types I to III OCs remain undetected in the solar neighbourhood (distance from the Sun $\\ds\\leq1.3$\\,kpc). A large fraction of the faint and/or poorly populated OCs end up drowned in the stellar background, particularly in directions intercepting the bulge.  \\begin{figure*} \\begin{minipage}[b]{0.50\\linewidth} \\includegraphics[width=\\textwidth]{fig1a.ps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.50\\linewidth} \\includegraphics[width=\\textwidth]{fig1b.ps} \\end{minipage}\\hfill \\caption[]{$\\rm20\\arcmin\\times20\\arcmin$ XDSS R images of M\\,52 (left panel) and NGC\\,3960 (right panel).} \\label{fig1} \\end{figure*}  Several decades of studies have provided fundamental parameters for a considerable fraction of the catalogued OCs (e.g. Lyng\\aa\\ \\cite{Lynga82}; WEBDA; Dias et al. \\cite{Dias2002}; Kharchenko et al. \\cite{Kharchenko05}; Bonatto et al. \\cite{BKBS06}). Some clusters have been studied twice or more and a disagreement of a factor of $\\sim2$ in parameters is not uncommon. Discrepancies arise mostly from different analytical methods dealing with optical or IR photometry, Colour-Magnitude Diagrams (CMDs), isochrone sets, photometric systems, etc.  Recently we started a program to derive fundamental parameters of OCs with near-IR photometry and a series of analytical procedures developed to minimise field-star contamination of CMDs. They have proven to be effective for parameter determination for a wide class of OCs. For spatial and photometric uniformity we work with \\jj, \\hh\\ and \\ks\\ 2MASS\\footnote {The Two Micron All Sky Survey, All Sky data release (Skrutskie et al. \\cite{2mass1997}), available at {\\em http://www.ipac.caltech.edu/2mass/releases/allsky/}} photometry. The 2MASS Point Source Catalogue (PSC) is uniform reaching relatively faint magnitudes covering nearly all the sky, allowing a proper background definition for clusters with large angular sizes (e.g. Bonatto, Bica \\& Santos Jr. \\cite{BBS2005}; Bonatto, Bica \\& Pavani \\cite{BBP2004}).  The use of field-star decontamination and colour-magnitude filters have produced more robust parameters (e.g. Bonatto \\& Bica \\cite{BB2005}; Bonatto et al. \\cite{BBOB06}). In particular, field-star decontamination constrains the age more, especially for low-latitude OCs (Bonatto et al. \\cite{BKBS06}). These procedures have proven useful in the analysis also of faint and/or distant OCs (Bica, Bonatto \\& Dutra \\cite{BBD03}; Bonatto, Bica \\& Dutra \\cite{BBD04}; Bica \\& Bonatto \\cite{LowC05}).  Another potential source of observational uncertainties in OC parameters is the presence of differential reddening. Spatial variation of reddening may affect the age determination by isochrone fitting and, consequently, the distance from the Sun. This in turn could affect the absolute values of core and limiting radii, and cluster mass.  Our goals in this paper are twofold: {\\em (i)} add to the previous techniques a procedure to take into account the differential reddening, and {\\em (ii)} derive more constrained parameters for two OCs of different ages. As targets we selected the relatively bright OCs M\\,52 and NGC\\,3960 that in the literature have uncertain age determinations. The age of M\\,52 has been estimated in the range $35 - 135$\\,Myr (Sect.~\\ref{M52}), while that of NGC\\,3960 varies from $0.5$ to $1.0$\\,Gyr (Sect.~\\ref{N3960}). Previous works have shown evidence of differential reddening in the fields of both clusters (e.g. Pandey et al. \\cite{Pandey2001}; Prisinzano et al. \\cite{Prisinzano04}). Besides reddening, age and distance from the Sun, in this paper we derive the core and limiting radii, and the stellar mass stored in the core and halo subsystems. The latter parameters for NGC\\,3960 are derived for the first time.  This paper is organized as follows. In Sect.~\\ref{M52N3960} we review previous findings on M\\,52 and NGC\\,3960. In Sect.~\\ref{2mass} we present the 2MASS data, subtract the field-star contamination, estimate the differential reddening, derive fundamental cluster parameters, and analyse the cluster density structure. In Sect.~\\ref{MF} we derive luminosity and mass functions (LFs and MFs), and compute stellar content properties. In Sect.~\\ref{dyna} we discuss the dynamical states of both OCs. In Sect.~\\ref{CWODS} we use diagnostic diagrams to analyse the structure and dynamical state of M\\,52 and NGC\\,3960 as compared to other OCs. Concluding remarks are given in Sect.~\\ref{Conclu}. ",
        "conclusions": "\\label{Conclu}  In this paper we analyzed the structure and distribution of stars in the OCs M\\,52 and NGC\\,3960 with \\jj, \\hh\\ and \\ks\\ 2MASS photometry, restricted to stars with observational uncertainties $\\epsilon_{J,H,K_S}<0.2$\\,mag. To derive cluster parameters we took into account the differential reddening and applied field-star decontamination and colour-magnitude filters.  Compared to previous works we found differences in some parameters, such as cluster radius, mass and MF slopes. Most of those studies were based on spatial and magnitude-limited data. However, the discrepancies can be largely accounted for by the absence of a proper field-star decontamination.  Reddening throughout the field of M\\,52 varies from $\\ebv\\approx0.50$ to $\\ebv\\approx0.93$, with a marked gradient from south to north and a near uniformity in right ascension. For NGC\\,3960 we derive the range $\\ebv\\approx0.03 - 0.34$, uniform both in right ascension and declination. The projected area considered in both fields was $35\\arcmin\\times35\\arcmin$.  For M\\,52 we derived an age of $60\\pm10$\\,Myr, absolute distance modulus $\\mMo=10.8\\pm0.1$, and distance from the Sun $\\ds=1.4\\pm0.2$\\,kpc. For NGC\\,3960 we found an age of $1.1\\pm0.1$\\,Gyr, $\\mMo=11.2\\pm0.1$, and $\\ds=1.7\\pm0.2$\\,kpc. M\\,52 is located $\\approx0.7$\\,kpc outside the Solar circle, while NGC\\,3960 is $\\approx0.5$\\,kpc inside it.  Radial density profiles built with colour-magnitude filtered photometry are much deeper and yield more constrained OC structural parameters than the observed ones. In particular, King's model provided an excellent fit to the differential-reddening corrected radial density profile of M\\,52, from the external parts to the inner region. We found core and limiting radii of $\\rc=0.91\\pm0.14$\\,pc and $\\rl=8.0\\pm1.0$\\,pc. The inner region of NGC\\,3960, on the other hand, presents a clear excess of the stellar density over King's profile, which may suggest post-core collapse in this 1.1\\,Gyr OC. Parameters derived are $\\rc=0.62\\pm0.11$\\,pc and $\\rl=6.0\\pm0.8$\\,pc. The tidal radii of M\\,52 and NGC\\,3960, computed with the differential-reddening corrected RDPs are $\\rt=13.1\\pm2.2$\\,pc and $\\rt=10.7\\pm3.7$\\,pc.  The present use of colour-magnitude filtered photometry to build RDPs excludes a large fraction of non-cluster stars. As a result, observed and colour-magnitude filtered RDPs present significant differences (especially depth and shape) as shown in Fig.~\\ref{fig5}. Because of the small residual field-star contamination in the present RDPs, they can be used to derive more realistic intrinsic structural parameters. Nilakshi et al. (\\cite{Nilakshi2002}) built RDPs using a photographic database (DSS photometry). They pointed out that the angular size was difficult to determine accurately due to the weak contrast between cluster and field stars. Kharchenko et al. (\\cite{Kharchenko05}) derived parameters based on stars with available proper motion, which restricts the analysis to brighter stars. For evolved clusters in which a significant fraction of low-mass stars has already been displaced from the center outwards, their brighter cutoff photometry may explain differences, together with lower statistics for a smaller sample of  cluster stars.  The mass stored in observed stars ($1.3\\leq m(\\ms)\\leq6.1$) in M\\,52 is $m_{\\rm overall}\\approx1.2\\times10^3$\\,\\ms; the core contains $m_{\\rm core}\\approx240$\\,\\ms. Extrapolating to the H-burning limit with Kroupa's IMF these values become $m_{\\rm overall}\\approx3.8\\times10^3$\\,\\ms\\ and $m_{\\rm core}\\approx480$\\,\\ms. For NGC\\,3960 we derived $m_{\\rm overall}\\approx1.3\\times10^3$\\,\\ms\\ and $m_{\\rm core}\\approx114$\\,\\ms.  We conclude that both structurally and dynamically, the core and overall parameters of M\\,52 are consistent with those expected of an OC more massive than 1\\,000\\,\\ms\\ and $\\approx60$\\,Myr old. The somewhat flat core MF slope ($\\chi_{\\rm core}\\approx0.89$) as compared to the steeper one at the halo ($\\chi_{\\rm halo}\\approx1.65$) and the Salpeter-like overall MF ($\\chi_{\\rm overall}\\approx1.48$) suggests mild mass segregation in the inner region. This picture is consistent with the low values of the core/overall dynamical-evolution parameters, $\\tau_{\\rm core}\\approx13 - 78$ and $\\tau_{\\rm overall}\\approx0.2 - 2.2$.  The parameters of NGC\\,3960, on the other hand, suggest that this OC - particularly the core - has reached an advanced dynamical state characterised by $\\tau_{\\rm core}\\sim10^3$ and $\\tau_{\\rm overall}\\sim11$, respectively for the core and overall cluster. NGC\\,3960 presents evident signs of mass segregation in the core/halo region, with a very flat MF slope ($\\chi_{\\rm core}\\approx-0.74$) and a steeper one at the halo ($\\chi_{\\rm halo}\\approx1.26$). Besides, the somewhat flat (compared to Salpeter) overall MF ($\\chi_{\\rm overall}\\approx1.07$) suggests low-mass star evaporation through the cluster rim. Tidal disruption due to the Galactic field may have boosted the dynamical evolution of NGC\\,3960, since it is located $\\approx0.5$\\,kpc inside the Solar circle.  Changes in parameters of M\\,52 and NGC\\,3960 because of differential-reddening correction are not significant, at least in the near-IR and/or because of the reddening ranges $\\ebv=0.50 - 0.93$ (M\\,52) and $\\ebv=0.03 - 0.34$ (NGC\\,3960). With respect to the analytical methods we conclude that colour-magnitude filters are essential to improve cluster parameters, particularly those related to structure."
    },
    "0608/astro-ph0608501_arXiv.txt": {
        "abstract": "Tissot indicatrices have provided visual measures of local area and isotropy distortions.  Here we show how large scale distortions of flexion (bending) and skewness (lopsidedness) can be measured.  Area and isotropy distortions depend on the map projection metric, flexion and skewness, which manifest themselves on continental scales, depend on the first derivatives of the metric.  We introduce new indicatrices that show not only area and isotropy distortions but flexion and skewness as well.  We present a table showing error measures for area, isotropy, flexion, skewness, distances, and boundary cuts allowing us to compare different world map projections.  We find that the Winkel-Tripel projection (already adopted for world maps by the National Geographic), has low distortion on most measures and excellent quality overall. ",
        "introduction": "Tissot (1881) indicatrices have been very useful for providing a visual presentation of local distortions in map projections in a simple and compelling fashion.  A small circle of tiny radius (say 0.1 degree of arc in radius) is constructed at a given location, and then enlarged and projected on the map at that location.  This always produces an ellipse.  Usually one favors conformal map projections that minimize the changes in scale factor, or equal area projections that minimize anisotropy, or recently, map projections that are neither conformal nor equal area, but which have a judicious combination of minimizing both scale and isotropy errors (like the Winkel-tripel used by the National Geographic Society for world maps).  The Tissot ellipse at a given location is specified by three parameters, the major axis, the minor axis, and the orientation angle $\\theta$ of the major axis of the ellipse.  Geometrically, from differential geometry (and General Relativity) we know that the measurement of local distances is measured by the metric tensor $g_{ab}$, where a and b can each take the values x or y, yielding three independent components.  Locally for two nearby points separated by infinitesimal map coordinate differences dx and dy, the true distance between these two points on the globe is given by: \\begin{equation} ds^2=g_{xx}dx^2+2 g_{xy}dx\\ dy +g_{yy}dy^2 \\end{equation} The Tissot ellipse (major axis, minor axis, and orientation angle, $\\theta$) can be calculated from the three components of the metric tensor.  Thus, the Tissot ellipse essentially carries the information on the metric tensor for the map.  It tells us how local infinitesimal distances on the map correlate with local infinitesimal distances on the globe.  As we will see, however, the Tissot ellipse does not carry all of the information related to distortions, which has been noted previously. Others (Stewart 1943; Peters 1975, 1978; Albinus 1981; Canters 1989, 2002; Laskowski 1997ab) have noted that there are finite distortions apart from those of the Tissot ellipse.  Previous authors have previously depicted finite size distortions by several methods, including showing faces on a globe (Reeves 1910; Gedymin 1946), 30 degree x 30 degree equiangular quadralaterals (Chamberlin 1947), a net of twenty spherical triangles, in an icosahedral arrangement (Fisher \\& Miller 1944), 150 degree great circle arcs, and circles of 15 degrees radius Tobler (1964).  In the Oxford Hammond Atlas of the World (1993) , new conformal map projections (\"Hammond Optimal Conformal Projections\") were designed for the continents.  Following the Chebyshev criterion (Snyder 1993), the rms scale factor errors were minimized by producing a constant scale factor along the boundary of the continent.  (For a circular region this conformal map would be a stereographic projection.)  By tailoring the boundary to the shape of each continent, the errors could be reduced relative to those in a simple stereographic projection.  In touting the advantages of their projection the Hammond Atlas did the following experiment.  Following Reeves (Reeves 1910), they constructed a face on the globe with a triangular nose, a straight (geodesic) mouth, and eyes that were pairs of concentric circles on the globe.  They then showed this face with various map projections. In the gnomonic, the mouth was straight, but the eye circles were not circular and were not concentric (what we would call skewness).  The Mercator projection had the mouth smiling (what we would call flexion) and although the eyes were circular they were not concentric.  The Hammond Optimal Conformal projections did a bit better on these qualities because the gradients of the scale factor changes were small so the flexion and skewness were small, although of course not zero.  In Sec. \\ref{sec:flexion} we introduce the concept of ``flexion'', by which a map projection can cause artificial bending of large structures.  In Sec. \\ref{sec:skewness}, we show another distortion: skewness, which represents lopsidedness and an asymmetric stretching of large structures.  We show a simple way to visualize these distortions in Sec. \\ref{sec:gg}, in which we introduce the ``Goldberg-Gott Indicatrices.'' In Sec. \\ref{sec:geometry} we derive a differential geometry approach to measuring flexion and skewness. While readers interested in computing flexion and skewness on projections not included in this paper should refer to Sec. \\ref{sec:geometry}, it is highly technical, and those interested in seeing results may skip directly to Sec. \\ref{sec:discussion}, in which we discuss Monte Carlo estimates of the distortions for a number of projections, a ranking of map projections, and our conclusions. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Numerical Analysis of Standard Map Projections}  \\label{sec:numerical}  Not all projections produce such simple results.  Thus, in general, we will want to compute the local flexion and skewness numerically.  Our approach is as follows: For each projection we chose 30,000 points selected randomly on the surface of a globe.  For each of these points, we chose a random direction to start a geodesic.  We follow that geodesic using small steps ($d\\tau\\simeq 10^{-5}$ rad) numerically, and use standard difference methods to compute the map velocity and acceleration along the geodesic.  We are thus able to compute the metric and the Christoffel symbols (and thus the flexion and skewness) directly.  We make our IDL (Interactive Data Language) code available to the interested reader at our projection webpage (see below).  Likewise, we also do a distance test, in which pairs of points, (i,j), are chosen at random and the distance is measured both on the globe and on the map.  This is a somewhat different perspective than simply inspecting the Goldberg-Gott indicatrices at a few locations, since we are now doing a uniform sample over the surface of the globe, rather than a uniform sampling over the {\\it map}.  When looking at the indicatrix map we can occasionally get a distorted view as to the quality of a particular projection.   Some (like the Mercator) have relatively good fits over most of the globe, but the high latitudes can, in principle, be projected to infinite areas, and thus, the reader may erroneously think the Mercator infinitely bad.  By sampling uniformly over the globe, we get a fair assessment of the overall quality of a particular projection.  We define a number of fit parameters: I, corresponding to errors in the local isotropy (zero for conformal projections), A, corresponding to errors in the Area (zero for equal area projections), F, corresponding to flexion (defined in the discussion of flexion, above), S, corresponding to skewness (also defined above), D, corresponding to distance errors, and B, corresponding to the average number of map boundary cuts crossed by the shortest geodesic connecting a random pair of points. \\begin{eqnarray} I&=&RMS\\left(\\ln \\frac{a_i}{b_i}\\right)\\\\ A&=& RMS\\left( \\ln a_i b_i-\\langle \\ln a_i b_i \\rangle  \\right)\\\\ F&=& \\langle |f_i|\\rangle\\\\ S&=& \\langle |s_i|\\rangle\\\\ D&=& RMS\\left( \\ln\\frac{d_{ij,map}}{d_{ij,globe}} \\right)\\\\ B&=&\\frac{L_{B}}{4\\pi} \\end{eqnarray} where $a_i$ and $b_i$ are the major and minor axes of the local Tissot ellipses of random point, $i$, $\\langle X_i \\rangle$ indicates the mean of property, X, and $L_B$ is the total length of the boundary cuts.  Logarithmic errors have been used before by Kavrayskiy (1958). Bugayevskiy and Snyder (1995) adopted schemes that weighted isotropy errors, $(a/b -1)^2$, and area errors, $(ab - 1)^2$, according to one's interest in isotropy and area. Airy (1861) had just given these terms equal weight.  \\begin{table}[h!] \\begin{tabular}{|l|c c c c c c|} \\hline Projection & I & A & F&  S & D  & B\\\\ \\hline\\hline {\\bf Non Interrupted Projections} & & & & & &\\\\ Azimuthal Equidistant & 0.87 & 0.60 & 1.0 & 0.57 & 0.356 & 0 \\\\ Gott-Mugnolo Azimuthal          & 1.2  & 0.20 & 1.0 & 0.59 & 0.341 & 0\\\\ Lambert Azimuthal     & 1.4  & 0    & 1.0 & 2.1  & 0.343 & 0 \\\\ Stereographic         & 0    & 2.0  & 1.0 & 1.0  & 0.714 & 0\\\\ \\hline {\\bf 1 180 deg. Boundary Cut} & & & & & & \\\\ Breisemeister   & 0.79 & 0   & 0.81 & 0.42 & 0.372  & 0.25\\\\ Eckert IV       & 0.70 & 0   & 0.75 & 0.55 & 0.390 & 0.25\\\\ Eckert VI       & 0.73 & 0   & 0.82 & 0.61 & 0.385 & 0.25 \\\\ Equirectangular & 0.51 & 0.41& 0.64 & 0.60 & 0.449 & 0.25 \\\\ Gall-Peters     & 0.82 & 0   & 0.76 & 0.69 & 0.390 & 0.25 \\\\ Gall Stereographic & 0.28 & 0.54 & 0.67 & 0.52 & 0.420 & 0.25\\\\ Gott Elliptical & 0.86 & 0   & 0.85 & 0.44 & 0.365 & 0.25 \\\\ Gott-Mugnolo Elliptical & 0.90 & 0 & 0.82 & 0.43 & 0.348 & 0.25\\\\ Hammer          & 0.81 & 0   & 0.82 & 0.46 & 0.388 & 0.25 \\\\ Hammer-Wagner   & 0.687 & 0 & 0.789 & 0.518 & 0.377 & 0.25 \\\\ Kavrayskiy VII  & 0.45 & 0.31& 0.69 & 0.41 & 0.405 & 0.25 \\\\ Lagrange        & 0    & 0.73& 0.53 & 0.53 & 0.432 & 0.25 \\\\ Lambert Conic   & 0    & 1.0 & 0.67 & 0.67 & 0.460 & 0.25 \\\\ Mercator        & 0    & 0.84& 0.64 & 0.64 & 0.440 & 0.25 \\\\ Miller          & 0.25 & 0.61& 0.62 & 0.60 & 0.439 & 0.25 \\\\ Mollweide       & 0.76 & 0   & 0.81 & 0.54 & 0.390 & 0.25 \\\\ Polyconic       & 0.79 & 0.49& 0.92 & 0.44 & 0.364 & 0.25 \\\\ Sinusoidal      & 0.94 & 0   & 0.84 & 0.68 & 0.407 & 0.25 \\\\ Winkel-Tripel   & 0.49 & 0.22& 0.74 & 0.34 & 0.374 & 0.25 \\\\ Winkel-Tripel (Times) & 0.48 & 0.24 & 0.71 & 0.373 & 0.39 & 0.25\\\\ \\hline   {\\bf 1 360 deg. Boundary Cut} && & & & &\\\\ Lambert Azimuthal (2 hemisphere) & 0.36 & 0    & 0.52 & 0.11 & 0.432&0.5\\\\ Stereographic (2 hemisphere)     & 0    & 0.39 & 0.37 & 0.37 & 0.692 &0.5\\\\ \\hline   {\\bf Multiple Boundary Cut Projections} & & & & & & \\\\  Gnomonic Cube   & 0.22 & 0.37 &   0.12  &   0.87  &  0.43 & 0.686 \\\\\\hline \\end{tabular} \\label{tab:global} \\caption{The errors in isotropy, area, flexion, skewness, distances , and boundary cuts for some standard projections.} \\end{table}  In Table~\\ref{tab:global} we compare a these measures for a number of standard projections, which, for fairness of comparison we divide into a number of categories.  The Gott-Elliptical, The Gott-Mugnolo Elliptical, and the Gott-Mugnolo Azimuthal have been discussed in Gott, Mugnolo \\& Colley (2007) and Gott, et al. (2007), where they have been applied to the earth, Mars, the moon, and the Cosmic Microwave all sky map.  First, we show projections which represent the complete globe without interrupts.  These projections are azimuthal and the average flexion over these maps is 1.  Second, we show the set of whole earth projections with one $180^\\circ$ interrupt.  These include rectangular and elliptical projections.  Note that among all of the complete projections with 0 or 1 $180^\\circ$ interrupt, there are a number of ``winners'' with regards to performance for flexion and skewness.  The Lagrange has the smallest flexion.  The Winkel-Tripel has the smallest skewness.  For all conformal projections, the skewness is equal to the flexion.  Of all of the whole earth projections, the most accurate for distance measure between points is the Gott-Mugnolo Azimuthal, followed very closely by the Lambert Azimuthal (Gott, Mugnolo \\& Colley 2007).  In the third and fourth groups, we show 2-hemisphere and other multiple cut projections, respectively.  In the final group, we have a projection with multiple interrupts, the gnomonic cube, which is defined piecemeal.  This is a particularly interesting projection since the gnomonic is locally flexion-free, but it is clear that geodesics will {\\it not} trace out straight lines in the gnomonic cube cube map because they bend when they cross an edge between faces.  The gnomonic cube is presented as a cross, so 5 edges are included in the map proper.  Geodesics bend when they cross an edge in this laid out cross configuration.  The above comparisons do not depend on how important each of the criteria are (i.e. what weighting to give each measure).  Laskowski (1997ab) suggested a means of ranking very disparate maps. Though his weighting scheme is not unique, as a simple illustration of how this can be done, we will minimize the sum of the squares of all 6 parameters, normalized to their values in the equirectangular projection: \\begin{equation} \\Sigma_\\varepsilon=\\left(\\frac{I}{N_i}\\right)^2+ \\left(\\frac{A}{N_a}\\right)^2+ \\left(\\frac{F}{N_f}\\right)^2+ \\left(\\frac{S}{N_s}\\right)^2+ \\left(\\frac{D}{N_d}\\right)^2+ \\left(\\frac{B}{N_b}\\right)^2 \\end{equation} Following Laskowksi (1997ab) we set the normalization constants equal to the values of these errors in the Equirectangular projection ($x=\\lambda$, $y=\\phi$): $N_i= 0.51$, $N_a = 0.41$, $N_f = 0.64$, $N_s = 0.60$, $N_d = 0.449$, $N_b = 0.25$.  The projections with the lowest values of $\\Sigma_\\varepsilon$ are: \\begin{enumerate} \\item Winkel-Tripel       (4.5629) \\item Winkel-Tripel (Times Atlas)  (4.5687) \\item Kavrayskiy VII  (4.8390) \\item Gall Stereographic (5.7582) \\item Hammer-Wagner (5.7847) \\item Eckert IV       (5.8519) \\end{enumerate} This approach is certainly not unique.  One may take issue with the weighting of the individual parameters, the domain over which they are applied (the whole earth, as opposed to continents only, for example), or even how the parameters are computed, we present it as a simple example of how our results may be combined with previous studies of map projections.  It is interesting that the ``best'' map, as selected by this criterion is the one already used by the National Geographic for it's whole world projection.  It is also interesting to note that the Winkel-Tripel has especially low skewness.   Another nice property of this weighting scheme is that it gives a low score to pathological projections.  For example, a series of n gores (made using the polyconic projection) arranged in a sunflower pattern would approximate the Azimuthal equidistant projection in distance errors as n became large but would have arbitrarily low values of I, A, F, and S. But if a boundary cut term B is included this term would blow up and save us from picking the bad subdivided map as better than the more visually pleasing projections described above.  Stretching individual pixels at the edges of the gores to fill the gaps between them would eliminate the boundary cuts.  However, it would also cause the skewness to blow up, preventing us from giving a good score to a bad projection.  Interested readers may visit {\\tt www.physics.drexel.edu/$\\sim$goldberg/projections/} to download a free IDL code to measure the flexion, area, and other measures discussed in this paper.  We have not done all known projections, but have covered ones that have available mathematical formulas and we thought likely to do well.  \\subsection{Conclusions}  We have developed two new measures of curvature distortions found in maps of the earth.  The Skewness and the Flexion can be used to identify particularly warped sections of a map, or to identify features (such as the U.S. Canada border), which appear as a straight line on some map projections but do not, in fact, follow geodesics. We have developed a new graphical tool, called the Goldberg-Gott indicatrices, which can be used to show these distortions, along with those in area and shape, at many points in a map, and have produced indicatrix maps for a number of popular projections.  Finally, we have used these measures to produce global distortion measures for different projections.  We have found that the Winkel-Tripel produces low distortion on most measures, and, in particular, has the lowest skewness of all projections in our sample with a $180^\\circ$ boundary cut."
    },
    "0608/astro-ph0608447_arXiv.txt": {
        "abstract": "We study 87 Extremely Red Objects (EROs), selected both to have color redder than~$R-[3.6]=4.0$, and to have confirmed spectroscopic redshifts. Together, these two constraints result in this sample populating a fairly narrow redshift range at $0.76 < z < 1.42$. The key new ingredient included here is deep \\emph{Spitzer} Space Telescope InfraRed Array Camera (IRAC) data. Based on $[3.6]-[8.0]$ color, we demonstrate that it is possible to classify EROs into early-type, dusty starburst, or power-law (AGN) types. We present ultraviolet to mid-infrared spectral energy distributions (SEDs) and Advanced Camera for Surveys (ACS) images, both of which support our simple IRAC color classification. ",
        "introduction": "\\label{sec:intro}  First discovered in the late 1980s \\citep{elston-88}, extremely red objects (EROs) are defined by their very red optical/near-infrared colors. It has been known for some time that the redness of their color constrains these galaxies to be either early-type galaxies, starburst galaxies reddened by dust, or AGN  (or a combination of these three classes). However, until recently, with observations limited to $K$-band (or shorter wavelengths), it has proven extremely challenging to accurately classify the EROs by type \\citep{mann-02}, even in combination with high-resolution HST imaging \\citep{mous-04}.   Although EROs appear to consist of a heterogeneous mix of galaxy classes, the emerging paradigm is that they may well be the high redshift counterparts and progenitors of local massive E and SO galaxies. Their reliable classification and study, especially at intermediate redshift ($z \\sim 1$) can provide crucial constraints on the evolution of massive, starburst,  dusty and/or ultraluminous infrared galaxies (ULIRGs) known to exist at higher redshift \\eg\\ BzKs \\citep{daddi-04}, BX/BMs \\citep{reddy-05}, Distant Red Galaxies  \\citep[DRGs]{franx-03, pap-06a, conselice-06}, submillimeter and IR-Luminous Lyman Break Galaxies \\citep[ILLBGs]{huang-05, rig-06}.      In the restframe near-IR, old stellar populations show a turndown at wavelengths longer than  the restframe $1.6 \\micron$ `bump', while dusty starburst populations show emission from small hot dust grains. AGN-dominated sources display a power-law spectral energy distribution. In \\cite{wilson-04} we showed how data from \\emph{Spitzer} could  begin  to help distinguish among different ERO populations. In this paper, we extend our ERO study to take advantage of the rich panchromatic dataset available from the All-wavelength Extended Groth strip International Survey (AEGIS). All magnitudes used in this letter are AB, unless otherwise specified. ",
        "conclusions": "\\label{sec:conc}    In this paper, we utilized the AEGIS dataset to explore the nature of 87 EROs with confirmed spectroscopic redshifts. IRAC imaging facilitated dividing this sample into three distinct classes using a simple [3.6] -[8.0] color selection technique. We presented SEDs and high-resolution ACS images supporting our (early-type, dusty starburst or power-law) classification.  We showed that the three classes of ERO and their redshift distribution were broadly consistent with what one would expect from a simple color-color diagram (Figure~\\ref{fig:R14}). We found 53 early-types, 32 dusty starbursts/LIRGs, and two obvious AGN.  Both of these AGN would be classified as ULIRGs.  The selection of this particular sample of EROs was subject to the spectroscopic biases  discussed in  \\S~\\ref{sec:zdist}. We next plan to extend this study to the  several thousand IRAC-selected EROs in the AEGIS \\emph{without} spectroscopic redshifts."
    },
    "0608/astro-ph0608671_arXiv.txt": {
        "abstract": "We argue on the basis of astrometric and spectroscopic data that the ultramassive white dwarf GD50 is associated with the star formation event that created the Pleiades and is possibly a former member of this cluster. Its cooling age ($\\sim$60Myrs) is consistent with it having evolved essentially as a single star from a progenitor with a mass M$>$6M$ _{\\odot}$ so we find no need to invoke a white dwarf-white dwarf binary merger scenario to account for its existence. This result may represent the first direct observational evidence that single star evolution can produce white dwarfs with M$>$1.1M$_{\\odot}$, as predicted by some stellar evolutionary theories. On the basis of its tangential velocity we also provisionally identify the ultramassive (M$\\sim$1.2M$_{\\odot}$) white dwarf PG0136+251 as being related to the Pleiades. These findings may help to alleviate the difficulties in reconciling the observed number of hot nearby ultramassive white dwarfs with the smaller number predicted by binary evolution models under the assumption that they are the products of white dwarf mergers. ",
        "introduction": "GD50 (WD0346-011) is a well studied nearby hot H-rich white dwarf with a mass of M$>$1.1M$_{\\odot}$ (e.g. Giclas, Burnham \\& Thomas 1965, Bergeron et al. 1991; Table 1). Extrapolation of recent determinations of the initial mass-final mass relation (IFMR; e.g. Dobbie et al. 2006) suggest that if it evolved via single star evolution then its progenitor likely had a mass M$\\simgreat$6M$_{\\odot}$. However, it appears to reside well away from any known star forming region, stellar association or young ($\\tau$$\\simless$200Myrs) open star cluster. This is somewhat surprising as current observational evidence indicates that massive stars are born predominantly within rich stellar associations or star clusters (e.g. Lada \\& Lada 2003, de Wit et al. 2004).  More generally, the observed number of massive white dwarfs, M$>$0.8M$_{\\odot}$, may be too large for them all to have evolved as single stars from OB type progenitors, for reasonable assumptions about the Galactic star formation history and the form of the IFMR. Accordingly, it has been suggested that GD50 and a large proportion of the massive white dwarfs may have formed through binary evolution ($\\approx$80\\%; Liebert et al. 2005a). For example, these objects may be born from the merger of two white dwarfs (either He+CO or CO+CO) each with a mass closer to the canonical value of 0.6M$_{\\odot}$ (e.g. Segretain, Chabrier \\& Mochkovitch 1997, Mochkovitch 1993). The detection in an Extreme Ultraviolet Explorer spectrum of GD50 of what appeared to be rotationally  broadened (1000kms$^{-1}$) absorption lines of photospheric helium lent credance to this hypothesis (Vennes et al. 1996) since it might be expected that these merger products are rapid rotators with unusual atmospheric compositions. However, a more recent and detailed high resolution study of the spectral energy distribution of this star centered on the H-$\\alpha$ line failed to corroborate this result and has set an upper limit on the rotational velocity of $v \\sin i$$\\simless$35kms$^{-1}$ (Vennes 1999). Furthermore, an extreme-ultraviolet spectroscopic survey of a further eight ultramassive white dwarfs failed to reveal compelling evidence for the presence of helium in the atmospheres of any of these objects (Dupuis, Vennes \\& Chayer 2002). Thus the origin of GD50 and ultramassive white dwarfs in general is still an open question.  In this short communication we argue, based on astrometric and spectroscopic data and our recent re-evaluation of the form of the IFMR, that GD50 is associated with the Pleiades open cluster and is possibly a former member with properties consistent with having evolved essentially as a single star. We briefly discuss the implications of this result for stellar evolution theory. \\begin{table*} \\begin{minipage}{170mm} \\begin{center} \\caption{Summary details of GD50 (WD0346$-$011) including coordinates, apparent visual magnitude (Marsh et al. 1997), distance, proper motion, radial velocity and heliocentric space velocity.} \\label{sum1} \\begin{tabular}{llcccccccc} \\hline \\multicolumn{1}{|c|}{RA}    &  \\multicolumn{1}{|c|}{Dec}   & V & d &  $\\mu_{\\alpha}\\cos\\delta$ & $\\mu_{\\delta}$ & RV & U & V & W \\\\ \\multicolumn{2}{|c|}{J2000.0} & & pc &\\multicolumn{2}{|c|}{mas yr$^{-1}$} & kms$^{-1}$ & \\multicolumn{3}{|c|}{kms$^{-1}$} \\\\ \\hline 03 48 50.20  & -00 58 31.2 & 14.04$\\pm$0.02  & 31.0$\\pm$1.7 & +86.6$\\pm$6.8 & -164.4$\\pm$5.0 & +13.8$\\pm$12.2 &-3.8$\\pm$9.2 & -28.0$\\pm$2.2 & -11.8$\\pm$7.9 \\\\% & +75.1$\\pm$16.4 & -155.9$\\pm$18.1 \\hline \\end{tabular} \\end{center} \\end{minipage} \\end{table*}  \\section[]{The heliocentric space velocity of GD50}  We have obtained estimates of the proper motion of GD50 from the SuperCOSMOS Sky Survey database ($\\mu_{\\alpha}\\cos\\delta$=+75.1$\\pm$16.4mas~yr$^{-1}$, $\\mu_{\\delta}$=$-$155.9$\\pm$18.1mas~yr$^{-1}$; Hambly et al. 2001), the Lick Northern Proper Motion Program ($\\mu_{\\alpha}\\cos\\delta$=+89.4$\\pm $5.5mas~yr$^{-1}$, $\\mu_{\\delta}$=$-$165.6.9$\\pm$5.5mas~yr$^{-1}$; Klemola et al. 1987) and from our own measurements of the Palomar Sky Sky (E932, 1953/12/31) and UK Schmidt Telescope (9597, 1986/09/22) survey plates ($\\mu_{\\alpha}\\cos\\delta$$\\approx$+86$\\pm$4mas~yr$^{-1}$, $\\mu_{\\delta}$$\\approx$$-$164$ \\pm6$mas~yr$^{-1}$). These values are consistent with each other within their respective uncertainties: we determine a weighted mean proper motion for GD50 of $\\mu_{\\alpha}\\cos\\delta$=+86.6$\\pm$6.8mas~yr$^{-1} $, $\\mu_{\\delta}$=$-$164.4$\\pm$5.0mas~yr$^{-1}$, where the final uncertainties have been estimated from the scatter of the individual measurements.  From a high S/N UVES spectrum of the white dwarf taken in the course of the SPY programme (Napiwotzki et al. 2003) we have determined the weighted mean redshift of the H-$\\alpha$ and H-$\\beta$ line cores to be +176.0$\\pm$4.3kms$^{-1}$, in excellent agreement with the earlier measurement of Vennes (1999). Additionally, by determining the weighted means of the robust estimates found in the refereed literature of the last 15 years and new measurements obtained from our independent analysis of the FORS1 spectrum of GD50 detailed in Aznar-Cuadrado et al. (2004) we have estimated the effective temperature and surface gravity of this object to be T$_{\\rm eff}$=41550$\\pm$720K and log g=9.15$\\pm$0.05, respectively. As listed in Table 2, the errors in the individual determinations of effective temperature and surface gravity have been assumed to be at least 2.3\\% and 0.07~dex, respectively (e.g. see Napiwotzki et al. 1999), while the scatter of these measurements has been used to estimate the final uncertainties in the mean parameters. The measurements of Barstow et al. (1993) were rejected from the calculation as they were found to lie $>$3$\\sigma$ in both effective temperature and surface gravity from an initial estimate of the means.  Using CO core, thick (thin) H-layer, white dwarf evolutionary models (e.g. Fontaine, Brassard \\& Bergeron 2001) we infer the mass, the radius and the cooling time of GD50 to be 1.264$\\pm $0.017 (1.253$\\pm$0.018)M$_{\\odot}$, 0.00495$\\pm$0.00034 (0.00493$\\pm$0.00034)R$_{\\odot}$ and 61$\\pm$6 (58$\\pm$6)Myrs, respectively. We note that the specific choice of core composition, CO or ONe, is believed to make little difference these values (e.g. van Kerkwijk \\& Kulkarni 1999, Hamada \\& Salpeter 1961). We determine the gravitational contribution to the line core redshift to be +162.2$\\pm$11.4 (+161.4$\\pm$11.4)kms$^{-1}$. Thus we estimate the radial velocity of GD50 to be +13.8$\\pm$12.2 (+14.6$\\pm$12.2)kms$^{-1}$. Based on a modern grid of white dwarf synthetic photometry (Holberg \\& Bergeron 2006; Bergeron et al. 1995), the above effective temperature and surface gravity estimates and the V magnitude from Marsh et al. (1997) we determine the distance of GD50 to be 31.0$\\pm$1.7 (30.9$\\pm$1.7)pc. Following the prescription given by Johnson \\& Soderblom (1987) we calculate the heliocentric space velocity of GD50 to be U=$-$3.8$\\pm$9.2 ($-$4.4$\\pm$9.2)kms$^{-1}$, V=$-$28.0$\\pm$2.2 ($ -$28.1$\\pm$2.2)kms$^{-1}$ and W=$-$11.8$\\pm$7.9 ($-$12.3$\\pm$7.9)kms$^{-1}$ (see Table 1).  The recent kinematical study performed by Pauli et al. (2006) of 398 field white dwarfs drawn from the catalogue of McCook \\& Sion (2000) reveals that $\\sim$75\\% have V$>$-25kms$^{-1}$. Additionally, the error bounds on the U,V and W velocity components of only 3 of these stars overlap with the range defined by the above values and their associated uncertainties (assuming a solar velocity with respect to the local standard of rest of U=10.00kms$^{-1}$, V=5.25kms$^{-1}$ and W=7.17kms$^{-1}$; Dehnen \\& Binney 1998). Thus we conclude that GD50 lies in a region of velocity space which is not heavily populated by white dwarfs. However, we note that the space velocity of this object is strikingly close to that of Pleiades open star cluster and also the AB Dor moving group (Zuckermann et al. 2004). The Pleiades is the closest rich young ($\\tau$$\\approx$125Myrs) open cluster to GD50. We determine its space motion to be U=$-$6.5$\\pm$0.5kms$^{-1}$, V=$-$27.8$\\pm$1.1kms$^{-1}$ and W=$-$14.6$\\pm $0.5kms$^{-1}$, for a distance of 134$\\pm$5pc (the weighted mean of recent ground based estimates e.g. Percival et al. 2005, Southworth et al. 2005, Pan et al. 2004, Munari et al. 2004, Zwahlen et al. 2004, Gatewood et al. 2000), a radial velocity of +5.7$\\pm$0.5kms$^{-1}$ and a proper motion of $\\mu_{\\alpha}\\cos\\delta$=+19.15.1$\\pm$0.23mas yr$^{-1}$,$\\mu_{\\delta}$=-45.72$\\pm$0.18mas yr $^{-1}$ (Robichon et al. 1999). Luhman et al. (2005) estimate the space motion of the AB Dor moving group, within which the Sun resides, to be U=$-$7.7$\\pm$0.4kms$^{-1}$, V=$-$26.0$\\pm$0.4kms $^{-1}$ and W=$-$13.6$\\pm$0.3kms$^{-1}$. ",
        "conclusions": "\\subsection{The Local Association}  The Pleiades and the AB Dor moving group are part of the Local Association, a large scale ``supercluster'' structure consisting of a gravitationally unbound collection of open clusters, stellar associations and moving groups with comparable space motions (e.g. Sco-Cen, IC2602, Per OB3, Cas-Tau and NGC2516; Eggen 1992). While this could be taken as evidence of a common ancestry (e.g. Weidemann et al. 1992), the age of the members of this supercluster is estimated to span the relatively broad range of $\\sim$few Myrs to $\\sim$few 100Myrs (e.g. Eggen 1992, Asiain et al. 1999). Chereul et al. (1998) argue that the Local Association may just be the result of the chance superposition in phase space of several smaller stellar streams, each associated with a star formation event which if sufficiently massive manufactured a gravitationally bound star cluster. Indeed, Hipparcos observations of open clusters within the Local Association reveal that at a level of $<$10kms$^{-1}$ their space velocities are quite distinct e.g.  Pleiades, U=$-$6.5$\\pm$0.5kms$^{-1}$, V=$-$27.8$\\pm$1.1kms$ ^{-1}$ and W=$-$14.6$\\pm$0.5kms$^{-1}$, Per OB3 ($\\alpha$-Per), U=$-$15.3$\\pm$0.7kms$^{-1}$, V=$ -$25.8$\\pm$1.0kms$^{-1}$, W=$-$7.9$\\pm$0.4kms$^{-1}$) and NGC2516, U=$-$17.4$\\pm$1.4kms$^{-1}$, V=$ -$23.7$\\pm$0.4kms$^{-1}$, W=$-$3.9$\\pm$0.4kms$^{-1}$ (Robichon et al. 1999).  Other recent extensive analyses of the velocity and age distributions of stars in the solar neighbourhood confirm the existence of this significant substructure within the Local Association. For example, based on a study of nearby early-type stars, Asiain et al. (1999) resolve the supercluster into 4 distinct components in U,V,$\\tau$ space, each with a lower velocity dispersion and a more restricted age range than the Local Association as a whole. Moving group B1, with U=$-$4.5$\\pm$4.7kms$^{-1}$, V=$-$20.1$\\pm$3.3kms$^{-1}$ and W=$-$5.5$\\pm$1.9kms$^{-1}$, consists of a very young population (20$\\pm $10Myrs) and is attributable to the Scorpio-Centaurus OB association, while moving group B4, with U=$-$8.7$\\pm$4.8kms$^{-1}$, V=$-$26.4$\\pm$3.3kms$^{-1}$ and W=$-$8.5$\\pm$4.7kms$^{-1}$, is substantially older (150$\\pm$50Myrs). Given the similarities between the space velocities and ages of this moving group and the Pleiades open cluster, it is likely the two are related (Asiain et al. 1999).   \\begin{table} \\begin{minipage}{84mm} \\begin{center} \\caption{Temperature and gravity determinations of GD50 in the refereed literature of the last 15 years used in deriving the means. The values of Barstow et al. (1993) were excluded from the calculation as they lie $>$3$\\sigma$ from an initial estimate of the means. We assume a minimum uncertainty in all temperature and gravity determinations of 2.3\\% and 0.07dex, respectively.} \\label{sum1} \\begin{tabular}{ccc} \\hline T$_{\\rm eff}$(K)   & log g   & Reference \\\\ \\hline 40538$\\pm$932 & 9.22$\\pm$0.08 &     1 \\\\ 38088$\\pm$876 & 8.87$\\pm$0.07 &     2 \\\\ 43102$\\pm$1982 & 9.09$\\pm$0.09 &    3 \\\\ 43200$\\pm$994  & 9.21$\\pm$0.07 &    4 \\\\ 39508$\\pm$908  & 9.07$\\pm$0.07 &    5 \\\\ 43170$\\pm$993  & 9.08$\\pm$0.07 &    6 \\\\ 41480$\\pm$954  & 9.19$\\pm$0.07 &    7 \\\\ \\hline $<$41550$\\pm$720$>$ &  $<$9.15$\\pm$0.05$>$ & \\\\ \\hline  \\end{tabular} \\end{center} (1) Bergeron et al. (1991), (2) Barstow et al. (1993), (3) Bragaglia et al. (1995), (4) Vennes et al. (1997), (5) Napiwotzki et al. (1999), (6) Aznar-Cuadrado et al. (2004) and (7) This work \\end{minipage} \\end{table}  Luhman et al. (2005) and Luhman \\& Potter (2006) have recently shown that the stellar sequences of the AB Dor moving group (including AB Dor Ba and Bb) and the Pleiades open cluster are coincident in M$_{\\rm K}$,V$-$K and M$_{\\rm K}$,J$-$K colour-magnitude diagrams. Thus not only do the AB Dor moving group members and the Pleiades have very similar space velocities, there is compelling evidence that the former also have an age in the range 100-125Myrs. It is argued by these authors that the AB Dor moving group is thus likely to be the remnants of an OB and T association related to the formation and evolution of the Pleiades open cluster, a more mature version of the relationship which exists today between the Per OB3 open cluster and the Cassiopeia-Taurus Association (de Zeeuw et al. 1999). In this framework, presumably the bulk of the B4 moving group is also part of this unbound remnant.  \\subsection{GD50 and the Pleiades star formation event}  We find, in the context of the substantial velocity substructure of the Local Association, the kinematics of GD50 are most closely matched to those of the B4 moving group/Pleiades open cluster stream. To further examine the possibility that GD50 is related to this particular episode of star birth we have examined it's location in initial mass-final mass space under the assumption that it is coeval with the Pleiades (125$\\pm $25Myrs; Ferrario et al. 2005). Noting that the determination of initial mass is extremely sensitive to the adopted age of the progenitor population (see Dobbie et al. 2006 for details), the proximity of GD50 to an extrapolation of the linear fit to the initial mass-final mass data for 27 white dwarfs in open clusters and the Sirius binary system (Dobbie et al. 2006, Liebert et al. 2005b, Williams, Bolte \\& Koester 2005, Claver et al. 2001, Koester \\& Reimers 1996), indicates that it's cooling age is entirely consistent with it being associated with this star formation event and having evolved essentially as a single star (see Figure 1).  To get further insights we calculated the Galactic orbits of GD50 and the Pleiades back in time. We used the code {\\sc orbit6} of Odenkirchen \\& Brosche (1992) for our calculations. {\\sc orbit6} computes the trajectory of a star in the Galactic potential for a given set of initial coordinates and space velocities. Please refer to Pauli et al.\\ (2003) for further details. For every time step we calculated the distance between the orbits of GD50 and the Pleiades. Errors in the estimated distance and space velocities were propagated by means of a Monte Carlo simulation. The results of this modelling are consistent with GD50 having formed within the Pleiades and having been subsequently ejected. Indeed, the escape velocity from the center of the present day cluster is only $\\sim$2kms$^{-1}$ (Dehnen, priv. comm) and while the Pleiades is relatively young and one white dwarf member has already been firmly identified (LB1497; Luyten \\& Herbig 1960), it is estimated, based around an extrapolation of the present day mass function, that a small number of other massive members could by now have evolved to this configuration (e.g. Williams 2004). As dynamical evolution would have likely led the relatively massive progenitor of this white dwarf to settle towards the most densely populated central regions of the cluster, it might be that, after formation, GD50 gained sufficient kinetic energy through an interaction with another star or a binary system to allow it to escape the Pleiades. Alternatively, some massive white dwarfs may receive a small recoil velocity kick during the final stages of their (super)-AGB evolution due to low level point-asymmetries in the outflowing material (e.g. Fellhauer et al. 2003). We note that GD50 need only have been moving at a mean velocity, post ejection, of $\\sim$2kms$^{-1}$ with respect to the cluster for $\\sim$55Myrs for the two to be separated by $>$100pc as observed today.  A further possibility, which is less favoured but not ruled out by our Galactic orbit modelling, is that GD50 or more particularly it's progenitor star, may have originated in an unbound OB association related to formation of the Pleiades open cluster. While the AB Dor moving group has only $\\sim$30 known members with spectral types from F5-M6 (Zuckermann et al. 2004, Luhman \\& Potter 2006), the B4 moving group is known to contain at least $\\sim$50 objects of spectral types A and B (Asiain et al. 1999). Furthermore, detailed simulations of the early evolution of a rich stellar aggregate like the nascent Pleiades, suggest that when the ignition of the OB stars drives away the mass dominating primordial gas, as many as 2/3 of the initial constituents may become part of an unbound and expanding association. Over a period comparable to the age of the B4 stream this structure can disperse over scales well in excess of 100pc (Kroupa, Aarseth \\&  Hurley 2001).  \\begin{figure} \\vspace{185pt} \\special{psfile=gd50.eps voffset=5 hscale=88 vscale=95 hoffset=1} \\caption{The white dwarf members of the Hyades, Praesepe, M35, NGC2516 and Pleiades open clusters and the Sirius binary system in initial mass-final mass space (see text for references). A linear fit to the data derived by Dobbie et al. (2006) and the relation of Weidemann (2000; dotted line) are overplotted as solid and dotted lines, respectively. The location of GD50, derived assuming it to be coeval with the Pleiades, is consistent with it having evolved essentially as a single star (open square). The approximate location of PG0135+251 is also shown (filled square).} \\end{figure}   \\subsection{Stellar evolution}  Although the available data do not allow us to pin down the details of the history of GD50, as a whole, the evidence presented here provides a relatively compelling argument that it is associated with the star formation event that created the Pleiades. Clearly there is no need to invoke a binary white dwarf merger scenario to account for its evolution, if as supported by its location in Figure 1, it has evolved essentially as a single star. Liebert et al. (2005a) have recently noted the apparent projected spatial coincidence between many of the hot massive white dwarfs detected in the ROSAT WFC and EUVE surveys and Gould's Belt, a rich band of OB stars tilted at 18$^{\\circ}$ with respect to the Galactic Plane which represents a ring of nearby recent star formation. This has led them to suggest that significantly more than 20\\% of these objects may be the progeny of single stars. Furthermore, based on their recent determination of the form of the IFMR, Ferrario et al. (2005) predict that $\\sim$28\\% of objects in the white dwarf mass distribution have M$\\simgreat$0.8M$_{\\odot}$.  Indeed, the formation of ultramassive white dwarfs from single stars is predicted by the stellar evolutionary modelling of a number of research groups. For example, Garcia-Berro et al. (1997) show that once the helium exhausted cores of some stars exceed $\\sim$1.1M$_{\\odot}$, a series of carbon-burning shell flashes can lead to a super-AGB phase of evolution and ultimately to the formation of degenerate ONeMg cores with masses in the range $\\sim$1.1-1.4M$_{\\odot}$. Alternatively, if a star has significant rotational angular momentum, then the pressure lifting effect of this can allow the CO core, which will ultimately become the white dwarf, to grow to M$>$1.1M$_{\\odot}$ (e.g. Dominguez et al. 1997). We believe the present result represents the first compelling observational evidence directly linking ultramassive white dwarfs and single star evolution.  We also note that PG0136+251, with T$_{\\rm eff}$=39640K, log g=8.99 and M=1.20$_{\\odot}$ (Liebert et al. 2005a), has, as listed in the USNO-B1.0 catalogue, a proper motion of $\\mu_{\\alpha}\\cos\\delta$=+52$\\pm$2mas yr$^{-1}$, $\\mu_{\\delta}$=$-$46$\\pm$1mas yr$^{-1}$ and is thus moving on a bearing which points within 10$^{\\circ}$ of the convergent point of the Pleiades (e.g. Makarov \\& Robichon 2001). Additionally, the distance calculated using the moving clusters method (Equation 1), which assumes the white dwarf to be affiliated kinematically to this cluster,  \\begin{equation} d_{\\rm mc}=v \\sin \\lambda/4.74\\mu \\approx 97{\\rm pc} \\end{equation}  where $v$ is the space velocity of the Pleiades ($\\sim$32kms$^{-1}$; Luhman et al. 2005), $\\lambda$ the angular separation between the white dwarf and the convergent point ($\\sim$95$^{\\circ}$) and $\\mu$ the tangential motion ($\\sim$0.07arcsec~yr$^{-1}$), is remarkably close to the spectrophotometrically derived value of $\\approx$93pc. The mass and the cooling time of this white dwarf can also be considered consistent with it having evolved essentially as a single star born at about the same time as the Pleiades open cluster (see Figure 1). A firm conclusion regarding the origins of PG0136+251 must, however, await a radial velocity measurement for this object.  Both GD50 and PG0136+251 are sometimes touted as possible examples of white dwarf binary mergers (e.g. Bergeron et al. 1991, Mochkovitch 1993, Segretain, Chabrier \\& Mochkovitch 1997). However, the observed number of nearby hot ultramassive white dwarfs is difficult to reconcile with the population predicted by binary evolution models which adopt a plausible Galactic white dwarf merger rate (10$^{-2}$--10$^{-3}$ yr $^{-1}$), under the assumption that these objects are the products of coalesced double degenerate systems. For example, for a Galactic merger rate of 10$^{-2}$yr$^{-1}$ and a cooling time scale of the order 20Myrs, Segretain, Chabrier \\& Mochkovitch (1997) estimate that the average distance to the closest hot ultramassive white dwarf merger product is $\\sim$70pc. The probability of there being 3 within 40pc (e.g. GD50, PG1658+441; Green, Schmidt \\& Liebert 1986, and  RE0317-854; Barstow et al. 1995) is a mere $\\sim$0.1\\%. However, this difficulty must be alleviated to some extent if a significant proportion of these hot ultramassive white dwarfs have instead formed via single star evolution (Segretain, Chabrier \\& Mochkovitch 1997).  A detailed study of the heliocentric space velocities of a large sample of hot young massive white dwarfs is clearly warranted. This would address additional questions regarding the end points of the evolution of intermediate mass stars, Type Ia supernovae progenitors and more generally the recent history of the local star formation rate."
    },
    "0608/astro-ph0608392_arXiv.txt": {
        "abstract": " ",
        "introduction": "Cataclysmic variables (CVs) are semi-detached binaries in which a red (K or M) main-sequence, or slightly evolved, low-mass star loses mass to a white-dwarf primary through overfilling its Roche lobe. They can be divided into two groups, the magnetic CVs (mCVs) and the non-magnetic CVs, according to whether the white-dwarf has a magnetic field that is strong enough to affect the accretion dynamics. Semi-detached binaries frequently possess an accretion disc, where gravitational energy is converted to kinetic energy and then to radiation. With mCVs the magnetic-field strength of the white-dwarf is high enough to disrupt the accretion disc at the magnetosphere. The white-dwarf's magnetosphere will determine the inner radius of the accretion disc and the tidal force exerted by the secondary star will truncate the outer region of the disc. If the magnetosphere is large enough, no disc will form.  MCVs without an accretion disc are known as polars (or AM Herculis systems). In a polar the white-dwarf magnetic field is sufficiently strong ($\\sim 10$~MG) to also lock the white-dwarf and its companion star into synchronous (or almost synchronous)  rotation with the binary orbit. The accretion flow near the white-dwarf is channelled by the field and a shock is often formed near the base of the the accretion column before the accreting gas settles onto the white-dwarf atmosphere. The shock heats the accreting gas, leading to the emission of bremsstrahlung X-rays that can subsequently photoionize the cooler gas upstream. For systems consisting of a white-dwarf with a moderately strong magnetic field ($\\sim 1$~MG), the white-dwarf spin and the orbital motion are not synchronous, although the secondary may still be tidally locked and rotate with the orbit. These moderately magnetic systems are known as intermediate polars (IPs). Most IPs possess an accretion disc, but  the inner region of the disc is truncated by the magnetic field, and the flow is channelled in the white-dwarf magnetosphere, similar to polars. A shock can be formed in the accretion column, and the X-rays from the post-shock region can photoionize both the pre-shock gases in the accretion column and the cool gas in the accretion disc. It is worth noting that some IPs may not have a fully developed accretion disc. They probably have a thin accreting annulus, which couples loosely to the magnetosphere of the white-dwarf.  Non-magnetic CVs form a heterogeneous class of objects, which include nova, dwarf nova and nova-like systems, primarily on the basis of their light curve behaviour (although strict classification is sometimes not possible, e.g. some novae are also polars). They are asynchronously rotating systems, and they have a fully developed accretion disc with an inner edge extending to the white-dwarf surface. A substantial fraction of the accretion energy is dissipated in a transition layer between disc and the white-dwarf. The strong UV photons emitted from the transition region can be a main cause of photoionization in the accreting gas.  CVs  are strong emission line objects, particularly in quiescence. They all show a clear H Balmer emission sequences. He I, He II and high ionization lines of other species are also prominent in the optical spectra of many systems. The high ionization lines of the CVs are believed to be emission from photoionized gas in the accretion disc (for IPs and non-magnetic CVs), the accretion stream, or irradiatively heated surface of the companion star. Some of the lines may originate from hot spots in the system, for example in regions where the accretion stream impacts onto the accretion disc and where the accretion flow begins to couple with the magnetic field. In these cases the gas is often heated by shocks that are formed as consequence of an abrupt change in the dynamics of the accretion flow. (For reviews of CVs and related accretion physics, see \\citet{warner}, \\citet{cropper_90}, \\citet {wu_00}, \\citet{wu_03}.)  The lines from photoionized gases and shock heat\\-ed gases have different properties. An example is that of the relative strengths of the H Balmer emission lines. For instance, simple photoionization-re\\-com\\-bin\\-a\\-tion models \\citep{osterbrock} generally predict ratios of $\\approx 3.5$ for H$\\alpha$/H$\\beta$; yet observations show H$\\alpha$/H$\\beta$ which deviate significantly from these model predictions. Such deviation implies that the line formation process is more complex than simple photoionization and subsequent recombination.  Here we analyse the H Balmer emission from CVs systematically using multivariate statistical methods. We consider the systems in \\cite{williams_gw83} as the sample for our analysis. We search for correlations between the emission line ratios and CV classification and review the line formation mechanisms that may be responsible for the correlations. The paper is organized as follows. In \\S \\ref{S2} we review the multivariate statistical methods used to analyze the H Balmer emission data. These methods are discriminant function and principal component analysis. In \\S \\ref{S3} we present the data in a four dimensional scatterplot space and give the results of the discriminant function and principal component analysis in that geometric space. We also review models of Balmer emission relevant to CVs. In section \\S \\ref{S4} we discuss possible physical causes, based on existing models, of the statistical behaviour of the data. Two significant directions of variation are found that correlate with the binary orbital period and the equivalent width of the h$\\beta$ line. Our conclusion is given in \\S \\ref{S5}.  \\hspace*{2cm} ",
        "conclusions": "\\label{S5}  Our analysis of hydrogen Balmer lines in CVs shows that the source of variation seen in the principal components analysis are correlated with the binary orbital period, and the source of variation seen in the discriminant function analysis are correlated with the equivalent width of the H$\\beta$ line. Comparing models to the data  produces one positive result. The ratios of hydrogen Balmer lines predicted by the photoionized model match the observed line ratios of the T CrB systems. In particular, the emission from RS Oph is consistent with photoionization from a $\\sim$5000 K optically thin (Case A) nebula. These systems, with a massive white dwarf and likely progenitors of type Ia supernovae, are hypothesized to be embedded in a nebula produced by the frequent novae outbursts \\citep{warner}. Our analysis supports this scenario, as the observed Balmer emissions could originate from a photoionized nebula. Otherwise, the theoretical models considered in this work do not produce the wide range of line ratios along the first principal component direction of the CV dataset. In particular, if we assume that the mass-transfer rate and the period are related by equation (\\ref{mdot}), the model of LTE emission from isothermal medium predicts an opposite correlation between the observed first principal component direction and mass transfer rate than what is observed.  Our analysis showed a wide variance of discriminant function values and a wide distribution of line ratios along the first principal component direction for dwarf novae. In light of the known thermal cycling in the accretion disc, regardless of whether the lines are emitted directly from the disc or are from reprocessed radiation, both the Balmer line ratios and the line equivalent widths would be expected to vary with the dwarf nova outburst cycle. This may be a cause of the wide variance in the discriminant function values and of the wide distribution along the first principal component direction. As far as we are aware, systematic observations of hydrogen Balmer emission lines throughout a dwarf nova cycle have not been reported.  Our analysis here shows that the range of Bal\\-mer ratios seen along the first principle component direction, including ``inverted'' Balmer decrements, can not be reproduced by the conventional models considered here. The more or less continuous variation of the line ratios from Balmer decrements to Balmer ``increments'' for all CVs as a group prevents any clustering of the data that clearly points to a particular single underlying physical process. We therefore conclude that either the emission from different CV sub-classes have different locations (and hence line formation processes) within the system, or  there are multiple regions for the hydrogen Balmer lines within a CV. The good matching between the data of T CrB systems and the photo-ionization nebular model supports the former. However, the latter is also likely, given that the UV radiation from the accreting white dwarf can ionize the cooler material in the outer accretion disc, the accretion stream and the companion star."
    },
    "0608/hep-ph0608344_arXiv.txt": {
        "abstract": "Considering gravitino dark matter scenarios, we obtain the full gauge-invariant result for the relic density of thermally produced gravitinos to leading order in the Standard Model gauge couplings. For the temperatures required by thermal leptogenesis, we find gaugino mass bounds which will be probed at future colliders. We show that a conceivable determination of the gravitino mass will allow for a unique test of the viability of thermal leptogenesis in the laboratory. ",
        "introduction": "The smallness of the neutrino masses can be understood naturally in terms of the see-saw mechanism~\\cite{Minkowski:1977sc} once the Standard Model is extended with right-handed neutrinos which have heavy Majorana masses and only Yukawa couplings. For a reheating temperature after inflation, $T_\\Reheating$, which is larger or not much smaller than the masses of the heavy neutrinos, these particles are produced in thermal reactions in the early Universe. The CP-violating out-of-equilibrium decays of the heavy neutrinos generate a lepton asymmetry that is converted into a baryon asymmetry by sphaleron processes~\\cite{Fukugita:1986hr}. This mechanism, known as thermal leptogenesis, can explain the cosmic baryon asymmetry for $T_{\\Reheating}\\gsim 3\\times 10^9\\,\\GeV$~\\cite{Buchmuller:2004nz}.  One will face severe cosmological constraints on $T_\\Reheating$ if supersymmetry (SUSY) is discovered.  An unavoidable implication of SUSY theories including gravity is the existence of the gravitino $\\gravitino$ which is the gauge field of local SUSY transformations. As the spin-3/2 superpartner of the graviton, the gravitino is an extremely weakly interacting particle with couplings suppressed by inverse powers of the (reduced) Planck scale $\\MPl=2.4\\times 10^{18}\\,\\GeV$~\\cite{Wess:1992cp}. In the course of spontaneous SUSY breaking, the gravitino acquires a mass $\\mgr$ and the couplings of the spin-1/2 goldstino which become dominant for small $\\mgr$. Depending on the SUSY breaking scheme, $\\mgr$ can range from the eV scale up to scales beyond the TeV region~\\cite{Martin:1997ns}. Gravitinos can be produced efficiently in the hot primordial plasma. Because of their extremely weak interactions, unstable gravitinos with $\\mgr \\lsim 5~\\TeV$ have long lifetimes, $\\tau_{\\gravitino} \\gsim 100~\\sec$, and decay after big bang nu\\-cleo\\-synthesis (BBN).  The associated decay products affect the abundances of the primordial light elements. Demanding that the successful BBN predictions are preserved, bounds on the abundance of gravitinos before their decay can be derived which imply $T_{\\Reheating}\\lsim 10^8\\,\\GeV$ for $\\mgr \\lsim 5~\\TeV$~\\cite{Kohri:2005wn}. Thus, the temperatures needed for thermal leptogenesis are excluded.  We therefore consider SUSY scenarios in which a gravitino with $\\mgr \\gsim 10~\\GeV$ is the lightest supersymmetric particle (LSP) and stable due to R-parity conservation. These scenarios are particularly attractive for two reasons: (i) the gravitino LSP is a compelling dark matter candidate and (ii) thermal leptogenesis can still be a viable explanation of the baryon asymmetry~\\cite{Bolz:1998ek}. ",
        "conclusions": "We provide the full SU(3)$_\\Color\\times$SU(2)$_\\Weak\\times$U(1)$_\\Hypercharge$ result for the relic density of thermally produced gravitino LSPs to leading order in the gauge couplings. Our result is obtained in a consistent gauge-invariant finite-temperature calculation and thus independent of arbitrary cutoffs. With this result, new gravitino and gaugino mass bounds emerge as a prediction of thermal leptogenesis. If supersymmetry is realized in Nature, these bounds will be accessible at the LHC and the ILC. In particular, with a charged slepton NLSP, there will be the exciting possibility to identify the gravitino as the LSP and to measure its mass. Confronting the measured gravitino mass with the predicted bounds will then allow for a unique test of the viability of thermal leptogenesis in the laboratory.  \\medskip"
    },
    "0608/astro-ph0608300_arXiv.txt": {
        "abstract": "We present images and slitless spectra which were obtained in {\\it HST} surveys of Planetary Nebulae (PNe) in both the Large and Small Magellanic Clouds, using the Space Telescope Imaging Spectrograph. These new data on 59 PNe (54 in the LMC and five in the SMC) permit us to determine the nebular dimensions and morphology in the monochromatic light of several emission lines: H$\\alpha$, [\\ion{N}{2}] $\\lambda$6583 and [\\ion{O}{3}] $\\lambda$5007, plus others of varying ionization, including [\\ion{O}{1}], \\ion{He}{1}, and [\\ion{S}{2}]. We describe the nebular morphology and related features in detail. This survey, when combined with similar data from our prior {\\it HST} programs and other archived PN images, brings the total of nebulae imaged with {\\it HST} to 114 in the LMC and 35 in the SMC. We describe various basic properties for the sample, including sizes, morphologies, densities, and completeness. Trends in [\\ion{O}{3}] $\\lambda$5007 flux, surface brightness, and electron density with physical radius suggest that many nebulae, particularly those with bipolar morphology, may be optically thick even at large size. Bipolars also show the most extreme values of [\\ion{N}{2}]/H$\\alpha$ flux ratios, which is a rough indicator N enrichment. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608136_arXiv.txt": {
        "abstract": "% The southern Pierre Auger Observatory, presently under construction in Malarg\\\"ue, Argentina, is nearing completion. The instrument is designed to measure extensive air-showers with energies ranging from $10^{18}$-$10^{20}$~eV and beyond.  It combines two complementary observation techniques; the detection of particles at ground and the coincident observation of associated fluorescence light generated in the atmosphere above the ground.  This is being realized by employing an array of 1600 water Cherenkov detectors, distributed over an area of 3000~km$^{2}$, and operating 24 wide-angle Schmidt telescopes, positioned at four sites at the border of the ground array.  The Observatory will reach its full size only in 2007 but data are routinely recorded already and have started to provide relevant science results.  This talk will focus on the detector characterizations and presents first results on the arrival direction of extremely-high energy cosmic rays, their energy spectrum, and on the upper limit of the photon fraction. \\vspace{1pc} ",
        "introduction": "Over the past decade, interest in the nature and origin of extremely high energy cosmic rays (EHECR) has grown enormously. Of particular interest are cosmic rays (CR) with energies $\\ga 10^{20}$~eV. There is a twofold motivation for studying this energy regime, one coming from particle physics because CRs give access to elementary interactions at energies much higher than man-made accelerators can reach, and another coming from astrophysics, because we do not know what kind of particles they are and where and how they acquire such enormous energies.  An excellent review, published by Michael Hillas 20 years ago, presented the basic requirements for particle acceleration to energies $\\ge 10^{19}$~eV by astrophysical objects.\\cite{Hillas84} The requirements are not easily met, which has stimulated the production of a large number of creative papers.  The problem is aggravated even more by the fact that at these energies protons and nuclei should interact with the Cosmic Microwave Background (CMB).  Above a threshold energy of $E_{\\rm GZK}\\simeq 5 \\times 10^{19}$~eV protons lose their energy over relatively short cosmological distances via photo-pion production $p+\\gamma_{\\rm CMB} \\to \\pi^{0}+p$ or $\\pi^{+}+n$. Accidentally, nuclei (He, \\ldots Fe) lose their energy at similar threshold energies and on even shorter length scales.  This is because of photodissociation (e.g. ${\\rm Fe} + \\gamma_{\\rm CMB} \\to X + n$) taking place mostly via giant nuclear resonances. Finally, photons interact even more rapidly in the CMB by producing $e^{+}e^{-}$-pairs.  Thus, particles that have traveled over distances of 50 or 100 Mpc are unlikely to retain an energy of $\\sim 10^{20}$~eV or more when they reach us.  This was already recognized in the 1960's shortly after the discovery of the CMB and is called the Greisen-Zatsepin-Kuzmin (GZK) cutoff.\\cite{GZK} Thus, not only do we not know how particles could obtain such extreme energies even in the most powerful astrophysical accelerators, these accelerators have to be located nearby on cosmological scales!  To solve this most pressing puzzle of high energy astroparticle physics, one either needs to invent nearby exotic EHECR sources or find ways of evading the GZK effect.  Top-Down models with decaying topological defects or decaying superheavy relic particles are typical representatives of the former group, as EHECRs would be produced nearby.  Typical representatives of the latter kind are violation of the Lorentz invariance, propagation of heavy supersymmetric particles, or the $Z$-burst model.  A comprehensive review, with emphasis placed on top-down models, is given by Ref.~\\cite{Bhattacharjee00}.  Generally, the top-down models predict a dominance of photons and neutrinos over protons or nuclei, so that measurements of the chemical composition become important also at the highest energies.  Furthermore, the $Z$-burst model cannot avoid producing a strong background of GeV energy photons leading to severe constrains due to the measured EGRET fluxes.\\cite{Sigl04} Such complications have recently given more emphasis again to astrophysical sources.  While the large magnetic rigidity of $\\sim 10^{20}$ eV protons gives rise to the problems of particle acceleration in astrophysical sources, it opens at the same time a new window for astronomy with CRs.  Since such particles cannot deviate much in the magnetic fields of the Galaxy and extragalactic space, they should point to their sources within a few degrees deviation only.  For example, using nominal guesses of 1~nG for the magnetic field strength of extragalactic space and 1 Mpc for the coherence length, deviations for protons on the order of $2.5^\\circ$ are expected after travelling 50 Mpc.\\cite{Cronin92}  Two types of experiments based on very different techniques have undoubtedly detected particles well exceeding the GZK cut-off.\\cite{NaganoWatson00,Takeda03,Abbasi04a} Unfortunately, despite 40 years of data taking the number of events is still small.  Also, the largest experiments so far disagree at an approx.\\ $2 \\sigma$ level on the flux and on arrival direction correlations.  The HiRes collaboration, employing the fluorescence technique, reported a suppression of the flux above the GZK-threshold, with no evidence for clustering in the arrival directions.\\cite{Abbasi04a,Abbasi04b} On the other hand, ground arrays have detected no GZK-cutoff.\\cite{NaganoWatson00,Takeda03} Furthermore, the the AGASA collaboration published results about seeing a clustering of the highest energy events \\cite{Takeda03} which, however, is not free of dispute.\\cite{Finley04} Clearly, the situation is very puzzling, and a larger sample of high quality data is needed for the field to advance. ",
        "conclusions": "The construction of the southern Pierre Auger Observatory is well underway.  About 1000 stations of the surface array and 18 telescopes of the fluorescence detector are in operation and taking data routinely.  Completion of the southern site is planned for 2007 and R\\&D work for the northern site to be located in south east Colorado has started.  Parallel to the completion of the observatory, first science results were already obtained on the energy spectrum, searches for localized anisotropies near the direction of the Galactic Center, and on setting upper limits on the photon fraction of the primary particles.  Considering the very limited statistics from 1.5 years of data taking during construction, being equivalent of only 3 months of a full array, it is not surprising that the emergence of a clear picture about the shape of the energy spectrum above the GZK threshold needs a little more time.  The 2.3 years of data used for the anisotropy searches in the Galactic Center region provides statistics much greater than those of previous experiments.  No evidence for a point-like source in the direction of Sagittarius\\,A$^{*}$ was found.  This excludes several scenarios of neutron sources in the GC suggested recently.  Our searches on larger angular windows in the neighborhood of the GC do not show abnormally over-dense regions. In particular, they do not support the large excesses reported in AGASA data (of 22\\% on $20^\\circ$ scales) and SUGAR data (of 85\\% on $5.5^\\circ$ scales).  The upper limit to the photon fraction above 10~EeV derived from a direct observation of the shower maximum confirms and reduces previous limits from ground arrays. Again, the current analysis is limited mainly by the small number of events.  \\begin{figure}[t] \\centerline{\\epsfxsize=8cm\\epsfbox{very-inclined.eps}} \\vspace*{-3mm}\\caption[xx]{Example of a near horizontal air shower as seen by the ground array. The shower has triggered 31 stations and extends over 30 km at ground. \\label{fig:inclined}} \\end{figure}  The number of hybrid events will considerably increase over the next years, allowing to set much stronger limits on the anisotropy and point source searches and on the photon limits. It will also reduce the uncertainties of the energy spectrum and will allow for further studies of EHECRs.  For example, the Pierre Auger Collaboration is developing the study of inclined events, and showers with zenith angles above $85^{\\circ}$ have been seen.  This was expected as they had been detected long ago with much smaller arrays, but the richness of the new data is impressive.  Figure \\ref{fig:inclined} shows an event at about $88^{\\circ}$ with 31 detectors, and even the present array is too small to contain it.  A preliminary estimate of its energy is around 30 EeV. An understanding of these events will lead to additional aperture for collection of the highest-energy particles and also give additional routes to understanding the mass composition.  Further, these events form the background against which a neutrino flux might be detectable.  There is an exciting future ahead.   \\vspace*{5mm} {\\small {\\bf Acknowledgement:} It is a pleasure to thank the organizers of the Symposium for their invitation to this Symposium.  The work of the group at University Wuppertal is supported in part by the German Ministry for Research and Education (Grant 05 CU5PX1/6).}"
    },
    "0608/astro-ph0608070_arXiv.txt": {
        "abstract": "The dense, clustered environment in which massive stars form can lead to interactions with neighboring stars.  It has been hypothesized that collisions and mergers may contribute to the growth of the most massive stars.  In this paper we extend the study of star-disk interactions to explore encounters between a massive protostar and a less massive cluster sibling using the publicly available SPH code GADGET-2. Collisions do not occur in the parameter space studied, but the end state of many encounters is an eccentric binary with a semi-major axis $\\sim 100$ AU.  Disk material is sometimes captured by the impactor.  Most encounters result in disruption and destruction of the initial disk, and periodic torquing of the remnant disk.  We consider the effect of the changing orientation of the disk on an accretion driven jet, and the evolution of the systems in the presence of on-going accretion from the parent core. ",
        "introduction": "In the standard picture of star formation, stars form from the collapse and fragmentation of turbulent molecular cloud cores on a timescale of $\\sim 10^{5}$ yr for individual stars to $\\sim 10^{7}$ yr for large OB associations.  Accretion onto the prestellar core from the surrounding envelope can occur spherically, or if the material has some angular momentum, from the circumstellar disk.  While the details of these processes are becoming understood for isolated low mass stars ($M_{\\star}\\lesssim10$ \\msun), the formation of clusters, multiple systems, and massive stars are less clear.  Simply extending the process for low mass stars to massive star formation presents some problems.  The Jeans mass of cloud cores containing massive stars is $\\sim0.3$ \\msun, while observed stellar masses are $\\sim50$ \\msuns \\citep{bon98a}.  \\citet{wol87} argue that radiation pressure from the central star may halt or reverse spherical accretion for masses greater than about 20 \\msun.  However, \\citet{mck03} suggest that high external pressure can overcome this obstacle.  Calculations of accretion from the envelope onto a disk, and from the disk onto the star \\citep{yor99,yor02} have resulted in protostars of $\\sim30$\\msun.  \\citet{kru05d,kru05b} argue that anisotropies in the stellar radiation pressure and instabilities in the accretion flow can circumvent the radiation pressure problem.  However, these scenarios deal solely with the formation of a star in isolation.  Young massive stars are usually found in dense clusters \\citep{cla00}, often near the central region \\citep{hil98}, and have a high multiplicity fraction \\citep{sta00, gar01}.  These factors suggest that interactions between neighboring stars are likely during cluster formation.  \\citet{bon02} performed numerical simulations of accretion and merging in a cluster of 1000 identical mass stars.  They found that 19 mergers occurred, driven by the dynamics of small subgroups within the cluster.  Though the stellar densities they required ($\\sim 10^8$ star pc$^{-3}$) are far in excess of those observed, it can be argued that such high densities exist briefly during the initial collapse of a high density molecular cloud core.  \\citet{bon98a} and \\citet{bon05} suggest a scenario wherein accretion onto a wide, low mass binary leads to a tight binary with a system mass of 30-50 \\msun, which, with its larger cross section, could be driven to collision by an encounter with a third body.  If one considers a cluster with number density $10^6$ pc$^{-3}$, a velocity dispersion of 3 \\kms, and populated by typical pre-main-sequence protostars with masses 0.3 \\msuns and radii 0.05 AU, one can estimate the mean time between collisions \\citep{bin87} as $t_{col} \\sim4 \\times 10^8$ yr.  This corresponds to a virialized cluster core with, e.g., 100 \\msuns ($\\sim$ 330 protostars) and radius $\\sim$ 0.04 pc.  Even at such an extreme number density, this rate is not high enough for collisions to be a likely star formation mechanism.  \\citet{bal05} suggest that circumstellar disks and envelopes could play a role beyond accretion by increasing the interaction radius, and therefore assisting stellar capture.  The cross section of a star-disk system is much larger than the cross section of a single star, with disks commonly a few$\\times10^2$ AU in radius. Material ejected during an encounter carries away kinetic energy from the impactor, which expands the range of initial conditions that could lead to a merger.  Additionally, by sweeping through the disk the impactor is slowed both by accreting material, and by the gravitational effects of wakes and spiral arms induced in the disk. Can these effects increase the merger radius of the protostars to $\\sim 100$ AU?  The increased gravitational focusing between a 20 \\msuns star in a mass-segregated cluster with density $n_{\\star} = 10^{5}$ pc$^{-3}$, similar to the density found in \\citet{bon03}, can drive the timescale for encounters within 200 AU to less than $10^{5}$ yr; encounters with stars more massive than 1 \\msuns take place on a timescale $\\sim 10^{5}$ yr (Moeckel \\& Bally 2006, in preparation).  Can the ensuing interactions lead to mergers, or to binaries which could evolve further through third body interactions or accretion onto the system?  Several authors have considered encounters between star-disk systems.  \\citet{hel95}, \\citet{lar97}, \\citet{bof98}, and \\citet{pfa05} used n-body and SPH codes, while \\citet{hal96} and \\citet{hal97} used a reduced-three body approach to study encounters ranging from penetrating to distant.  \\citet{cla91} approached the question of binary formation due to star-disk encounters analytically, while \\citet{ost94} analytically treated the case of more distant encounters.  \\citet{wat98,wat98a} performed SPH simulations of encounters wherein both stars have circumstellar disks.  Of the numerical studies dealing with close encounters, none consider stellar mass ratios $m_{2}/m_{1} < 0.5$, and \\citet{hal97} and \\citet{pfa04} are the only studies that consider repeated encounters. In this paper we present the results of numerical simulations concerning repeated, penetrating encounters between a massive protostar with a disk, and a less massive stellar impactor, to extend the parameter space considered in previous studies toward a regime suitable for the early life of a high mass star in a star forming cloud core.  The primary system is simulated in a late stage of formation, when the ratio of disk to stellar mass $M_d/M_{\\star} = 0.1$, though we consider one case with an accreting envelope surrounding the system. ",
        "conclusions": "\\begin{figure*} \\centering \\plottwo{f10a.eps}{f10b.eps} \\caption{\\textit{Left}: Semi-major axis evolution with time for i=180, with (\\textit{solid}) and without (\\textit{dot}) an accreting envelope. Apastra are marked with symbols.  \\textit{Right}: Impactor and primary masses, and mass of gas associated with the primary for the two cases.} \\label{envelope_results} \\end{figure*} We have presented a suite of simulations which consider the dissipative effects of a circumstellar disk on a capture formed binary, in the context of a massive protostar encountering a less massive cluster member.  The systems were followed for 10 or more orbits, in many cases leading to almost total destruction of the original disk.  We find that for the parameters studied, the interactions result in a soft, eccentric binary system.  The energy change associated with different inclination angles is in agreement with previous simulation results.  The periodic torquing of the disk by the passage of the impactor may be observable in the orientation of a jet or outflow.  These simulations apply to a massive protostar that has accreted most of its final mass.  The disk and envelope mass is $\\sim$ 10 \\% of the stellar mass, analogous to the class I phase of a low mass protostar's evolution.  Previous studies of stellar encounters with star-disk systems \\citep{hel95, bof98} concluded that disk capture rates were too low to contribute significantly to the observed multiplicity of low mass stars.  All but the least dissipative encounters simulated here result in a binary with a semi-major axis $\\lesssim$ 500 AU within $10^{4}$ yr.  Extending the analysis of binary capture rates to the regime where the disk mass and impactor masses are comparable, and much less than the primary, may show that the high multiplicity of massive stars such as the Trapezium stars in Orion could result from disk encounters in the late stages of the cluster's formation.  In order to follow the encounters through several orbits, the parameter space covered here is small; only a single impactor mass and initial periastron were considered.  The mass of the impactor was chosen as the minimum such that a merger would lead to significant growth.  As the impactor's mass approaches the primary's, disk destruction occurs over fewer orbits and less time is available to harden the binary. The apastra of the resultant binaries are on the order of 100s of AU, and the periastra of the most dissipative cases remain above $\\sim 25$ AU.  For the parameters studied, it seems that in the absence of further effects, disk dissipation alone cannot lead directly to mergers.  Accretion onto the system from the parent core is unlikely to lead to significant hardening of the binary over any short timescale, as the stellar masses considered are too high for reasonable mass accretion rates to have a large effect.  Thus if disk mediated collisions are to contribute to the growth of a massive star, they must occur earlier in the cluster's evolution, when the disk and envelope are more massive and perhaps more dissipative.  Even if this is the case, there must be an efficient disk replenishment mechanism in order for more than a single merger to occur.  However, at least two further mechanisms seem likely to come into play.  The disks considered in this work are Toomre stable; if a colder disk is considered, gravitational fragmentation within the disk will change the nature of the interactions, potentially leading to a more rapid decrease in the binary separation as smaller bodies are ejected from the system.  In a similar dynamical effect, a binary with a semi-major axis $\\sim 100$ AU presents a large target for interactions with other encounter partners in the dense cluster core.  In this scenario the capture and gradual tightening of the initial binary, modeled here, provides the initial conditions for encounters with other cluster stars.  The dynamics of $n>2$ bodies plus the remnants of a disk are beyond the scope of this work, but are an enticing direction for further study of systems such as these.  In summary, it seems that disk interactions in the late phase of protostellar evolution considered here do not provide a direct path to merger.  Further accretion from the natal cloud will not be enough to induce collisions from the resultant binaries.  In order for disks such as those modeled here to play a role in any collisional growth of massive stars, the interactions must involve third bodies acting on the capture formed binary, through either induced fragmentation of the disk or from other cluster members.  Another, yet unmodeled possibility is that the encounters leading to a merger take place during an evolutionary phase more like a class 0 protostar, with the mass of the primary comparable to that of the disk surrounding it.  Encounters with such systems have only been considered with impactors of similar mass to the primary, and not in the limit where the disk mass is greater than the impactor.  We thank the referee, Ian Bonnell, whose comments helped improved the manuscript.  This work was supported by the NASA Astrobiology Institute through a Cooperative Agreement with the University of Colorado, through grant NNA04CC11A."
    },
    "0608/astro-ph0608593_arXiv.txt": {
        "abstract": "Mid-infrared observations of the Andromeda galaxy, M31, obtained with the Infrared Array Camera on board the Spitzer Space Telescope are presented. The image mosaics cover areas of approximately $3\\fdg7 \\times 1\\fdg6$ and include the satellite galaxies M32 and NGC~205. The appearance of M31 varies dramatically in the different mid-infrared bands, from the smooth bulge and disk of the old stellar population seen at 3.6\\mic to the well-known `10 kpc ring' dominating the 8\\mic image. The similarity of the 3.6\\mic and optical isophotes and nearly constant optical-mid-infrared color over the inner 400\\arcsec\\ confirms that there is no significant extinction at optical wavelengths in M31's bulge. The nuclear colors indicate the presence of dust but not an infrared-bright active galactic nucleus. The integrated 8\\mic non-stellar luminosity implies a star formation rate of $0.4 M_{\\sun}$~yr$^{-1}$, consistent with other indicators that show M31 to be a quiescent galaxy. ",
        "introduction": "The proximity of the Andromeda galaxy, M31, has long made it a prime target for studies of galaxy structure and stellar populations. One technique has been to isolate particular populations: for example, studies of luminous stars \\citep{hmf90} showed that M31 has a lower massive star formation rate than M33 and the LMC. Global studies of particular components have also been valuable: H~I maps provide estimates of the mass distribution \\citep{bt04} and have recently revealed a population of high-velocity clouds \\citep{thil04}. {\\it ISOPHOT} \\citep{haas98}, {\\it MSX} \\citep{kraemer02}, and {\\it IRAS} \\citep{habing84} observations showed a bright ring of infrared emission with radius 10~kpc coincident with many of the H~II regions.  Mapping the entire disk of M31 at mid-infrared wavelengths allows {\\em both} local and global studies of the galaxy.  Observations with the Infrared Array Camera \\citep[IRAC;][]{irac} on the {\\it Spitzer Space Telescope} simultaneously trace the dust in the spiral arms at 8\\mic and the oldest stars in the disk and bulge at 3.6 and 4.5\\mic without the complicating extinction or distance effects that make such studies in the Milky Way difficult. IRAC observations of M31 are complemented by deep data now available at many other wavelengths. They also complement {\\it Spitzer} studies of other nearby galaxies.  This paper presents an initial look at the IRAC observations of M31, focusing on the surface brightness profiles and extended emission. Companion papers discuss longer-wavelength MIPS observations of M31 \\citep{gordon06}, IRAC and MIPS observations of the M31 satellite galaxy NGC~205 \\citep{marleau06}, and the implications of the galaxy's morphology as seen in 8\\mic non-stellar emission \\citep{block06}. A distance to M31 of 783~kpc \\citep{sg98} is assumed throughout. All magnitudes are on the Vega system, using the calibration given by \\citet{reach05}. ",
        "conclusions": ""
    },
    "0608/astro-ph0608520_arXiv.txt": {
        "abstract": "{Realistic atmospheric models that link the properties and the physical conditions of supernova ejecta to observable spectra are required for the quantitative interpretation of observational data of type~Ia supernovae (\\snia) and the  assessment of the physical merits of theoretical supernova explosion models. The numerical treatment of the radiation transport -- yielding the synthetic spectra -- in models of {\\snia} ejecta in early phases is usually carried out in analogy to atmospheric models of `normal' hot stars. Applying this analogy indiscriminately leads to inconsistencies in {\\snia} models because a diffusive lower boundary, while justified for hot stars, is invalid for hydrogen and helium-deficient supernova ejecta. In type~Ia supernovae the radiation field does not thermalize even at large depths, and large optical depths are not reached at all wavelengths.} {We aim to derive an improved description of the lower boundary that allows a more consistent solution of the radiation transfer in {\\snia} and therefore yields more realistic synthetic spectra.} {We analyze the conditions that lead to a breakdown of the conventional diffusion approximation as the lower boundary in {\\snia}.  For the radiative transfer, we use a full non-LTE code originally developed for radiatively driven winds of hot stars, with adaptations for the physical conditions in SN~Ia.  In addition to a well-tested treatment of the underlying microphysical processes, this code allows a direct comparison of the results for {\\snia} and hot stars.} {We develop a semi-analytical description that allows us to overcome some of the limiting assumptions in the conventional treatment of the lower boundary in {\\snia} radiative transfer models. We achieve good agreement in a comparison between the synthetic spectrum of our test model and an observed spectrum. } {} ",
        "introduction": "Type~Ia supernovae (\\snia) have become an invaluable tool for the determination of the cosmological parameters \\citep{riess98, perlmutter99, riess01, tonry03} as their exceptional brightness makes them observable even at large cosmological distances.  Using {\\snia} for distance determination requires knowledge of their absolute luminosities; however, {\\snia} are not perfect ``standard candles'' in this respect because they show an intrinsic scatter in their properties, in particular in the peak brightness. The application of {\\snia} for cosmology therefore relies on empirical relationships between the peak brightness and other observed characteristics \\citep[e.g.,][]{phillips93, hamuy96b, riess96, perlmutter97}.   To first order the shape of the light curves and certain spectral properties is determined by the mass of synthesized {\\nifs} and its distribution within the ejecta \\citep{nugent95b,hoeflich95,pinto00a}. However, the details of the physical processes that cause the observed variation are still unclear (see \\citealt{hillebrandt00} for a review).  This uncertainty is an essential problem because the different calibration methods used to derive the peak brightness yield partly different results, implying an unaccounted source of systematic error \\citep{leibundgut04}.  In addition, the application of {\\snia} for cosmological distance measurement relies on the crucial assumption that objects observed at high redshifts have the same properties as objects in the nearby universe, so that the same calibration method for the luminosity differences can be applied to all objects.  It is, therefore, of fundamental interest to derive a physical model that can explain the explosion mechanism of {\\snia} in detail, including the observed intrinsic variability. This will allow a more reliable estimate of the systematic errors of the distance measurement.  The currently favored models for the explosion mechanism of {\\snia} involve the thermonuclear explosion of a carbon-oxygen white dwarf (WD) in a binary system. Two progenitor scenarios are generally considered: the ``single degenerate'' scenario and the ``double degenerate'' scenario. In the single degenerate scenario the WD accretes  mass from a red-giant companion star. When the WD reaches a mass close to the Chandrasekhar mass ($M_{\\rm Ch}\\approx 1.4\\msun$) the compressional heating at the center of the star triggers carbon burning. After a period of a few thousand years of quiet burning a thermonuclear runaway disrupts the star \\citep{iben84, webbink84, woosley04, han04}. In the ``double degenerate'' scenario the companion star is also a WD with a mass such that the total mass of the system exceeds $M_{\\rm Ch}$.  Due to energy loss by gravitational waves the orbital separation of the binary system gradually decreases, leading to a merger, which triggers the thermonuclear runaway that explodes the star \\citep{whelan73, nomoto82}. Potential progenitor systems have been found in recent years; their numbers, however, are too low to explain the observed rates of SN Ia \\citep{cappellaro99}. \\citet{pauldrach05}  suggests on basis of the results of \\citet{pauldrach04} a potential connection of {\\snia} progenitors to a subgroup of central stars of planetary nebulae (CSPN).  The details of the explosion process itself are also still subject to a lively debate.  The general picture is that a subsonic deflagration wave (``flame'') is ignited near the center of the star. The flame travels outward, burning part of the star to nuclear statistical equilibrium (NSE). Because the flame propagates subsonically, the star can expand while undergoing burning. This allows partial burning of C and O to intermediate mass elements (Si, S, Mg, Ca) which dominate the composition. In contrast to the deflagration, a prompt, supersonic detonation of the star would generate predominantly iron group elements. This outcome contradicts the observed composition.  No agreement has been reached in the debate whether the explosion continues subsonically until the end of the explosion \\citep{nomoto84,woosley84,niemeyer95,reinecke02,roepke05,roepke05c} or if a (yet unknown) mechanism triggers the deflagration to turn into a supersonic detonation toward the end of the explosion (delayed detonation transition, DDT) \\citep{hoeflich96,iwamoto99}. At present, the resulting composition of DDT models seem to favor the latter scenario; however, those models are not self-consistent and depend on ad hoc assumptions of the occurrence of the DDT.  Judging the validity of numerical explosion models requires a comparison of the observational consequences predicted by the explosion models to the {\\em de facto} observations of {\\snia}. Realistic radiative transfer models provide the crucial link between explosion models and observations. It is possible to predict the observational implications of the hydrodynamical models only if the radiative transfer models are sufficiently realistic and detailed.  Conversely, such radiative transfer models make it possible to establish constraints on the composition and structure of {\\snia} from the spectroscopic interpretation of observed spectra.  {\\snia} in phases before and shortly after maximum exhibit a spectrum dominated by a few very broad absorption features in a non-thermal continuum.  These absorption features are mostly due to blends of several lines, while the ``continuum'' itself is formed by the overlap of a large number of Doppler-broadened metal lines. The true continuum opacities and emissivities that determine the overall shape of the spectrum in stars are of minor significance in supernovae. At later epochs, the ``pseudo-continuum'' photosphere recedes deeper and deeper into the ejecta and eventually disappears as the ejecta become transparent. Unlike stars, the radiation in {\\snia} is generated within the expanding medium itself by the deposition of the energy of $\\gamma$-photons that result from the decay chain {\\nifs}\\ra{\\cofs}\\ra{\\fefs} \\citep{colgate69}.  The early epochs where a photosphere is still present are referred to as the photospheric phase. In the photospheric phase the ejecta can be treated in analogy to hot stars as expanding, extended atmospheres. Only the radiative transfer in the photosphere and in the outer envelope above the photosphere needs to be considered. In this setup a steady state is assumed because the photon escape timescale in the thin medium generally will be much shorter than the expansion timescale.  In principle, a complete radiative transfer model for {\\snia} would require consistent, time-dependent solutions of the populations of all atomic levels and the continuum and line transfer, including the treatment of energy deposition by the decay products of {\\nifs} and {\\cofs}.  Current models implement various simplifications according to the specific model's purpose.  These simplifications are necessary because the solution of time-dependent radiative transfer in three dimensions, including the full coupling of radiation and matter, is not yet feasible and some of the terms involved in such a consistent solution are shown to be or regarded to be of second order.  In recent decades synthetic spectra of {\\snia} have been modeled by several groups with a variety of approaches involving different levels of complexity depending on the application \\citep{branch85, mazzali93a, mazzali93b, eastman93, hoeflich95, nugent95a, pauldrach96, nugent97, lentz01, hoeflich05, stehle05, baron06,kasen06}.  Highly parametrized models, which implement a simplified treatment of physical processes to achieve short run-times seem to be suitable for the comparative analysis of a large number of observed spectra, while more realistic models are required for a deeper understanding of the physical effects leading to specific observed properties. In particular, judging the validity of hydrodynamic explosion models, as mentioned above, can only be performed using radiative transfer models that include a very detailed treatment of relevant physical processes.  Such detailed models may also be used to validate or invalidate specific simplifying assumptions used in less elaborate models.  In spite of the analogy mentioned before, {\\snia} ejecta differ from ordinary stellar atmospheres in several important aspects. Techniques and methods that are adopted for stellar atmosphere modeling must be carefully checked to verify whether the applied approximations are still justified for {\\snia}. In this work we present an improved description for the lower boundary of the radiative transfer calculation.  This is required because the assumption of thermalization at large depths, justified for stellar atmospheres, breaks down for {\\snia}.  In the work we present here, we use the computer program {\\em WMbasic} \\citep{pauldrach94, pauldrach01} to obtain a consistent solution of the full non-LTE rate equations and a detailed observer's frame solution of the radiative transfer.  This code was originally designed for the analysis of the spectra of hot stars with radiatively driven winds, but an earlier version has already been used by \\citet{pauldrach96} (Paper~I) to quantitatively investigate the effects of line blocking in {\\snia}. While they used a consistent treatment of line blocking, the back-reaction of the line opacities on the temperature structure (line blanketing) was not taken into account.  In our present work the current version of the code has been further adapted to treat the radiative transfer in supernovae in a more sophisticated way.  Section~\\ref{sec:model} describes the setup of the non-LTE model and introduces the numerical scheme used to solve the radiative transfer. In Section~\\ref{sec:photosphere} the physical conditions in the ``pseudo-photosphere'' of {\\snia} are discussed with respect to the solution of the radiative transfer and are compared to the situation in normal stars (i.e., stars that have a well defined photosphere).  Section~\\ref{sec:bc} describes the derivation of an improved treatment of the inner boundary for the numerical solution of the radiative transfer. The results are discussed in Section~\\ref{sec:discussion}, and a comparison of a model spectrum with an observed {\\snia} spectrum is shown.  The conclusions are provided in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have shown that the physical conditions in the expanding atmospheres of {\\snia} are such that even at early times a complete thermalization of photons cannot be assumed.  Therefore the common approach using an LTE diffusive boundary condition at the inner point of the computational grid does not provide a consistent description of the radiation field and may lead to non-realistic synthetic spectra.  Observationally this can be seen in the red and infrared wavelength bands where the spectral slope of typical {\\snia} spectra deviates significantly from a thermal continuum. The assumption of such a thermal continuum in radiative transfer models generally results in an overestimate of the radiation flux in those wavelengths. We have developed a theoretical framework to eliminate some of the restrictions that are imposed by the assumption of LTE conditions.  With the formalism discussed in this work we are able to self-consistently derive the isotropic term of the boundary condition. Only the flux term has to be specified analytically by taking  the physical conditions in the photospheric region into account.  Of course, removing constraints in the boundary conditions may cause the system to become numerically less stable. For this reason the consistency of the solution is much more important than in the classical case where the explicit specification of the Planck-function forces the system into LTE-conditions naturally.  Our modifications to the LTE-diffusion approximation have been derived in a very general way. Therefore, this formalism can also be applied to other objects (such as Wolf-Rayet stars with very extended atmospheres) where the physical conditions are such that the assumption of LTE at the photosphere is not justified.  The comparison of the synthetic spectrum from a {\\snia} test model to an observed {\\snia} spectrum shows that the overall shape and prominent features of observed {\\snia} are well reproduced by the model. A more detailed analysis of observed spectra will be the subject of forthcoming publications.  Late epochs of {\\snia} have not been considered here in more detail because the energy deposition by $\\gamma$-photons above the photosphere has not yet been fully implemented in our code.  The description developed here nevertheless provides a basis for a reliable implementation of this energy deposition, since even if a fraction of the radiative energy is created above the photosphere the remaining radiation has to originate from below the computational grid. However, given that the photospheric conditions are such that thermalization occurs only partially, it is impossible that the energy deposited in the outer ejecta will be completely thermalized. Instead, excitation and ionization by fast electrons and $\\gamma$-photons above the photosphere may provide a non-thermal contribution to the spectrum.  This might already be an issue even for epochs around maximum, as recent three-dimensional explosion models indicate extensive mixing of {\\nifs} into the outer layers of the ejecta.  Furthermore, to derive a model with a luminosity that is consistently determined from the {\\nifs} distribution of an explosion model, the time-dependent effects of photon trapping in earlier epochs have to be incorporated into the boundary luminosity \\citep{arnett82,hoeflich96,nugent97}."
    },
    "0608/gr-qc0608084_arXiv.txt": {
        "abstract": "We analyze two kinds of matched filters for data output of a spherical resonant GW detector. In order to filter the data of a real sphere, a strategy is proposed, firstly using an omnidirectional in-line filter, which is supposed to select periodograms with excitations, secondly by performing a directional filter on such selected periodograms, finding the wave arrival time, direction and polarization. We point out that, as the analytical simplifications occurring in the ideal 6 transducers TIGA sphere do not hold for a real sphere, using a 5 transducers configuration could be a more convenient choice. ",
        "introduction": "Spherical resonant gravitational wave (GW) detectors are the evolution of bar resonators. For an equivalent size they have a higher-cross section than other geometries and are the only kind of GW detectors able to to determine the signal arrival direction and polarization, even operated at a single location on Earth. Nowadays two spheres are under development in the Netherlands (MiniGRAIL)\\cite{MG} and in Brasil (MARIO SCHENBERG)\\cite{MS}. Their experimental situation at present requires an efficient and concrete data-analysis strategy, \\flo{which supplement the available literature (see for exemple \\cite{GT},\\cite{Merkowitz:1997qs})}. Works in this direction have already been done, see for example \\cite{Caesar}.   In this paper we will focus on the case of matched filtering for such a detector, starting from some basic notions of \\cite{Stevenson}. Matched filtering has been shown to be a key point in the burst oriented data analysis for resonant detectors \\cite{[PIA]} and for interferometers~\\cite{Schutz}. We will discuss the question from a theoretical point of view, keeping in mind the technical implementation of our propositions. We propose a strategy in order to filter the sphere data, firstly by an omnidirectional in-line filter, which we call ``refined energy matched-filter\", in order to select signal excitations, then by performing a ``directional matched-filter\" on events selected by the first filter, extracting the wave arrival time, direction and polarization.  In actual experiments, the preferred configuration for the transducer placement is the truncated ichosaedral (TI) proposed by Merkowitz and Johnson \\cite{Merkowity95}. We find that when data analysis is performed for a realistic sphere with resonators and electronic readout, the analytical simplifications of an ideal sphere with transducers placed in the TI configuration do not hold anymore. Therefore using a 5 transducer configuration as the pentagonal hexacontahedron (PHC) proposed by Lobo and Serrano \\cite{Lobo} can be a more convenient possibility, because simpler to build and manage.  This paper is structured as follows: in the next section we introduce notations and a schematic model of spherical detector. In section \\ref{SF1} and \\ref{SF2} we present two different ways of implementing a matched filter with a description of advantages and limitations for each one. In the last section we draw a strategy in order to perform data analysis with these two filters and end with our main conclusions. ",
        "conclusions": "In this paper we have analyzed two ways for matched-filtering data output of a spherical GW detector. In section \\ref{SF2} the ``refined energy matched filter\" has been presented, whose efficiency does not depend on the GW direction.  For the specific setup of each existing GW spherical detector, it would be worth to investigate in more details the performance of the pipeline data analysis proposed in this work: firstly using the REMF as an in-line filter in order to select periodograms with excitations. REMF requires some computer memory, since the set of matrices $\\hat\\mtx{B}(\\omega)$ has to be stored, but only few more stages of computation than the matched filter used for the cylindrical detectors.  Then, on periodograms selected by REMF, the directional matched filter presented in section \\ref{SF1} of this paper can be used in order to extract the arrival time, direction and polarization of the GW signal. Furthermore, once we have obtained (for a chosen $\\psi$) $\\tilde{h}_+$ and  $\\tilde{h}_\\times$ from the DMF procedure, we have a check of the coherence of the whole procedure by making the comparison with the value of $\\vert\\tilde{h}_+\\vert^2+\\vert\\tilde{h}_\\times\\vert^2 $ obtained from the refined energy matched filter.  Finally we stress that the mode-channel simplifications occurring when the 6-transducers TI configuration is considered do not hold in a realistic model of a spherical antenna. Therefore, for the data analysis strategy proposed in this work, it may be convenient to choose working with only five transducers.  In such a case the $\\hat\\mtx{G}$ is a $5\\times5$ matrix and, finding a configuration of the transducer such that it is invertible for all values of $\\omega$, we can perform a REMF with $\\hat{\\mtx{B}}=\\hat{\\mtx{G}}^{-1}$. \\\\  We would like to thank C. A. Costa and A. Malaspinas. This work is partially supported by the Swiss National Science Foundation (FNS)."
    },
    "0608/astro-ph0608650_arXiv.txt": {
        "abstract": "The X-ray spectra of the magnetar candidates are customarily fitted with an empirical, two component model: an absorbed blackbody and a power-law. However, the physical interpretation of these two spectral components is rarely discussed. It has been recently proposed that the presence of a hot plasma in the magnetosphere of highly magnetized neutron stars might distort, through efficient resonant cyclotron scattering, the thermal emission from the neutron star surface, resulting in production of non-thermal spectra. Here we discuss the Resonant Cyclotron Scattering (RCS) model, and present its XSPEC implementation, as well as preliminary results of its application to Anomalous X-ray Pulsars and Soft Gamma-ray Repeaters. ",
        "introduction": "\\label{intro} Anomalous X-ray Pulsars (AXPs) and Soft Gamma-ray Repeaters (SGRs) are a small class of slowly rotating (5-12 s) neutron stars with emission properties much at variance with those of ordinary X-ray pulsars, both the young radio pulsars and the X-ray binary pulsars. They are called ``anomalous'' because their high X-ray luminosity ($10^{34} - 10^{36}$ erg/s) cannot be easily explained in terms of the conventional processes which apply to other classes of pulsars, i.e. accretion from a binary companion or injection of rotational energy in the pulsar wind/magnetosphere. On the other hand, measurements of spin periods and period derivatives, when the latter are interpreted as due to electromagnetic dipolar losses, suggest that these objects may host ``magnetars'', i.e. neutron stars endowed with an ultra-strong magnetic field ($10^{14} - 10^{15}$\\,G, see Duncan \\& Thompson 1992; Thompson \\& Duncan 1993, 1996). The magnetar scenario appears so far very promising in explaining both the main energy source of these objects (the decay of the super-strong field) and the emission of the short, energetic bursts. Moreover, the magnetar scenario can account for the properties of Giant Flares, extremely energetic transient events ($L\\approx 10^{44}-10^{47}$ erg/s) detected from SGRs. However, alternative scenarios to explain the enigmatic properties of these sources have been invoked. Among these, models involving accretion from a fossil disk, left by the supernova event which gave birth to the neutron star, are still largely plausible (van Paradijs et al. 1995; Chatterjee, Hernquist \\& Narayan 2000). Very recently such a (maybe passive) disk has been indeed observed in the IR around AXP 4U 0142 + 61 (Wang, Chakrabarty \\& Kaplan 2006).  Magnetars candidates are strong X-ray sources, and their X-ray spectra in the 0.5-10~keV band are usually fit with a two component model consisting of (besides interstellar absorption) a thermal blackbody with a typical temperature $kT\\sim$0.3--0.4 keV and a power-law with a relatively steep photon index, $\\Gamma\\sim$2--4 (see Kaspi 2006 in this volume, and Woods \\& Thompson 2006 for recent reviews). In some cases, SGR spectra have been fit with a single power-law component, but recent longer observations showed that, also for these sources, a blackbody component is often required (Mereghetti et al. 2005).  Despite the fact that the blackbody plus power-law spectral model has been largely applied to magnetar spectra for many years, a reliable physical interpretation of these two components is still missing. Quite recently, an attempt to interpret the observed spectrum and long term variations of 1E 1048.1--5937 in terms of a (non resonant) Comptonization model has been presented by Tiengo et al. (2005).  Furthermore, the recent discovery of magnetar counterparts in the radio to the $\\gamma$-ray bands (Camilo et al. 2006; Hulleman et al. 2000; Kuiper et al. 2004) enforced the idea that their multiwavelength spectral energy distribution is by far more complicated than a simple blackbody plus power-law distribution (see also Den Hartog et al. contribution in this volume).      \\begin{figure} \\centering \\includegraphics[height=0.35\\textwidth,width=0.45\\textwidth]{ioutlog.ps} \\caption{Modification of an initially Plankian spectrum (dashed line) by multiple cyclotron scattering for different values of $\\tau_0$ and $\\beta_{th}$=0.3 (from Lyutikov \\& Gavriil 2006).} \\label{fig:models} \\end{figure}   \\begin{figure}[t] \\centering \\vbox{ \\includegraphics[height=0.3\\textwidth,angle=270,width=0.45\\textwidth]{1048_bbpl_eufs_c.ps} \\includegraphics[height=0.3\\textwidth,angle=270,width=0.45\\textwidth]{1048_rcs_eufs_c.ps}} \\caption{The XMM--Newton spectrum of 1E 1048--5937 fitted with an absorbed power-law plus a blackbody (top panel) and with the RCS model (second panel).} \\label{fig:1e1048}       % \\end{figure} ",
        "conclusions": "The preliminary results reported here show that the X-ray spectra of the AXP 1E\\,1048--56 can be successfully reproduced by the RCS model.  The quality of the fit is comparable to that obtained by using the canonical blackbody plus power-law model, with the advantage that the RCS model spectra provide a better physical understanding of the presence of the thermal/ non-thermal components in the observed spectrum. On the other hand, this is not totally unexpected  since this source has a non-thermal tail which is relatively soft with  respect to other AXPs, and a similar result has already been found by  Lyutikov \\& Gavriil, although their fitting technique was less accurate. For the other two sources considered here, 1RXS\\,J1708--4009 and SGR\\,1806--20, we find that the RCS model alone cannot account for the relatively hard non-thermal component observed below $\\sim$10~keV.   On the other hand, in both cases a composite model consisting of a RCS model plus a power-law can fit the X-ray spectra in the 1--100~keV band. This may  suggest that, in the 6--10\\,keV energy range, the X-ray  emission of these sources is contaminated by the hard non-thermal component which has beed recently observed with INTEGRAL at much higher energies (Kuiper et al. 2006). A similar hard tail has not been observed, at least so far, in the case of 1E\\,1048--56. A detailed paper with a systematic application of the RCS model to the whole class of magnetars and a study of their spectral and flux variability in this scenario is in preparation, accompanied by a  more complete discussion (Rea et al.~2006, in prep.)."
    },
    "0608/astro-ph0608699_arXiv.txt": {
        "abstract": "If cosmic ray protons interact with gas at roughly the mean density of the interstellar medium in starburst galaxies, then pion decay in starbursts is likely to contribute significantly to the diffuse extra-galactic background in both $\\gamma$-rays and high energy neutrinos.  We describe the assumptions that lead to this conclusion and clarify the difference between our estimates and those of Stecker (2006). Detection of a single starburst by {\\it GLAST} would confirm the significant contribution of starburst galaxies to the extra-galactic neutrino and $\\gamma$-ray backgrounds. ",
        "introduction": "\\label{section:introduction}  Inelastic proton-proton collisions between cosmic rays and ambient nuclei produce charged and neutral pions, which subsequently decay into electrons, positrons, neutrinos, and $\\gamma$-rays.  Neutral pion decay accounts for most of the $\\gamma$-ray emission from the Milky Way at GeV energies (e.g., Strong et al.~2004).  The contribution of pion decay to high energy emission from other galaxies --- and thus also to the extra-galactic neutrino and $\\gamma$-ray backgrounds --- remains uncertain, however, because of the current lack of strong observational constraints.  For example, aside from the LMC (Sreekumar et al.~1992), EGRET did not detect $\\gamma$-ray emission from star formation in any other galaxy.  In two recent papers (Loeb \\& Waxman 2006; Thompson, Quataert, \\& Waxman 2006; hereafter LW and TQW, respectively), we have estimated that starburst galaxies contribute significantly to the extra-galactic neutrino and $\\gamma$-ray backgrounds, with most of the contribution arising from starburst galaxies at $z \\sim 1$.  TQW also provide detailed predictions for the $\\gamma$-ray fluxes from nearby starbursts, several of which should be detectable with {\\it GLAST}.  Using what seem to be similar arguments, however, Stecker (2006) presents an upper limit on the starburst contribution to the neutrino and $\\gamma$-ray backgrounds that is a factor of $\\approx 5$ below the predictions of LW and TQW. Here we clarify the difference between these predictions (see \\S\\ref{section:other}).  The timescale for cosmic ray protons to lose energy via $p$-$p$ collisions with gas of number density $n$ cm$^{-3}$ is $\\tau_{pp}\\approx7\\times10^7/n$\\,yr.  The importance of pion losses depends on the ratio of $\\tau_{pp}$ to the cosmic-ray proton escape timescale, $\\tau_{\\rm esc}$.  The escape timescale for cosmic rays is quite uncertain, but in starbursts it is probably set by advection in a galactic superwind, rather than diffusion, as in the Galaxy.  If correct, the escape timescale from a starburst with a wind of velocity $V\\approx300$\\,km\\,s$^{-1}$ and a gas scale height of $h\\approx100$\\,pc is of order $\\tau_{\\rm esc}\\approx h/V\\approx3\\times10^5$\\,yr.  This implies that if cosmic rays interact with gas of mean density $n\\gtrsim$100\\,cm$^{-3}$, then $\\tau_{pp}\\lesssim\\tau_{\\rm esc}$ and we expect that essentially all cosmic ray protons are converted into charged and neutral pions and their secondaries before escaping. In this limit, the galaxy is said to be a ``cosmic ray proton calorimeter.''  If escape is instead due to diffusion rather than advection, then the density threshold required for a galaxy to be a proton calorimeter is likely to be significantly lower.  In the Milky Way, few GeV cosmic-ray protons only lose $\\approx 10\\%$ of their energy to pion production before escaping the galaxy.  For this reason, LW and TQW focused on the high energy emission from luminous starbursts.  These galaxies have much higher gas densities and likely dominate the contribution of star-forming galaxies to the total $\\gamma$-ray and neutrino backgrounds.  For reference, we note that the local starbursts NGC\\,253 and M\\,82 have $n\\gtrsim400$\\,cm$^{-3}$, and the nuclei of Arp\\,220  have gas densities of order $10^4$\\,cm$^{-3}$, well in excess of that required by the above arguments for  $\\tau_{pp}<\\tau_{\\rm esc}$.  Since the fraction of the proton energy supplied to secondary electrons, positrons, neutrinos and $\\gamma$-rays is determined by the microphysics of the Standard Model, the $\\gamma$-ray and neutrino fluxes from starbursts depend primarily on astrophysical properties of star forming galaxies via  \\begin{enumerate}  \\item{the fraction of a SN's canonical $10^{51}$ ergs of energy supplied to cosmic ray protons, $\\eta$,}  \\item{the spectral index $p$ of the injected cosmic ray proton distribution,}  \\item{the evolution of the comoving star formation rate density of the universe, $\\dot{\\rho}_\\star(z)$ (M$_\\odot$ yr$^{-1}$ Mpc$^{-3}$),}  \\item{and the fraction of all star formation that occurs in the proton calorimeter limit as a function of cosmic time, $f(z)$.}  \\end{enumerate}  For any $\\eta$ and $p$, the $\\gamma$-ray and neutrino emission from a given proton calorimeter can be readily calculated, while the cumulative extra-galactic backgrounds also depend on the star formation history of the universe via $\\dot{\\rho}_\\star(z)$ and $f(z)$.  As TQW and LW show, in the proton calorimeter limit the total $\\gamma$-ray and neutrino luminosities for an individual starburst can be related to $L_{\\rm TIR}$, the galaxy's total IR luminosity, by \\beq L_\\gamma \\approx(2/3) L_{\\rm nu}\\approx 1.5 \\times 10^{-4} \\eta_{0.05}L_{\\rm TIR}, \\label{gamma_neutrino_fir} \\eeq where $\\eta_{0.05}=\\eta/0.05$ (i.e., $5\\times10^{49}\\eta_{0.05}$\\,ergs per supernova is supplied to cosmic ray protons).  $\\eta$ is constrained by the requirement that the radio emission from secondary electrons/positrons in starbursts be consistent with the observed FIR-radio correlation (see, e.g., Yun et al.~2001).  For proton calorimeters, $\\eta \\gtrsim 0.05$ is only possible if ionization, bremsstrahlung, and/or inverse-Compton losses dominate synchrotron losses for secondary electrons and positrons in starbursts (see TQW; Thompson et al.~2006).  \\begin{figure} \\centerline{\\hbox{\\psfig{file=plot_sfr.ps,width=9cm}}} \\figcaption[x]{The comoving star formation rate density $\\dot{\\rho}_\\star(z)$ for a number of models and fits to existing data: RSF1 ({\\it dashed black}), RSF2 ({\\it solid black}), and RSF3 ({\\it dotted black}) are from Porciani \\& Madau (2001); the ``Fossil'' model ({\\it solid red line}) is a simple broken power-law fit to Nagamine et al.~(2006) (their Fig.~5); the ``GALEX'' model ({\\it solid blue line}) is the dust-corrected $\\dot{\\rho}_\\star(z)$, adjusted to the same assumed IMF as the other models presented here, from Figure 5 of Schiminovich et al.~(2005) ($\\dot{\\rho}_\\star(z)$ assumed constant for $z>3$); the ``H\\&S'' model ({\\it solid green line}) is an analytic function from Hernquist \\& Springel (2003).  The total TIR background for these models as computed using eq.~(\\ref{tir}) is $F^{\\rm TIR}=19.7$, 20.4, 23.0, 15.9, 19.5, and 17.3\\,nW\\,\\,m$^{-2}$\\,\\,sr$^{-1}$ for RSF1, RSF2, RSF3, Fossil, GALEX, and H\\&S models, respectively. These numbers for $F^{\\rm TIR}$ can be compared with observational constraints that indicate  $F^{\\rm TIR}\\sim24-27.5$\\,nW\\,\\,m$^{-2}$\\,\\,sr$^{-1}$ (Dole et al.~2006). Taking $f(z)=\\min(0.9z+0.1,1)$, $f_{\\rm cal}=0.79$, 0.81, 0.81 0.76, 0.75, and 0.73 for these same models. Using these values of $F^{\\rm TIR}$ and $f_{\\rm cal}$, the total $\\gamma$-ray and neutrino backgrounds can be calculated from eq.~(\\ref{gamma_neutrino_fir}) as in eq.~(\\ref{background_total}). Figure \\ref{fig:b} shows the cumulative fraction of $F^{\\rm TIR}$ as a function of redshift, $\\zeta(z)$, for each of these models (eq.~\\ref{zeta}).\\\\ \\label{fig:sfr}} \\end{figure} ",
        "conclusions": "\\label{section:conclusion}  The magnitude of the observed diffuse $\\gamma$-ray background is quite uncertain, primarily as a result of foreground subtraction. Whereas Sreekumar et al.~(1998) (compare with Strong et al.~2004) find a specific intensity of $1.4\\times10^{-6}$\\,GeV\\,\\,s$^{-1}$\\,\\,cm$^{-2}$\\,\\,sr$^{-1}$ at GeV energies, Keshet et al.~(2004) derive an {\\it upper limit} of $\\nu I_{\\nu_\\gamma}<5\\times10^{-7}$\\,GeV\\,\\,s$^{-1}$\\,\\,cm$^{-2}$\\,\\,sr$^{-1}$. Future observations by {\\it GLAST} should alleviate some of the uncertainty in the background determination.  A comparison of the upper limit by Keshet et al.~(2004) and equation (\\ref{background}) indicates that starburst galaxies may indeed contribute significantly to the diffuse $\\gamma$-ray background.   Even in stacking searches, {\\it EGRET} did not detect a single starburst galaxy (Cillis et al.~2005).  In TQW (their Figure 1) we show that the {\\it EGRET} non-detections are required by predictions that are correctly calibrated to the FIR or radio emission from galaxies.  Hence, the {\\it EGRET} non-detections do not provide a significant constraint on the contribution of starburst galaxies to the $\\gamma$-ray or neutrino backgrounds.  {\\it GLAST} will, however, be much more sensitive.  Using equation (\\ref{gamma_neutrino_fir}) and numerical calculations for different $p$, TQW show that {\\it GLAST} should be able to detect several nearby starbursts (NGC\\,253, M\\,82, IC\\,342; see TQW Table 1) if they are indeed proton calorimeters.  The dominant uncertainty in assessing the starburst contribution to the $\\gamma$-ray and neutrino backgrounds lies in whether a significant fraction of cosmic-ray protons in fact interact with gas of average density or whether the cosmic rays escape in a galactic wind without coupling to the high density ISM (the star formation history of the universe is comparatively well understood).  Detection of a single starburst galaxy by {\\it GLAST} would strongly support the estimates of LW and TQW and thus the importance of starburst galaxies for the extra-galactic neutrino and $\\gamma$-ray backgrounds.  Such a detection would also provide an important constraint on the physics of the ISM in starburst galaxies."
    },
    "0608/astro-ph0608466_arXiv.txt": {
        "abstract": "We revisit a method to obtain upper limits on the jet matter content combining synchrotron self-absorption constraints and the large scale bubble energy.  We use both X-ray observations, which give limits on the jet power from the energies and timescales of bubbles found in clusters of galaxies, and radio observations, which give limits on the magnetic field in the jets.  Combining the two imposes constraints on the particle number density, and hence the jet content.  Out of a sample of clusters which have clear radio bubbles, there are only two which have sufficient resolution in the radio images to give significant constraints, under the assumption that the jets are fairly steady.  The results for M87 and Perseus indicate that the radio emitting region of the jet is electron-positron dominated, assuming that the minimum of the electron energy distribution, $\\gamma_{\\rm min}\\sim 1$. ",
        "introduction": "The radio synchrotron emission from jets associated with active galactic nuclei (AGN) indicate that they must contain relativistic leptons and magnetic fields.  The matter content of these jets, however, is still a mystery (e.g. \\citealt{Homan05}).  Assuming that these jets are electrically neutral, are they electron-positron (``light'') or electron-proton (``heavy'')?  This paper aims to determine whether the positively charged particles are positrons or protons, in the assumption that the magnetic field is not energetically dominant.  Previous attempts to determine the matter content of the jets have used Synchrotron Self-Absorption (SSA) arguments (e.g. \\citealt{CelottiACFjets,ReynoldsACFjets}).  Measurements of the core radio flux and the size of the emitting region lead to constraints on the number density of the relativistic particles responsible for the emission, $n$, and on the magnetic field present in the jet, $B$.  The effect of any protons in the jet cannot be determined from radiative information.  As a result the large scale properties of the source are required to place independent estimates on the total power.  In the past this has been difficult because the total energy in the radio lobes has not been known; equipartition between the energy in the magnetic field and the radio-emitting particles has been assumed in order to be able to estimate the energy present \\citep{Burbidge}.  With the advent of the {\\it Chandra} and {\\it XMM-Newton} X-ray observatories the interaction of radio sources in the centres of clusters of galaxies with the surrounding Intra-Cluster Medium (ICM) has been observed in great detail.  One of the first observed was the Perseus Cluster, where decrements in the X-ray emission from the thermal ICM were first seen by {\\it ROSAT} \\citep{Bohringer}.  These have been interpreted as ``bubbles'' blown in the ICM by the central radio source.  Strong support for this interpretation came with {\\it Chandra} imaging and spectral information \\citep{ACF_complex_PER00}. Subsequently many other examples of such bubbles have been found (e.g. Hydra A, \\citealp{McNamaraHydra00}; A2052, \\citealp{Blanton01}; A2199, \\citealp{JohnstoneA2199}; Centaurus, \\citealp{SandersCent02}).  As these radio lobes are embedded in the ICM, their energy content can be more accurately determined using arguments of pressure balance between the thermal ICM and the relativistic plasma within the bubbles, as well as bubble dynamics.  As a result, the connection between the kinetic luminosity of the jet and the energy in the bubble can be calculated. The potential of this approach is that it does not require assumptions on equipartition for the estimate of the jet power.  In this work we combine the information from X-ray and VLBI observations of radio sources embedded in clusters of galaxies, where clear bubbles are visible in the X-ray images.  We apply the calculations performed for the jet in M87 by \\citet{ReynoldsACFjets} to calculate the matter content of the jet present in this and other sources.  In essence the method determines the lepton content of the base of the jet from synchrotron emission and absorption, and the total particle energy content from power required to create the observed bubbles.  The comparison between the two, sets constraints on the particle content.  In Section \\ref{SynSelAbs} we outline the models and assumptions used in the calculations.  The source parameters and the results are presented in Sections \\ref{sourceparam} and \\ref{results}.  The assumed lepton energy distribution is discussed in Section \\ref{gamma_min}, and further complications of this method and comparisons with other methods are discussed in Section \\ref{discussion}.  The implications of this work on the particle content of radio lobes is investigated in Section \\ref{radio_particle}.  We use \\mbox{$H_0= 70 \\kmpspmpc$} throughout. ",
        "conclusions": "If the conditions observed at the base of the jet are typical over the timescale for bubble creation, and all the caveats discussed in the previous section hold, then the jets in M87 and Perseus are likely to be dominated by an electron-positron plasma.  Other radio sources in clusters analysed are too distant to afford the resolution necessary to measure the synchrotron flux from the base of the jet.  Despite the significant potential offered by estimates of the jet power from its interaction with the ICM, this method is currently limited to determining the matter content of extra-galactic radio jets in sources in sufficiently nearby clusters of galaxies.  A concerted effort in monitoring the cores of nearby AGN at a variety of frequencies will help in constraining the jet matter content.  Clear identification of the cores, measurements of the jet velocity and limits on the variability will aid in improving the parameters and assumptions involved in these calculations. Future perspectives include high accuracy measurement of resolved radio cores via the space VLBI, and higher frequencies observations (though the flux should not be dominated by optically thin emission).  A definitive result on the matter content of jets remains, however, elusive for now."
    },
    "0608/astro-ph0608185_arXiv.txt": {
        "abstract": "We perform Monte-Carlo calculations of the radiative transfer of Ly$\\alpha$ photons emitted by a source embedded in a neutral collapsing gas cloud. This represents a young galaxy or quasar during the early stages of the epoch of reionisation (EoR). After computing the Ly$\\alpha$ spectrum as function of radius and time, we find that the Ly$\\alpha$ color temperature may be negative in large volumes surrounding the central source. Motivated by this result, we explore the prospects for a population inversion in the hyperfine levels of atomic hydrogen via the Wouthuysen-Field (WF) effect. The reason for this exploration is clear: if 21-cm masers exist during the EoR, they could greatly boost the expected 21-cm flux from this epoch. We find that population inversion is unlikely to occur for several reasons: (1) the required Ly$\\alpha$ luminosities of the central source exceed $\\sim 10^{45}$ erg s$^{-1}$. The radiation pressure exerted by such a large Ly$\\alpha$ flux likely halts the collapse of the cloud; (2) When quantum corrections to the WF-coupling strength are applied, the required Ly$\\alpha$ luminosities are (even) larger by orders of magnitude; (3) A relatively low flux of Ly$\\alpha$ photons that is produced via other channels (x-ray heating, collisional excitation of hydrogen) prevents the Ly$\\alpha$ color temperature from becoming negative. ",
        "introduction": "\\label{sec:intro}  Detecting the redshifted HI 21-cm line from the epoch of reionisation (EoR) is among the greatest observational challenges in present-day cosmology. Experiments such as the Low Frequency Array (LOFAR)\\footnote{www.lofar.org}, the Primeval Structure Telescope \\citep[PAST,][]{Pe06} and the Murchinson Widefield Array (MWA)\\footnote{http://www.haystack.mit.edu/ast/arrays/mwa/Site/index.html}, are currently being developed with the goal of detecting this EoR 21-cm signal within the next few years.  The strength and sign of the 21-cm signal depend on the spin temperature of the hydrogen gas, $T_s$, which is defined through the ratio of atoms in the ground and excited states of the hyperfine transition in the electronic ground state of hydrogen, $n_2/n_1 \\equiv g_2/g_1$ exp$[-T_*/T_s]$, where the ground and excited states are denoted by the labels '1' and '2', respectively and $T_*=0.068$ K. If we denote the temperature of the Cosmic Microwave Background (CMB) by $T_{\\rm CMB}$, then neutral hydrogen gas can be seen in emission if $T_s > T_{\\rm CMB}$. When $T_s < T_{\\rm CMB}$, the neutral hydrogen gas can be seen in absorption against the CMB.  The distribution of atoms in the hyperfine levels is affected by collisions between atoms, absorption and re-emission of the Cosmic Microwave Background (CMB) photons, and absorption and re-emission of Ly$\\alpha$ photons \\citep{review}. These Ly$\\alpha$ photons were either emitted by nearby sources, produced as local recombination emission, or in cascades following photo-excitation by photons with energies $E\\in [10.2-13.6]$ eV \\citep{PF06a}. When the Ly$\\alpha$ scattering rate dominates the collisional and CMB scattering rate, the hydrogen spin temperature, $T_s$, approaches the Ly$\\alpha$ color temperature, $T_{\\alpha}$ \\citep{Wouthuysen52}. The color temperature has been shown to approach the gas temperature, $T_{\\alpha}\\rightarrow T_{\\rm gas}$,  because of the frequency redistribution that Ly$\\alpha$ photons experience while resonantly scattering through an optically thick medium \\citep{Field59}. Because of this 'Wouthuysen-Field' effect, it is commonly assumed that the presence of Ly$\\alpha$ photons strengthens the coupling of the spin to the gas temperature \\citep[for the most detailed calculation see][]{Hirata06} and that neutral hydrogen gas is detectable in emission.\\\\  In this paper, we examine whether there are physical conditions under which the WF-effect produces population inversion (i.e. $T_s <0$) in the hyperfine levels. This investigation is motivated by recent Monte-Carlo calculations of Ly$\\alpha$ radiative transfer, which show that Ly$\\alpha$ photons emerge from neutral collapsing gas clouds with a systematic blueshift \\citep[e.g.][]{Zheng02,Dijkstra06}. This blueshift was found to increase with Ly$\\alpha$ optical depth and with the infall speed of the gas \\citep{Dijkstra06}. In this paper we show that this may lead to a negative Ly$\\alpha$ color temperature, and potentially to a negative spin temperature. The possibility of population inversion via the WF-effect and the accompanying stimulated 21-cm emission has been addressed in previous work \\citep[e.g.][]{Va67,S67,V68,Kaplan79}, but (obviously) not in the light of the results of these recent radiative transfer calculations, and not in the context of the EoR \\footnote{\\citet{SN97} have investigated whether hydrogen masers can exist during EoR. However, they focused on hydrogen  $n\\alpha$ recombination lines, which correspond to the transitions of the form $n\\rightarrow n-1$, where $n$ is the principal quantum number of atomic hydrogen.  \\citet{SN97} found that maser lines may be produced for $n\\sim 120$ along the edges of early HII regions.}. Exploring in more detail whether $T_s<0$ barely needs justification: if high redshift 21-cm masers exist, they could boost the 21-cm EoR signal by orders of magnitude. Interestingly, \\cite{CS06} recently showed that scattering of Ly$\\alpha$ and Ly$\\beta$ photons by deuterium can cause its spin temperature to be negative. Technically, this produces a 92-cm maser, although the associated gain is too small to provide a dramatic effect.  Before proceeding, we point out that \\citet{Deguchi85} found that Ly$\\alpha$ scattering in expanding gas clouds results in $T_s \\rightarrow T_{\\rm gas}$. One expects the same to apply to a collapsing gas cloud. However, the analysis presented in \\citet{Deguchi85} focused on Ly$\\alpha$ that was generated throughout the cloud. On the other hand, the young pre-reionisation universe may contain bright Ly$\\alpha$ sources, such as young galaxies and/or (mini)-quasars, that are surrounded by cold, neutral, collapsing gas. In such scenarios the Ly$\\alpha$ emission is concentrated strongly towards the center of the collapsing cloud. This introduces an important difference with the calculations of \\citet{Deguchi85}. Photons that are emitted in the center of the cloud must propagate through a larger column of collapsing material. This causes more energy to be transferred from the infalling gas to the photons, which results in a larger overall blue-shift of the Ly$\\alpha$ line \\citep{Dijkstra06}, and we will show in this paper that this large blueshift is the main reason that the WF-effect may cause a population inversion.  The outline of this paper is as follows. In \\S~\\ref{sec:gain} we discuss the expected maser amplification associated with a given negative hydrogen spin temperature. In \\S~\\ref{sec:color} we derive an expression for the Ly$\\alpha$ color temperature, which can be applied to arbitrary Ly$\\alpha$ frequency distributions, and its relation to the hydrogen spin temperature. In \\S~\\ref{sec:collapse} we discuss an example of a Ly$\\alpha$ source surrounded by cold collapsing gas, in which the Ly$\\alpha$ color temperature is negative. This is followed in \\S~\\ref{sec:pressure}-\\S~\\ref{sec:sourcereq} by a discussion of the factors that determine whether 21-cm masers can actually exist inside such clouds, before presenting our conclusions in \\S~\\ref{sec:conclusion}. The parameters for the background cosmology used throughout this paper are $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$, $\\Omega_b=0.044$, and $h=0.7$ \\citep{Spergel03}. ",
        "conclusions": "\\label{sec:conclusion} We have performed Monte-Carlo calculations of the radiative transfer of Ly$\\alpha$ photons emitted by a source embedded in a neutral collapsing gas cloud and computed the Ly$\\alpha$ spectrum as function of radius and time. This source represents a young galaxy or quasar during the early stages of the epoch of reionisation (EoR). After deriving an expression for the Ly$\\alpha$ color temperature $T_{\\alpha}$ for arbitrary frequency distributions of Ly$\\alpha$ photons (Eq.~\\ref{eq:tcolor}), we show that neutral collapsing gas clouds surrounding a central Ly$\\alpha$ source may see a negative Ly$\\alpha$ color temperature (Fig~\\ref{fig:sim}) under certain circumstances.  We have investigated whether a population inversion in the hyperfine transition of atomic hydrogen via the Wouthuysen-Field effect can be established in such scenarios, thus producing a 21-cm maser. The reason for this investigation is clear: if 21-cm masers exist during the epoch of reionisation (EoR), this could greatly boost the expected 21-cm emissivity of neutral gas during the EoR. Unfortunately, we have found that in practice a population inversion does not occur for several reasons:  (1) The required Ly$\\alpha$ luminosities exceed $\\sim 10^{45}$ erg s$^{-1}$. Although this can be obtained with physically possible star formation rates ($\\dot{M}_*\\gsim 10^2M_{\\odot}$ yr$^{-1}$) the radiation pressure exerted by such a large Ly$\\alpha$ flux likely halts the collapse of the cloud (\\S~\\ref{sec:pressure}).  (2) Quantum interference can strongly reduce the strength of the WF-effect in the Ly$\\alpha$ line wings \\citep{Hirata06}. The probability that a spin-flip occurs when a Ly$\\alpha$ wing photon scatters is reduced by a factor of $Q=(x_{45}/x)^2$ (see Appendix~\\ref{app:qi}) when quantum interference is included. In the context of this paper, this factor can be as small as $Q=10^{-4}-10^{-6}$. Therefore, when this quantum correction to the WF-coupling strength is applied, the required Ly$\\alpha$ luminosities are larger by orders of magnitude.  (3) A key requirement for having $T_{\\alpha}<0$ is a very low level of Ly$\\alpha$ photons in the line center, as these would scatter and rapidly rearrange the level populations in the hyperfine transition such that $T_s \\rightarrow T_{\\rm gas}$. We discussed in \\S~\\ref{sec:sourcereq}, however, that Ly$\\alpha$ core photons can be injected at a given location via different channels including x-ray heating, collisional excitation and conversion of higher Lyman series photons. The negative impact of these newly produced Ly$\\alpha$ photons (both by x-rays and soft-UV photons) is only overcome when the central source has an extremely prominent Ly$\\alpha$ line with EW$\\gsim 10^6/Q$ \\AA, which is several orders of magnitude larger than that of currently known sources in the Universe.  Unless the high redshift Universe contains some truly remarkable sources that can overcome the three obstacles mentioned above simultaneously, the Wouthuysen-Field effect will not produce a population inversion in the 21-cm transition. Detecting the redshifted 21-cm from the epoch of reionisation will therefore remain among the greatest observational challenges in cosmology.   {\\bf Acknowledgments} We thank George Rybicki and Stuart Wyithe for useful discussions. MD would like to thank Harvard's Institute for Theory \\& Computation for its hospitality. We thank an anonymous referee for providing us with some useful references to earlier work on 21-cm Masers. MD is supported by the Australian Research Council. AL was supported in part by NASA grants NAG 5-1329 and NNG05GH54G."
    },
    "0608/astro-ph0608191_arXiv.txt": {
        "abstract": "An X-ray compact source was discovered with {\\it Chandra} in a supernova remnant (SNR) N23, located in the Large Magellanic Cloud. The compact source (CXOU J050552.3$-$680141) is seen in only the hard band ($>$ 2 keV) image of N23, while the soft band image ($<$ 2 keV) shows diffuse emission of the SNR, with an extent of $\\sim 60 \\arcsec \\times 80\\arcsec$. The compact source is located at almost the center of N23, and there is no identifiable object for the source from previous observations at any other wavelength. The source spectrum is best explained by a power-law model with a photon index of $\\Gamma$ = 2.2 $^{+0.5}_{-0.3}$ and an absorption-corrected luminosity of $L_{x}$ = 1.0 $\\times$ 10$^{34}$ ergs s$^{-1}$ in the 0.5--10 keV band for a distance of 50 kpc. Neither pulsation nor time variability of the source was detected with this observation with a time resolution of 3.2 sec. These results correspond with those of \\citet{hughes06}, who carried out analysis independently around the same time as our work.  Based on information from the best-fit power-law model, we suggest that the source emission is most likely from a rotation-powered pulsar and/or a pulsar wind nebula. It is generally inferred that the progenitor of N23 is a core-collapsed massive star. ",
        "introduction": "The Large Magellanic Cloud (LMC) quite well meet the requirements of a systematic study of supernova remnants (SNRs), in view of the facts concerning the known distance (50 kpc; Feast 1999) and the small absorption because it is the nearest face-on galaxy to us. Over 30 SNRs have been discovered in the LMC so far, and categorized \\citep{williams99,hughes98}. Besides, due to the limited spatial resolution of previous observations, it is reported that only six of pulsar/SNR associations in the LMC, PSR B0540$-$69 in N158A \\citep{gotthelf00}, PSR J0537$-$6910 in N157B \\citep{marshall98}, a pulsar in N206 \\citep{klinger02}, SGR 0526$-$66 in N49 \\citep{kulkarni03,rothschild94}, XMMU J053559.3$-$673509 in DEM L241 \\citep{bamba06}, and a pulsar wind nebula (PWN) in SNR B0453$-$68.5 \\citep{gaensler03a}. A PWN is a synchrotron emission from high pressure plasma swept by the pulsar's relativistic wind.  Even though all of PWNe are powered by a pulsar, there is a case of a PWN without detection of a pulsar \\citep{gaensler03a}. The pulsar/SNR or PWN/SNR associations provide meaningful information, such as the age of the systems, and the type of remnants.  The SNR N23 (0506$-$68.0) is one of the SNRs in the LMC that \\citet{milne80} identified as an SNR with radio observations at 5 and 14.7 GHz for the first time. It is classified by its morphology as being a middle-aged ($\\sim$3000 year) and an irregular-shaped SNR in \\citet{hughes98}. SNR N23 has not been well studied, so there is still no absolute solution for its explosion type (a core-collapse or a Ia). \\citet{cohen88} insisted the possibility that N23 is a remnant from a massive star SN, since a CO cloud with high-density OB stars was observed within $\\sim$30 pc of N23. On the other hand, \\citet{banas97} suggested that the remnant does not have any association with others according to their own observations with better spatial resolutions than that of the Cohen survey.  We present the results of an X-ray observation of N23 with {\\it Chandra} in this paper. The observation details are summarized in $\\S$2, and the analysis results in $\\S$3. Using those results, we discuss the source discovered at the center of the remnant in $\\S$4, and $\\S$5 is a summary of this paper. \\citet{hughes06} reported the same results based on their independent analysis. ",
        "conclusions": "Although a physical association between Source and the SNR N23 is strongly supported by the fact that the Source location is almost at the center of the SNR, our fitting results of the $N{_{\\rm H}}^{\\rm LMC}$ do not determine whether Source is in LMC or not due to the error. We thus verified the case that Source is (1) an object in Galaxy, (2) the object in LMC is unrelated with SNR, (3) a background active galactic nucleus (AGN), (4) the stellar remnant of N23.  \\subsection{If Source is in Galaxy}  If Source is in Galaxy, possible candidates are an active RS CVn binary ($L_x$=29$-$32), a Cataclysmic Variables (CV; $L_x$=30$-$33), a Young Stellar Object (YSO; $L_x$=29$-$31), or a Low Mass X-ray Binary (LMXB; $L_x=$30$-$39), because their luminosities are consistent with the Source luminosity, assuming that Source is in Galaxy.  Since plasma models with any of three backgrounds did not provide an acceptable reduced $\\chi^2$, we concluded that Source is unlikely to be a RS CVn, a CV, or a YSO, which are well explained by plasma models. The emission from a LMXB can explained by a power-law model, but the photon index is typically 1--2 \\citep{muno03}, which is harder than the best-fit power-law model of Source.  As a result, we can say that Source would not be a considerable object in the case that Source is in Galaxy.  \\subsection{Source possibility of being a HMXB in the LMC}  If Source is in the LMC and unrelated with the SNR, it could be a HMXB.  The luminosity of HMXB (log$ L_x$ = 33--38; Muno et al. 2003) corresponds to the luminosity of Source, if Source is in the LMC. The observed spectrum of a HMXB is explained with a photon index of 0.5--2.5 \\citep{muno03}, which is also the same as the best-fit power-law of Source. However, the HMXB should be observed with the companion star in the optical band, and the expected optical limit of the companion star in the LMC is magnitude 16. The work by \\citet{hughes06} using the data of the MACHO Project indicates that the optical limit in the Source region is $>$ 17. We can thus reject the possibility of being a HMXB in LMC.  \\subsection{Source possibility of being an AGN}  The spectrum of an AGN has a photon index of $\\sim$1--3 \\citep{beckmann06}, and the photon index of Source is in between. We calculated the expected number of AGNs in the back of N23 using a log{\\it N}$-$log{\\it S} function measured by the {\\it ROSAT} deep survey in the Lockman Hole \\citep{hasinger98}. As a result, 24 AGNs, the flux of which is equal or higher than Source, are spread per square degree in the sky. Then, because the possibility of the AGN coincidence in the SNR region (elliptical; $\\sim 60\\arcsec \\times 80\\arcsec$) is only 0.8\\%, we can hardly find an AGN in the back of N23.  \\subsection{Source as a stellar remnant of SNR N23}  It is highly possible that Source associates with the SNR due to its location. The existence of an associating stellar remnant suggests that N23 is outcome of gravitational core-collapse of a massive star.  The possible candidates for Source as a stellar remnant of N23 would be a neutron star, a compact central object (CCO), an isolated pulsar, or a PWN.  The emission by a neutron star can be explained by a blackbody model with $kT$ = 0.15--0.25 keV \\citep{page96,page98}, but the blackbody model was not able to represent the Source radiation with any backgrounds. Thus, Source is hardly said to be a usual neutron star.  CCOs have recently been discovered in several SNRs, like the central source in Cassiopeia A \\citep{chakrabarty01} or in Puppis A \\citep{zavlin99}. CCOs' spectra can be fitted with either a blackbody or a power-law model, but the blackbody temperatures of CCOs are excessively high ($kT$ = $\\sim$0.3--0.6 keV; Chakrabarty et al. 2001, Kaspi \\& Roberts 2004) to consider that they are the usual neutron stars; also, the photon indices of power-law fitting is too soft ($\\Gamma$ = $\\sim$2.6--4.1; Kaspi \\& Roberts 2004) to be pulsars. The luminosity of CCOs is log{\\it L}$_x$ = 31--35 ergs s$^{-1}$, as shown by the triangle symbols in Figure~\\ref{fig:pi_lum}. Even though the Source luminosity is consistent with those of CCOs, Source could not be acceptably fitted with a blackbody model, and the photon index of a power-law model is harder than the typical photon index of CCOs. Thus, we consider that Source in N23 seems not to be a CCO.  A pulsar typically has the power-law spectrum with a photon index of $\\Gamma$ = $\\sim$0.6--1.9 \\citep{gotthelf00}, which is harder than the best-fit photon index of Source. Meanwhile, for a PWN, we would observe a softer photon index, generically, $\\Gamma$ = $\\sim$1.1--2.3 \\citep{cheng04}. The luminosity of the PWN is log$ L_x$ = 29--37, which is brighter than that of only the pulsar \\citep{cheng04}. This is indicated in Figure~\\ref{fig:pi_lum} by circle symbols. Due to the spectrum agreement, Source likely included the emission from a PWN. In this case, Source would be a rotation-powered pulsar plus a PWN, or a PWN powered by an unseen pulsar. For any conclusion, the pulsation searches with the better time resolution are required.  There is no counterpart of Source in the ATCA radio observation (1380, 2387, 4790, and 8640 MHz \\citep{dickel98}), and the flux limit was estimated as $\\sim$2 mJy in any of four frequencies. Source is quite similar to the radio-quiet PWN around PSR B1509$-$58, which flux in the X-ray band is factor of 2 lagrer than Source and the radio flux limit is $\\sim$1.3 mJy \\citep{gaensler02,gaensler99}. Therefore, the radio limit is consistent with our conclusion that Source is a radio-quiet PWN."
    },
    "0608/astro-ph0608644_arXiv.txt": {
        "abstract": "{Wolfgang Priester was one of Germany's most versatile and quixotic astrophysicists, re-inventing himself successively as a radio astronomer, space physicist and cosmologist, and making a lasting impact on each field. We focus in this personal account on his contributions to cosmology, where he will be most remembered for his association with quasars, his promotion of the idea of a nonsingular ``big bounce'' at the beginning of the current expansionary phase, and his recognition of the importance of dark energy (Einstein's cosmological constant $\\Lambda$) well before this became the standard paradigm in cosmology. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0608/astro-ph0608534_arXiv.txt": {
        "abstract": "{Near-infrared and optical spectroscopy of the nuclear and extended emission region of the Seyfert~2 galaxy Mrk~1210 is presented. This galaxy is outstanding because it displays signatures of recent circumnuclear star formation and a high-level of X-ray activity, in addition to the classical spectral characteristics typical of an AGN. The NIR nuclear spectrum, which covers the interval 0.8--2.4~$\\mu$m, is dominated by \\ion{H}{i} and \\ion{He}{i} recombination lines as well as [\\ion{S}{ii}], [\\ion{S}{iii}] and \\fe2 forbidden lines. Coronal lines of [\\ion{S}{viii}], [\\ion{S}{ix}], [\\ion{Si}{vi}], [\\ion{Si}{x}], and [\\ion{Ca}{viii}] in addition to molecular H$_{2}$ lines are also detected. Outside the nuclear region, extended emission of [\\ion{S}{iii}] and \\ion{He}{i} is found up to a distance of 500~pc from the center. An estimate of the contribution of the stellar population to the continuum is made by means of the $^{12}{\\rm CO(6-3)}$~1.618~$\\mu$m overtone bandhead. It was found that $83\\pm8\\%$ of the H-band continuum is of stellar origin. It improves previous estimates, which claimed that at least 50\\% of the observed continuum was attributed to the AGN. The analysis of the emission line profiles, both allowed and forbidden, shows a narrower (${\\rm FWHM} \\sim 500$~\\kms) line on top of a broad (${\\rm FWHM} > 1\\,000$~\\kms) blue-shifted component. The latter seems to be associated to a nuclear outflow. This hypothesis is supported by 6~cm VLBI observations, which show a radio ejecta extending of up to $\\sim$30~pc from the nucleus. This result does not required the presence of a hidden BLR claimed to be present in previous NIR observations of this object. Internal extinction, calculated by means of several indicators including [\\ion{Fe}{ii}] flux ratios not previously used before in AGNs, reveals a dusty AGN, while the extended regions are little or not affected by dust. Density and temperature for the NLR are calculated using optical and NIR lines diagnostic ratios. The results show electronic temperatures from 10\\,000~K up to 40\\,000~K and densities between $10^3 - 10^5$~cm$^{-3}$. The larger temperatures points out that shocks, most probably related to the radio outflow, must contribute to the line emission. ",
        "introduction": "Mrk~1210, also known as the Phoenix Galaxy, is a nearby Sa Seyfert~2 galaxy at $z=0.01350$. It has been extensively studied in the optical region, mainly because of the Wolf-Rayet features found within the central 200~pc (Storchi-Bergmann et al. 1998), indicating the presence of a circumnuclear starburst, and the detection of broad H$\\alpha$ and H$\\beta$ components in polarized light (Tran et al. 1992; Tran 1995). The widths of these last components reach 2380~km~s$^{-1}$ at FWHM. Near-infrared (NIR) spectra reported by Veilleux et al. (1997) show that the Pa$\\beta$ profile is characterized by a strong narrow component on top of a broad base with ${\\rm FWHM} \\sim 1600$~km~s$^{-1}$, suggesting the presence of a hidden broad line region (BLR). However, a comparison of the Pa$\\beta$ profile with that of \\fe2~1.256~\\micron\\ indicates that at least part of the broad emission is also present in the forbidden \\fe2\\ line. The moderate critical densities of the NIR \\fe2\\ lines, $n_\\mathrm{c} \\sim 10^{4-5}$~cm$^{-3}$ (Nisini et al. 2002), and the broad emission observed in \\fe2~1.256~\\micron\\ suggest that the broad permitted line is not produced in a genuine high-density BLR. The possible presence of a broad component to Pa$\\beta$ and Br$\\gamma$ and the absence of broad H$\\alpha$ in direct light would imply an ${\\rm E(B-V)_b} >$ 1.6 and 1.1, respectively, meaning that this object is strongly reddened.  Additional evidence of the dusty nature of Mrk~1210 is abundant. It is a member of the sample of galaxies called 60~$\\mu$m peakers (60PKs, Heisler \\& De Roberties 1999) because of its ``warm'' far-infrared color and the spectral energy distribution peaking near 60~$\\mu$m. Over the past few years, evidence has been accumulating indicating that these properties can be attributed to dust-obscured active galactic nucleus (e.g., Keel et al. 1994; Hes et al. 1995; Heisler \\& Vader, 1995). Whether or not this obscured material is the putative dusty torus as postulated by the unified model (e.g., Antonucci 1993) is highly debatable although evidence from spectropolarimetry points out to this scenario. In fact, Heisler et al. (1997) showed that Seyfert~2 galaxies for which polarized broad lines have been detected have warmer far-infrared colors ($f_{60}/f_{25} < 4$) compared to Seyfert~2 galaxies for which polarized broad lines have not been detected. The warm IRAS Seyferts are believed to be viewed at an angle that is more face-on than that of cool IRAS Seyferts. Moreover, in the near-infrared, the degree of polarization of Mrk~1210 rises toward longer wavelengths, reaching $\\sim 5$\\% at 2.15~$\\mu$m (Watanabe et al. 2003). Observational evidence suggest that this polarization is due to dichroic absorption by aligned grains. Watanabe et al. (2003) also suggest the existence of an additional unpolarized component, most probably associated to electron scattering.  The picture of a highly reddened nucleus derived from polarimetry is also strongly supported by optical multicolor imaging. Indeed, Heisler \\& Vader (1994) showed that Mrk~1210 displays a smooth global light distribution but morphologically resembles a typical elliptical galaxy. Masked frames display arcs or shells surrounding the main body of the galaxy. More recently, Martini et al. (2003) classified Mrk~1210 as a tightly wound nuclear dust spiral with the individual dust lanes traced for over a full rotation about the nucleus.  In X-rays, Mrk~1210 is peculiar. It is one of the very few cases of an AGN caught in a dramatic transition between a Compton-thick, reflection dominated state, and a Compton-thin state (Guainazzi et al. 2002). This result was possible after comparing XMM Newton observations with ASCA observations made six years earlier. The transition is attributed to either a change of the properties of the absorbers or a ``switching-off'' of the nucleus during the ASCA observations. In the former scenario, it may indicate a clumpy or patchy structure of the torus to which the absorbers are associated, while the later may provide clues about the duty-cycle of the AGNs phenomenon.  From all said above, Mrk~1210 is a key source where the interplay between an active nucleus and the presence of strong star formation can be studied in detail. Therefore, it is a very important object to investigate the nuclear structure of Seyfert~2 galaxies. In order to achieve this goal, it is necessary to characterize these two components (i.e., nuclear and stellar) as completely as possible. The presence of broad emission line components, a high degree of polarization, power law near infrared colors, warm far-infrared spectrum, high radio brightness temperature in the milli-arcsecond central structure and, finally, hard X-ray emission, suggest that a significant contribution comes from the active nucleus. This paper is the first attempt to study in Mrk~1210 the properties that can be attributed to the AGN, although not leaving aside the stellar contribution. Because of the strong evidence of dust obscuration, the NIR region is most appropriate because it attenuates only a fraction ($\\sim 1/7$) of the light compared to that in the optical, allowing to obtain an unbiased estimate of the obscured AGN.  Here, we will study Mrk~1210 by means of JHK spectroscopy covering the interval 0.8--2.4~$\\mu$m, complemented with optical spectroscopy, aimed at investigating {\\it (i)} the extinction affecting the nuclear and extended emitting gas; {\\it (ii)} the kinematics of the narrow line region (NLR) and {\\it (iii)} the physical properties and conditions of that gas. To meet the first goal, a variety of extinction diagnostic ratios based on permitted and forbidden lines, some of them never used before in an AGN, will be used. This will allow us to make a detailed map the dust distribution on that object. For the second goal, a detailed comparison of the emission line profiles, both in the optical and NIR will be carried out. It will also allow us to state whether or not a hidden broad line region is in fact present in the spectra of this object. For the third goal, density and temperature sensitive line ratios will be employed to construct a detailed picture of the ionization structure of emitting gas. In \\S ~\\ref{data} we describe the observations and data reduction process. \\S ~\\ref{features} describes the most important emission features found in the NIR spectra, some of them reported for the first time in this work.  \\S ~\\ref{kinematics} discusses the kinematics of the NLR gas while \\S ~\\ref{redd} evaluates and discusses the dust distribution in the nuclear and extended emission regions of Mrk~1210; \\S ~\\ref{physprop} explore the physical condition throughout the narrow line region and concluding remarks are in \\S ~\\ref{conclusions}. A Hubble constant of 75~km~s$^{-1}$~Mpc$^{-1}$ will be used in this work. ",
        "conclusions": "\\label{conclusions} We have carried out a detailed study of the physical properties of the nuclear and extended emission line regions of the Seyfert~2 galaxy Mrk~1210. For this purpose, we presented spectroscopic observations covering simultaneously the interval 0.8--2.4~$\\mu$m. The spectrum in the range 0.81--1.03~$\\mu$m is the first one published for this object. For that reason, all lines detected in this region were not previously identified in this source.  The spectra of Mrk~1210 display a plethora of emission lines, from those emitted from molecular H$_{2}$, prominent mostly in the K-band, along with low and medium ionization lines such as [\\ion{C}{i}], [\\ion{S}{ii}] and [\\ion{S}{iii}] and up to very high ionization lines of [\\ion{S}{ix}] and [\\ion{Si}{x}]. In addition, Mrk~1210 is characterized by a rich [\\ion{Fe}{ii}] spectrum, not previously reported in the literature for other AGN. Extended emission of [\\ion{S}{iii}], [\\ion{Fe}{ii}] and \\ion{He}{i} is found up to a distance of 500~pc from the nucleus.  The large amount of spectral information, at medium resolution, covering the optical and NIR regions allow us to study the physical conditions in the inner few hundred parsecs of this galaxy. Our main results can be summarized as follows:  {\\it (i)} From the depth of the $^{12}{\\rm CO(6-3)}$ overtone bandhead at 1.618~$\\mu$m, we have estimated that $83\\% \\pm 8\\%$  of the H-band continuum is of stellar origin. This value improves previous estimates, which claimed that at least 50\\% of the H-band continuum was attributed to the AGN and agrees with results found from optical spectroscopy. It also suggests the usefulness of the NIR in determining the percentage contribution of the stellar population to the integrated continuum emission. After correcting for the increased contribution of the stellar population to the continuum, Mrk~1210 continues displaying one of the highest star formation efficiency among water megamaser galaxies in spite of its relatively low molecular gas content. This result provides further support to previous findings showing that the molecular lines in Mrk~1210 were mainly excited by UV heating from stars.  {\\it (ii)} The analysis of the emission line profiles, both allowed and forbidden, shows that most prominent lines are characterized by a narrow (${\\rm FWHM} \\sim 500$~\\kms) component on top of a broad (${\\rm FWHM} > 1\\,000$~\\kms) blue-shifted component ($\\Delta{\\rm V} \\sim 150$~\\kms). The latter seems to be associated to a nuclear outflow, instead of the hidden BLR claimed to be present in previous NIR observations of this object. Moreover, the  differences in form and width of the line profiles of [\\ion{S}{iii}], [\\ion{Fe}{ii}] and H$_2$ imply that the emission seen from each line originates in different volumes of gas.  {\\it (iii)} We have examined the internal extinction affecting the NLR within the inner few hundred parsec using several indicators, including [\\ion{S}{ii}] and [\\ion{Fe}{ii}] line ratios. The results reveal a dusty AGN while the extended regions are very little or not affected by dust. In the inner 250~pc we found an ${\\rm E(B-V)}=0.5$ from the \\ion{H}{i} lines ratios. From the [\\ion{Fe}{ii}] line ratios, a larger extinction value, ${\\rm E(B-V)}=1.5$, was derived. This supports the hypothesis that the [\\ion{Fe}{ii}] lines are formed in a separate region, different from the partially ionized zone that exist in AGN. At 250~pc from the center, no extinction is found from our data.  {\\it (iv)} Electronic densities and temperatures for the nuclear gas of Mrk~1210 were determined by means of several diagnostic line ratios. The results found from the different indicators show that the NLR of Mrk~1210 cannot be characterized by a single temperature and density, since it presents a noticeable gradient in the physical conditions. The temperatures are from about 10\\,000~K up to the unusual value of 40\\,000~K determined from the [\\ion{S}{iii}] lines. This very high temperature, together with the high temperature derived from the [\\ion{O}{iii}] line ratio and the fact that these two lines display broad blueshifted components are evidence that shocks contribute to the ionization structure of the NLR of this galaxy. The range of NLR gas densities extends over three orders of magnitude and points out to a low density region ($n_\\mathrm{e}=1\\,000$~cm$^{-3}$) in which the [\\ion{S}{ii}]~\\l6716, \\l6730 lines (and probably the [\\ion{N}{i}] and [\\ion{O}{i}] lines) are emitted, up to a high density region ($n_\\mathrm{e}=10^{5.2}$~cm$^{-3}$) where the [\\ion{Fe}{ii}] lines are emitted."
    },
    "0608/gr-qc0608059_arXiv.txt": {
        "abstract": "Our  results concern  the transition  of a  quantum string through the singularity of the compactified Milne (CM) space. We restrict our analysis to the string  winding around the compact dimension (CD) of spacetime. The CD undergoes contraction to a point followed by re-expansion. We demonstrate that both classical and quantum dynamics of considered string are well defined. Most of presently available calculations strongly suggest that  the singularity of a time dependent  orbifold is useless as a model of the cosmological singularity. We believe that our results bring, to some extent, this claim into question. ",
        "introduction": "In the cyclic model of the evolution of the universe the cosmic singularity (CS) plays the key role because it joins  each two consecutive classical phases. We propose the following  basic criterion for the choice of the model of the neighborhood of the CS: reasonable model should allow for propagation of  quantum p-brane (i.e., particle, string, membrane,...)   from the pre-singularity to the post-singularity epoch. If the CS constitutes an insurmountable obstacle for  quantum elementary objects, the cyclic evolution cannot be realized.  One of the simplest models of the neighborhood of the CS, inspired by string/M theory, is the compactified Milne (CM) space. It was used recently \\cite{Khoury:2001bz,Steinhardt:2001vw,Steinhardt:2001st} in the cyclic universe scenario.  We have already applied the above criterion for the propagation of a test particle across the CS of the CM space \\cite{Malkiewicz:2005ii,Malkiewicz:2006wq}. The results we have obtained suggest that the CM space is a promising model deserving further investigation by making use of strings and membranes.  In this article we apply our criterion to the evolution of a test string in the CM space. We examine the propagation of a string in the so-called winding mode \\cite{Pioline:2003bs,Turok:2004gb}. It is defined to be a state in which a string is winding around the compact dimension undergoing contraction to a point followed by re-expansion. It turns out that the propagation is well defined both at classical and quantum levels.  In Sec. II we recall the classical formalism concerning the propagation of a test $p$-brane in a fixed background spacetime. Classical evolution of the string in the winding mode is presented in Sec. III. Quantization of the dynamics is carried out in Sec. IV. We conclude in Sec. V. ",
        "conclusions": "Our results  demonstrate that the dynamics of a {\\it test} string\\footnote{A test string, by definition, does not change the background spacetime.} winding around the compactified dimension in the {\\it zero-mode} state is well defined. At the classical level there are no instabilities near the singularity (characteristic for particle dynamics) and our non-perturbative method of quantization leads to a non-singular quantum evolution of a string\\footnote{ For the discussion of some physical aspects of propagation of winding states through the Milne space singularity we recommend \\cite{Russo:2003ky}. One shows that the simple string cosmology based on the Milne space can be used to obtain interesting cosmologies in four dimensions.}.  A great challenge is  examination of the dynamics of a {\\it physical} string, i.e. a string that may change its state and modify the background spacetime during the evolution of the entire system. The issue of the backreaction has been intensively studied recently (see \\cite{Craps:2006yb} for review). Calculations of the S-matrix amplitudes in  time dependent singular orbifolds are plagued by infinities (see, e.g. \\cite{Liu:2002kb,Lawrence:2002aj,Berkooz:2002je}). However, there exist some attempts to cure the divergencies of scattering amplitudes (see, e.g. \\cite{Cornalba:2003kd}). It seems to us that such results only mean that in these type of orbifolds the string {\\it perturbation} theory breaks down and one should apply {\\it non-perturbative} methods for analyses.  It is claimed in \\cite{Horowitz:2002mw} that the pessimistic results concerning the S-matrix can be confirmed by non-perturbative analyses.  One shows in \\cite{Horowitz:2002mw} that, for instance, a single particle added to a time dependent singular orbifold causes the orbifold to collapse into a large black hole, which leads finally to the creation of a big-crunch that would not be followed by a big-bang\\footnote{ See \\cite{Nekrasov:2002kf} for related considerations, where analyses of the bosonic and fermionic spectrum in the Milne orbifold lead the author to the conclusion that the resolution of the cosmological orbifold singularity is impossible.}. This result suggests that the compactified Milne space is useless as a model of the neighbourhood of the cosmological singularity in the context of the cyclic universe scenario.  It seems to us that the result of \\cite{Horowitz:2002mw} should be treated with caution. It is so because it was obtained by making use of general relativity, i.e. a classical theory, which  is not suitable for {\\it complete} understanding of the physics of a black hole.  In our article we have proposed a new non-perturbative method of analysis of the singularity of a time dependent orbifold. We believe that it can be further applied to analyses of the propagation of a winding string in {\\it non-zero} modes (excited states). It may initiate  a new direction in a systematic examination of the backreaction phenomena in the context of time dependent orbifolds.  We are conscious that {\\it real} understanding of the problem would require quantization of the entire system, including string and  orbifold, but it seems to be beyond the scope of the current stage of quantum gravity."
    },
    "0608/astro-ph0608702_arXiv.txt": {
        "abstract": "{} {We have measured the deuterium fractionation and the CO depletion factor (ratio between expected and observed CO abundance) in a sample of high-mass protostellar candidates, in order to understand whether the earliest evolutionary stages of high-mass stars have chemical characteristics similar  to those of low-mass ones. In fact, it has been found that low-mass starless cores on the verge of star formation have large values both of the column density ratio $N({\\rm N_2D^+})/N({\\rm N_2H^+})$ and of the CO depletion factor.} {We have observed with the IRAM-30m telescope and the JCMT two rotational lines of \\H\\ and \\D , the (2--1) line of \\CIII\\ and \\DCO , and the sub-millimeter continuum towards a sample of 10 high-mass protostellar candidates.} {We have detected \\D\\ emission in 7 of the 10 sources of our sample, and found an average value $N({\\rm N_2D^+})/N({\\rm N_2H^+})\\sim 0.015$. This value is $\\sim 3$ orders of magnitude larger than the interstellar D/H ratio, indicating the presence of cold and dense gas, in which the physical--chemical conditions are similar to those observed in low--mass pre--stellar cores. Also, the integrated CO depletion factors show that in the majority of the sources the expected CO abundances are larger than the observed values, with a median ratio of 3.2.} {In principle, the cold gas that gives origin to the \\D\\ emission can be the remnant of the massive molecular core in which the high-mass (proto-)star was born, not yet heated up by the central object. If so, our results indicate that the chemical properties of the clouds in which high-mass stars are born are similar to their low-mass counterparts. Alternatively, this cold gas can be located into one (or more) starless core (cores) near the protostellar object. Due to the poor angular resolution of our data, we cannot decide which is the correct scenario.} ",
        "introduction": "\\label{intro}  The initial conditions of the star formation process are still poorly understood. In recent years, studies of starless low-mass cores have begun to unveil the physical and chemical features that lead to the formation of low-mass stars (Kuiper et al.~\\cite{kuiper}; Caselli et al.~\\cite{caselli99}; Evans et al.~\\cite{evans}; Caselli et al.~\\cite{casellia};~\\cite{casellib}; Bergin et al.~\\cite{bergin02}; Tafalla et al.~\\cite{tafalla02};~\\cite{tafalla04};~\\cite{tafalla06}; Harvey et al.~\\cite{harvey}; Crapsi et al.~\\cite{crapsi}). On the other hand, the characterisation of the earliest stages of the formation of high-mass stars is more difficult than for low-mass objects, given their shorter evolutionary timescales, larger distances, and strong interaction with their environments.  In the last years various authors (Molinari et al.~2000,~2002; Sridharan et al.~2002; Beuther et al.~\\cite{beuther}; Fontani et al.~2005) have performed extensive studies aimed at the identification of precursors of ultracompact (UC) \\HII\\ regions, i.e. very young ($<10^{5}$ yr), massive (M$>8$\\solm ) objects which have not yet ionised the surrounding medium. In particular, from CS and mm continuum observations, it has been noted that in the earliest stages prior to the onset of massive star formation, the radial distribution of the intensity is quite flat, resembling the structure of starless cores in more quiescent and less massive molecular clouds (Beuther et al.~2002). More evolved objects show more centrally-peaked density structures, with power-law indices once again resembling those found in low-mass cores (i.e. \\dens$\\propto r^{-1.6}$; Shu 1977; Motte \\& Andr\\'e~2001). These results strongly suggest to apply the investigative techniques that are successful in the study of low-mass star forming cores, to the high-mass regime.  It has been found that when the starless core is on the verge of dynamical collapse, most of the high density tracers, including CS and the CO isotopologues, are frozen onto dust grains. CS observations of low-mass cores clearly show that this molecule is avoiding the central high density core nucleus, where the mm continuum peaks (Tafalla et al.~2002). The morphology of the \\CIII\\ (1--0) and (2--1) line emission of IRAS 23385+6053, a reliable example of precursor of an UC \\HII\\ region (Fontani et al.~2004a), suggests that CO is depleted in the high--density nucleus traced by the continuum emission. On the other hand, species such as NH$_3$, \\H\\ and \\D\\ are good tracers of the dust continuum emission (Tafalla et al. 2002; Caselli et al. 2002b; Crapsi et al. 2004, 2005), suggesting that they are not significantly affected by freeze-out (see Bergin et al.~2001 and Crapsi et al.~2004 for some evidence of \\H\\ depletion toward the center of B68 and L1521F, respectively). This is probably due to the fact that the parent species N$_2$, unlike CO, maintains a large gas phase abundance at densities around 10$^6$ cm$^{-3}$, despite the recent laboratory findings that N$_2$ and CO binding energies and sticking coefficients are quite similar (\\\"Oberg et al.~2005; Bisschop et al.~2006). The authors above mentioned also found a relation between the column density ratio D$_{\\rm frac}$=N(\\D )/N(\\H ) and the core evolution: D$_{\\rm frac}$ is predicted to increase when the core evolves towards the onset of star formation, and then it is expected to drop when the young stellar object starts to heat up its surroundings. Summarising, pre--stellar cores or cores associated with very young massive stars should show strong emission in the \\H\\ and \\D\\ lines, with large values of D$_{\\rm frac}$, and little or no emission of molecular species such as CO and CS, if they evolve like their low-mass counterparts.  Although recently Sridharan et al.~(2005) and Beltr\\'an et al.~(2006) have detected several candidate high-mass starless cores in the neighbourhood of some massive protostar candidate, a sample of well--identified high-mass pre--stellar cores is lacking. At present, the objects that are believed to be the closest to the earliest stages of the high-mass star formation process are those of two samples of high-mass protostellar candidates, selected from the IRAS-PSC by Molinari et al.~(\\cite{mol96}) and Sridharan et al.~(\\cite{sridharan}), and then investigated by various authors (Molinari et al.~\\cite{mol98},~\\cite{mol00}; Brand et al.~\\cite{brand}; Fontani et al.~2004a,b; Beuther et al.~\\cite{beuther}; Williams et al.~\\cite{williams}; Fuller et al.~\\cite{fuller}). These studies have shown that a large fraction of these sources are indeed very young massive objects. Also, assuming that high-mass stars form in clusters (Kurtz et al.~\\cite{kurtz}), we may expect to find other cold and dense cores located close to our target sources. For this reason, we have observed with the IRAM-30m telescope and SCUBA at JCMT a selected sample of candidate precursors of UC \\HII\\ regions in \\H , \\D , \\CIII\\ and sub-mm continuum. The sources have been selected from the two samples mentioned above on the basis of observational features indicative of very little evolution: they are associated with massive ($M\\sim 10^{2} M_{\\odot}$) and compact (diameters less than 0.1 pc) molecular cores, in which the kinetic temperatures are around $\\sim 20$ K, and show faint or no emission in the radio continuum. Also, they have distances less than 5 kpc.  Sect.~\\ref{obs} describes the observations and the data reduction, while the results are presented in Sect.~\\ref{res} and discussed in Sect.~\\ref{discu}. A summary of the main findings is given in Sect.~\\ref{conc}. ",
        "conclusions": "\\label{conc}  We have observed several rotational lines of \\H , \\D , \\CIII, \\DCO\\ and sub-mm continuum emission with the IRAM-30m telescope and the JCMT in 10 high-mass protostellar candidates. Our main goal is to measure the deuterium fractionation through the ratio $N({\\rm N_2D^+})/N({\\rm N_2H^+})$, and the CO depletion factor (ratio between 'expected' and observed CO abundance), in order to shed light on the chemical and physical properties of the molecular clouds in which high-mass stars are born. The main results of this study are the following: \\begin{itemize} \\item We have detected \\D\\ emission in 7 out of 10 sources of our sample, with values of $N({\\rm N_2D^+})/N({\\rm N_2H^+})$ between $\\sim 0.004$ and $\\sim 0.02$, with an average value of 0.015. This value is 3 orders of magnitude larger than the interstellar value for the relative D/H abundance ($\\sim 10^{-5}$), and close to that found by Crapsi et al.~(\\cite{crapsi}) in their sample of low-mass starless cores. \\item For almost all the observed sources the sub-mm continuum map and the maps of the \\H\\ (1--0) and (3--2) integrated line intensity show a similar distribution of the emission, reinforcing the idea that \\H\\ and sub-mm continuum trace the same material. \\item The CO 'expected' abundances are comparable to or marginally higher than the values measured from the observations in 8 sources, and much higher in 2 sources (\\vcvs\\ and \\vtqt ). The median value is 3.2, which is smaller than that found in low-mass pre--stellar cores. However, the non-depleted gas along the line-of-sight may strongly affect these results, since the source distances are larger than 1 kpc. We thus believe that the 'real' CO depletion factors estimated by us are lower limits. \\item A simple chemical model suggests that the majority of the observed sources have conditions similar to low-mass pre--stellar cores (large CO depletion factors and deuterium fractionation). The observed deuterium fractionation and CO depletion can be reproduced if the gas and dust temperature are $\\sim 20$ K, i.e. close to that observed by \\AMM\\ or \\ace\\ observations and to the \\H\\ (1--0) excitation temperature. Although some of the studied sources (IRAS~20126+4104 and IRAS~19092+0841) may indeed host dense and cold clumps with physical and chemical characteristics identical to low-mass pre--stellar cores, on average the present observations are consistent with envelope material recently heated by the central massive star, suggesting that the studied sources may have experienced a pre--stellar phase with gas and dust temperatures  quite similar to those observed in the nearby low-mass cores ($\\sim 10$ K), despite the larger volume densities. It appears that large fractions of CO may be trapped in H$_2$O ice, although high angular resolution observations are needed to confirm this.  \\end{itemize}  Our results allow us to conclude that most of our sources host cold and dense gas with properties similar to those found in low--mass starless cores on the verge of the gravitational collapse. The problem to solve now is the location and the origin of this gas. We propose two scenarios: in the first one, the cold gas is distributed in an external shell not yet heated up by the high-mass protostellar object, a remnant of the parental massive starless core. In the second one, the cold gas is located in cold and dense cores close to the high-mass protostar but not associated with it. This scenario is supported by high--angular resolution observations of some of our objects, which have revealed their clumpy structure. Clearly, observations at higher angular resolutions will be very helpful in deciding which solution is correct.  \\newpage \\newpage \\newpage \\newpage"
    },
    "0608/astro-ph0608228_arXiv.txt": {
        "abstract": "{Clouds containing molecular dark matter in quantities relevant for star formation may exist in minihaloes of the type of cold dark matter included in many cosmological simulations or in the regions of some galaxies extending far beyond their currently known boundaries.} {We have systematically explored parameter space to identify conditions under which plane-parallel clouds contain sufficient column densities of molecular dark matter that they could be reservoirs for future star formation. Such clouds would be undetected or at least appear by current observational criteria to be uninteresting from the perspective of star formation.} {We use a time-dependent photon-dominated region code to produce theoretical models of the chemistry and emission arising in clouds for our chosen region of parameter space. We then select a subset of model clouds with levels of emission that are low enough to be undetectable or at least overlooked by current surveys.} {The existence of significant column densities of cold molecular dark matter requires that the background radiation field be several or more orders of magnitude weaker than that in the solar neighbourhood. Lower turbulent velocities and cosmic ray induced ionization rates than typically associated with molecular material within a kiloparsec of the Sun are also required for the molecular matter to be dark.} {We find that there is a large region within the parameter space that results in clouds that might contain a significant mass of molecular gas whilst remaining effectively undetectable or at least not particularly noticeable in surveys. We note briefly conditions under which molecular dark matter may contain a dynamically interesting mass.} ",
        "introduction": "The mapping of external galaxies in the 21\\,cm \\ion{H}{i} \\citep[e.g.][]{Noordermeer2005,Walter2005} and the 2.6\\,mm CO \\citep[e.g.][]{Helfer2003,Lundgren2004} lines provides information about the distributions of atomic and molecular gas in those galaxies. Many studies of these distributions are undertaken in the context of star formation. For example, one commonly addressed question concerns whether a column density of gas in excess of a local critical value leads to star formation \\citep{Kennicutt1989}.  Observations of far infrared fine structure lines indicate that in metal-poor regions, like those found in some dwarf galaxies, 2.6\\,mm CO observations alone lead to a significant underestimate of the mass in molecular gas \\citep{Madden1997,Madden2002}. This is consistent with the results of multiline CO studies showing that the 2.6\\,mm integrated antenna temperature to molecular hydrogen conversion factor in such regions is much larger than is appropriate for Galactic giant molecular clouds, and also with the known difficulty in detecting CO emissions in some dwarf and compact galaxies compared to large spiral galaxies \\citep[e.g.][]{Israel2005}. The possibility that molecular gas may be present in large quantities in such regions has been postulated \\citep[see, e.g.,][and references therein]{Pfenniger1994,Combes2000,Kalberla2000}. Previous authors have concentrated on clumps with extents of the order of 10\\,AU and masses very much smaller than a solar mass each. Such clumps would have a covering factor (or beam filling factor) that is much smaller than unity. Their existence might have a range of observational consequences. For instance, they may be responsible for ``extreme scattering events'', causing large variations in radio signals from more distant sources \\citep{Walker1998}. We focus on lower density, larger objects having a covering factor near unity.  Of course, no survey will lead to the detection of all neutral atomic or molecular gas in a galaxy. Sensitivity limits lead to mapped distributions having boundaries, even though the real gas distributions may be more extended. The purpose of the present investigation is the determination of parameters characterizing the gas and environment that could result in significant hidden quantities of gas, with column densities comparable to those often found by surveys to be associated with star formation. Such gas could provide a reservoir for eventual star formation which would in turn create conditions leading to the gas being more easily detected. In many cases, the observed regions will be sufficiently distant that individual clouds are not observed. We have assumed, in effect, that all clouds have the same total column density and that, on average, any line of sight will pass through one cloud.  Therefore, we consider the possible existence of molecular clouds that contain a significant column density of $\\mathrm{H_2}$ whilst producing sufficiently little emission that they are undetectable or at least ignored by those conducting most surveys. We produce theoretical models of the chemistry and emission arising in molecular clouds and conduct a parameter space search for physical conditions that might give rise to undetectable quantities of molecular gas. In Sect.~\\ref{Model} we describe the photon-dominated region (PDR) code used to calculate the chemical and emission properties of model clouds and outline the range of parameter space examined. In Sect.~\\ref{Constraints} we review the sensitivity limits of current telescopes appropriate to the detection of emission (in the IR/sub-mm/mm/radio) from molecular clouds and specify the constraints we impose for a cloud to be considered undetectable. The results of our parameter space search are presented in Sect.~\\ref{Results} and we summarize our findings in Sect.~\\ref{Conclusion}. ",
        "conclusions": "We have conducted a large study of the parameter space associated with regions where the radiation field is weak relative to that in the solar vicinity to search for conditions that give rise to significant quantities of molecular gas with emission strengths that are comparable to the lower limits considered in standard surveys. Having explored parameter space to identify the time-dependent cloud models that satisfy these limits, we selected those models in which a significant column density of molecular hydrogen, $N_\\mathrm{H_2}\\ge10^{21}\\,\\mathrm{cm^{-2}}$ was reached.  We find that there is a large region within the parameter space that meets these requirements, resulting in clouds that might contain a significant mass of molecular gas whilst remaining effectively undetectable or at least not particularly noticeable in surveys. This region is characterized by low turbulent velocities ($\\le$$0.1\\,\\mathrm{km\\,s^{-1}}$), low radiation fields ($\\le$$10^{-4}$\\,Draine), normal to low metallicities ($\\sim$1--$10^{-2}\\,Z_\\odot$) and intermediate pressures ($\\sim$$10^{3}$--$10^{4}\\,\\mathrm{cm^{-3}\\,K}$). Clouds forming from molecular gas under these conditions are capable of maintaining large amounts of $\\mathrm{H_2}$ without becoming detectable in emission. Initially atomic gas can also form significant column densities of molecular material with undetectable emission, but the time-scales required are much longer ($>$$10^{7}$\\,yr).  \\citet{Sternberg2002} conducted a study of PDRs in which they assumed the present metagalactic radiation field to be incident on clouds in work on minihaloes composed of cold dark matter of the variety considered in many cosmological studies and on compact high velocity \\ion{H}{i} clouds. As mentioned above, the radiation field that they assumed is of comparable strength at 100\\,nm to fields for which we found large amounts of dark molecular gas to exist. Hence, there is some overlap between their work and ours. However, their emphasis is on \\ion{H}{i} clouds detectable in 21\\,cm emission, whereas ours is on the possible existence of dark molecular gas with substantial enough column densities to serve as reservoirs for potential star formation. They address observations of \\ion{H}{i} clouds having total hydrogen nuclei column densities of an order of magnitude and more below those with which we are concerned. Our interest is in whether regions of much higher column density may harbour sufficient molecular gas to become sites of stellar birth while they remain hidden in current surveys. We also wonder whether molecular dark matter existing in regions far away from bright stars may lead to future star formation in some galaxies occurring significantly beyond the currently detected bounds of those galaxies.  Finally, we note that an $\\mathrm{H_2}$ column density of $10^{21}\\,\\mathrm{cm^{-2}}$ distributed evenly over a disc with a radius of 40\\,kpc would correspond to a mass of $10^{11}\\,M_\\odot$, sufficient to be of dynamical importance in some galaxies.  A means of placing limits on the product of the covering factor of the type of molecular dark matter we are considering and the mean surface area per galaxy over which such dark matter is distributed would be through millimetre-wave absorption observations of other galaxies similar to those made by \\citet{Liszt1998} to study diffuse clouds in our galaxy. Observations towards many background sources would almost certainly be required to set meaningful limits on that product, as a detection rate of about 4 optical QSO absorption systems per unit redshift at a redshift of $\\sim$1 implies that intervening galaxies must have radii of almost 100\\,kpc if they are randomly oriented discs \\citep[e.g.][]{Sargent1979}. \\citet{Zwaan2006} have argued that a redshift interval of $\\Delta z\\approx100$ has been covered in the CORALS survey of radio-selected quasars \\citep{Ellison2001}, whereas a $\\Delta z$ range on the order of 3300 near $z=0$ should be required to discover a high $\\mathrm{H_2}$ column density absorber at the present epoch, if $\\mathrm{H_2}$ has the sort of distributions in galaxies that are inferred from the BIMA SONG maps of CO in nearby galaxies \\citep{Helfer2003}."
    },
    "0608/astro-ph0608587_arXiv.txt": {
        "abstract": "In this paper we specifically studied one aspect of foreground primordial matter density perturbations: the relative gravitational time delay between a pair of light paths converging towards an observer and originating from two points on the last scattering surface separated by the physical scale of an acoustic oscillation. It is found that time delay biases the size of acoustic oscillations {\\it systematically} towards smaller angles, or larger harmonic numbers $\\ell$, i.e. the mean geometry as revealed by CMB light becomes that of an open Universe if $\\Omega=1$. Since the effect is second order, its standard deviation $\\delta\\ell/\\ell \\sim (\\delta\\Phi)^2$ where $(\\delta\\Phi)^2 \\sim 10^{-9}$ is the normalization of the primordial matter spectrum $P(k)$, the consequence is too numerically feeble to warrant a re-interpretion of WMAP data.  If, however, this normalization were increased to $\\delta\\Phi \\gtrsim 0.01$ which is still well within the perturbation limit, the shift in the positions of the acoustic peaks would have been substantial enough to implicate inflationary $\\Lambda$CDM cosmology.  Thus $\\Omega$ is not the only parameter (and by deduction inflation cannot be the only mechanism) of relevance to the understanding of {\\it observed} large scale geometry.  The physics that explains why $\\delta\\Phi$ is so small also plays a crucial role, but since this is a separate issue independent of inflation, might it be less artificial to look for an alternative solution to the flatness problem altogether? ",
        "introduction": "Among the so-called `secondary' physical mechanisms that re-process the CMB anisotropy, such as gravitational lensing, time delay, and the Sunyaev-Zel'dovich effect, time delay by foreground inhomogeneities in the matter distribution appears to be investigated least.  In the recent period a detailed treatment of the problem was provided by Hu \\& Cooray (2001, hereafter HC01), although the general framework for calculating gravitational perturbation effects on the CMB as published in the review of CMB lensing by Lewis \\& Challinor (2006) could also be employed to carry the study further.  In HC01 the authors found an infra-red logarithmic divergence in the variance of the {\\it absolute} (or total) time delay along a randomly chosen direction to the LSS. They `renormalized' this infinity by subtracting the contribution from the `monopole' term of the matter power spectrum, corresponding to the removal of a constant uniformly across the sky.  Nevertheless, the remaining (finite) quantity still carries the divergence trend.  More precisely the variance is dominated by long wavelength fluctuations, prompting HC01 to consider it as an effect of very large coherence length, $\\Delta \\ell \\approx$ 2 which does not affect our interpretation of the CMB anisotropy, because the net outcome is simply a gentle and arbitrary distortion of the spherical shape of the LSS, completely negligible over the size of one (or a few) cycles of CMB acoustic oscillations.  The coherence length inferred by HC01 should be viewed with some caution, however, because the large delay excursion of $\\sim$ 1 Mpc/c calculated there, which involved only the {\\it zeroth} order term, the path integral of the perturbing potential itself, stems from the part of the matter power spectrum which carries the scale-invariant Harrison-Zel'dovich dependence $P(k) \\sim k$, i.e. there is a danger that the large coherence length may simply be due to the infra-red divergence of the variance rather than any genuine physical scale in the matter spectrum.  To find the coherence length of relevance to the question of CMB anisotropy distortion, it is necessary to pursue the perturbation expansion to the next two orders.  Not only are both results free from divergences, but also it is only through a comparison of these two terms that the true coherence scale for variations in the {\\it relative} delay between two light paths separated by a small angle $\\theta$ would become transparent.  We shall find that this scale is defined by a physically significant parameter, viz. the characteristic wavenumber at which departures of $P(k)$ from the Harrison-Zel'dovich behavior occurs for the first time.  The consequence is that appreciable distortion of the LSS radius, with both amplitudes and wavelengths on par with the dimension of the primary acoustic oscillations at the time of last scattering, can exist in principle. ",
        "conclusions": ""
    },
    "0608/astro-ph0608064_arXiv.txt": {
        "abstract": "{} {The maximum-likelihood method is the standard approach to obtain model fits to observational data and the corresponding confidence regions. We investigate possible sources of bias in the log-likelihood function and its subsequent analysis, focusing on estimators of the inverse covariance matrix. Furthermore, we study under which circumstances the estimated covariance matrix is invertible.}  {We perform Monte-Carlo simulations to investigate the behaviour of estimators for the inverse covariance matrix, depending on the number of independent data sets and the number of variables of the data vectors. } {We find that the inverse of the maximum-likelihood estimator of the covariance is biased, the amount of bias depending on the ratio of the number of bins (data vector variables), $p$, to the number of data sets, $n$. This bias inevitably leads to an -- in extreme cases catastrophic -- underestimation of the size of confidence regions. We report on a method to remove this bias for the idealised case of Gaussian noise and statistically independent data vectors. Moreover, we demonstrate that marginalisation over parameters introduces a bias into the marginalised log-likelihood function. Measures of the sizes of confidence regions suffer from the same problem.  Furthermore, we give an analytic proof for the fact that the estimated covariance matrix is singular if $p>n$. } {} ",
        "introduction": "The maximum-likelihood method \\citep[e.g.][]{barlow} is common practice to obtain the best-fit parameters $\\vec{\\pi}_0$ and confidence regions from a measured data vector $\\vec{d}\\in \\mathbb{R}^{p}$ for a model $\\vec{m}(\\vec{\\pi})$. It usually consists of finding the maximum of the log-likelihood function \\begin{equation} \\label{logL} \\mathcal{L}(\\vec{d}|\\vec{\\pi})\\propto-\\frac{1}{2} \\left[\\vec{d}-\\vec{m}(\\vec{\\pi})\\right]^{\\mathrm t} \\tens{\\Sigma}^{-1} \\left[\\vec{d}-\\vec{m}(\\vec{\\pi})\\right]\\, , \\end{equation} where $\\vec{\\pi}$ is the parameter vector and a Gaussian distribution of the measurement errors is assumed. The confidence regions for the maximum-likelihood fit are then defined by the surfaces of constant $\\Delta \\mathcal{L} \\equiv \\mathcal{L}_{\\rm max}-\\mathcal{L}$, where $\\mathcal{L}_{\\rm max}$ is the maximum value of the log-likelihood function.  For the evaluation of the log-likelihood the population covariance matrix $\\tens{\\Sigma}$ and its inverse $\\tens{\\Sigma}^{-1}$ or estimates thereof are needed. In most cases, no exact analytical expression for $\\tens{\\Sigma}$ can be given, although numerous authors make use of analytical approximations. An example from the field of weak gravitational lensing is \\citet{CFHTLSdeep}, who use the Gaussian approximation to the covariance matrix of the shear correlation functions given by \\citet{schneiderCov}. Other possibilities are to estimate $\\tens{\\Sigma}$ from the data themselves \\citep[e.g.][]{GABODS6,budavari03} or to obtain it from a simulated data set whose properties are comparable to the original data \\citep[e.g.][]{szapudi_mono}. In the latter paper, the authors observed that the estimated covariance matrix becomes singular if $p$, the number of entries of the data vectors, exceeds the number of observations / simulated data vectors. As a remedy, they propose to use the Singular Value Decomposition \\citep[SVD, ][]{nr} to obtain a pseudo-inverse of the covariance matrix, but do not investigate the properties of the resulting estimate of $\\tens{\\Sigma}^{-1}$ in detail. In this paper, we prove analytically that the rank of the standard estimator of the covariance matrix cannot exceed the number of observations. We then point out that, even if this estimator is not singular, simple matrix inversion yields a biased estimator of $\\tens{\\Sigma}^{-1}$. This may, if not corrected for, cause a serious underestimate of the size of the confidence regions.\tThis problem has also been noticed by \\citet{HirataGGL} and \\citet{MandelbaumIntrinsic}, who use Monte-Carlo simulations to determine the correct confidence contours in cases where the covariance matrix is noisy. We report on the existence of a simpler method to remove this bias, which can be derived for Gaussian noise and statistically independent data vectors, and test the validity of this method when these assumptions are violated. ",
        "conclusions": "We have given a proof for the fact that the standard estimator of the covariance matrix (\\ref{Cml}) is singular for \\mbox{$p>n$} if the mean vector used in (\\ref{Cml}) is known, and for \\mbox{$p>n-1$} if the mean is estimated from the same data set as the covariance matrix. Furthermore, we noted that the inverse of the maximum-likelihood estimator of the covariance matrix is a biased estimator of the inverse population covariance matrix. This bias depends basically on the ratio of the number of bins to the number of independent observations and can be quite severe as these two numbers become comparable. If uncorrected for, it will lead to a significant underestimation of the size of confidence regions derived from maximum-likelihood fits. The bias can be corrected for \\mbox{$p<n-2$} by using the estimator (\\ref{invCunbiased}) instead, which was derived by \\citet{TWA} under the assumption of Gaussian errors and statistically independent data vectors. We stress that there is no contradiction between the foregoing two statements: The singularity of $\\hat{\\tens{C}}_{\\rm ML}$ for \\mbox{$p>n-1$} derives from linear algebra alone, whereas the fact that $\\hat{\\tens{C}}^{-1}$ is zero for \\mbox{$p=n-2$} is due to the statistical distribution of the covariance matrix. Going beyond \\mbox{$p=n-1$}, we find that it is not advisable to use the Singular Value Decomposition to invert the estimated covariance matrix, since the bias of the pseudo-inverse does not seem to be controllable and depends strongly on the population covariance matrix, which is a-priori unknown.  Given the unbiased estimator $\\hat{\\tens{C}}^{-1}$, we argue that also the log-likelihood function is unbiased. However,  great care has to be taken if one wishes to perform further analysis of $\\mathcal{L}$: since it is a statistical variable, any nonlinear operation on it has the potential to cause a bias. We demonstrate this for the case of marginalisation over certain parameters, where the bias is relatively mild for the examples we chose. The situation is much worse if one tries to quantify the size of the confidence regions. The square root of the determinant of the inverse Fisher matrix shows a significant amount of bias, and therefore should not be used as absolute measures of measurement uncertainty.  The upshot of all this is the following: avoid to use more bins for your likelihood fit than you have realisations of your data. If your errors are Gaussian and the data vectors are statistically independent, use the estimator $\\hat{\\tens{C}}^{-1}$ to obtain an unbiased estimate of the inverse covariance matrix and the log-likelihood function. If one or both of these two requirements are not fulfilled, the estimator is not guaranteed to work satisfactorily; this should be checked from case to case.  Finally, we note that the estimates of the covariance matrix and its inverse can be quite noisy. If one has prior knowledge about the structure of the covariance matrix, one can develop estimators with a much lower noise level. Since this noise is responsible for most of the problems discussed in this paper, these improved estimators may also be useful in situations where the requirements for the use of $\\hat{\\tens{C}}^{-1}$ are not fulfilled. We will explore these possibilities in a future paper."
    },
    "0608/astro-ph0608252_arXiv.txt": {
        "abstract": "We present calculations of the reflection of the cosmic X-ray background (CXB) by the Earth's atmosphere in the 1--1000 keV energy range. The calculations include Compton scattering and X-ray fluorescent emission and are based on a realistic chemical composition of the atmosphere. Such calculations are relevant for CXB studies using the Earth as an obscuring screen (as was recently done by INTEGRAL). The Earth's reflectivity is further compared with that of the Sun and the Moon -- the two other objects in the Solar system subtending a large solid angle on the sky, as needed for CXB studies. ",
        "introduction": "Having a mass column density of $\\sim 10^3 {\\rm g~cm^{-2}}$ at the sea level the Earth's atmosphere completely blocks the X-rays from celestial sources. At the same time the outer layers of the Earth's atmosphere reflect part of the incident X-ray photons due to Compton scattering. The physical picture is very similar to the well-studied case of the reflection from a star surface (e.g. Basko, Sunyaev \\& Titarchuk, 1974) or an accretion disk (e.g. George \\& Fabian, 1991) except for the different chemical composition of the reflecting medium.  The reflection of X-rays by the Earth's atmosphere is important for e.g. evaluating the echo produced by gamma-ray bursts (Willis et al., 2005) or for studies of the cosmic X-ray background (CXB). The present study was particularly initiated by recent observations of the Earth with the INTEGRAL observatory (Churazov et al., 2007) aimed at determining the CXB intensity near the peak of its luminosity distribution (i.e. around 30--40 keV). In these observations the Earth disk blocked the X-rays coming from distant objects, causing a decrease in the observed flux. The energy dependent reflection of X-rays by the Earth's atmosphere reduces the amplitude of this decrement. The purpose of this paper is to provide a simple recipe for the calculation of the reflected flux in the energy range from few to few hundred keV.  Since photons change their energy during Compton scattering, the reflected emission at a given energy depends on the overall shape of the input spectrum. Generally, one needs to calculate a Green function describing the reflection of monochromatic radiation (White, Lightman \\& Zdziarski 1988, Poutanen, Nagendra \\& Svensson 1996) and convolve the incident spectrum with this function to calculate the reflected spectrum. We instead calculate the energy dependent ``albedo'' $A(E)$ -- the ratio of the flux reflected by the Earth disk at a given energy to the flux of the incident spectrum at the same energy, using the canonical approximation of the CXB spectrum (Gruber et al., 1999) as the incident spectrum. The resulting function $A(E)$ may be used for calculating the reflected emission for an incident spectrum that has a shape similar to that assumed in our simulations. We also estimate the uncertainties introduced in $A(E)$ by variations in the input spectral shape.  Two other bodies in the Solar system subtend a large solid angle on the sky for telescopes in near-Earth orbits -- the Moon and the Sun. We therefore also calculate the X-ray albedo for them and compare it with the reflectivity of the Earth's atmosphere. ",
        "conclusions": "We calculated the Earth atmospheric albedo for the CXB radiation in the 1--1000 keV energy range. An analytic approximation for the angle averaged albedo is provided, which is especially useful when the whole Earth disk is observed. These calculations (along with the calculations of the cosmic ray induced atmospheric emission by Sazonov et al., 2007) were used in the analysis of the INTEGRAL observations of the Earth (Churazov et al., 2007) and a good agreement between predictions and measurements was found.  We further compared the Earth albedo to the CXB albedo for the Sun and the Moon and discuss other components contributing to the X-ray emission of these objects in the 1-1000 keV band.  The night-side of the Moon should be the X-ray darkest object in the solar system (subtending substantial solid angle for the telescope at the Earth orbit) in the energy range 1-30 keV, except at the energies of fluorescent lines of (e.g. Si or Fe).  The reflectivity of the Sun is on the contrary the highest and for exceptionally strong gamma-ray bursts or other transient events the signal reflected by the Sun below 100 keV (at the level of few $10^{-6}$ of the direct signal) will be stronger than by the Moon (see Mazets et al., 2005, Frederiks et al., 2007 for the discussion of a recent outburst from SGR1806-20 reflected by the Moon).  The Earth, because of the large solid angle subtended by its disk (for a wide-field telescope on the Earth orbit) is the most useful object for the CXB obscuration studies and potentially for the flux calibration of hard X-ray and gamma-ray missions at energies higher than 50-100 keV."
    },
    "0608/astro-ph0608408_arXiv.txt": {
        "abstract": "The galaxy cluster 1E0657-56 ($z = 0.296$) is remarkably well-suited for addressing outstanding issues in both galaxy evolution and fundamental physics. We present a reconstruction of the mass distribution from both strong and weak gravitational lensing data. Multi-color, high-resolution HST ACS images allow detection of many more arc candidates than were previously known, especially around the subcluster. Using the known redshift of one of the multiply imaged systems, we determine the remaining source redshifts using the predictive power of the strong lens model. Combining this information with shape measurements of ``weakly'' lensed sources, we derive a high-resolution, absolutely-calibrated mass map, using no assumptions regarding the physical properties of the underlying cluster potential. This map provides the best available quantification of the total mass of the central part of the cluster. We also confirm the result from \\citet{clowe04,clowe06b} that the total mass does {\\it not} trace the baryonic mass. ",
        "introduction": "}  The cluster of galaxies {\\bullet} is one of the hottest, most X-ray luminous clusters known. Since its discovery by \\citet{tucker95}, it has been the subject of intense and ongoing research \\citep{mehlert01,markevitch02,barrena02,clowe04,markevitch04,gomez04}. In particular, Chandra observations by \\citet{markevitch02} revealed the cluster to be a supersonic merger in the plane of the sky with a textbook example of a bow shock, making this cluster a unique case in which to study hydrodynamical properties of interacting systems. The optical images show that the cluster has two distinct components, and the X-ray analysis reveals that the lower mass sub-cluster has recently exited the core of the main cluster with a relative velocity of $4500^{+1100}_{-800} \\kms$. Although this relative velocity appears to be unusually large, an analysis of cosmological simulations demonstrates that it is well within the predicted range of the currently favored cosmological model \\citep{hayashi06}.   Due to its unique geometry and physical state, this cluster is the best known system in which to test the dark matter hypothesis \\citep{clowe04}. The observed offsets between the weak gravitational lensing mass peaks and the X-ray gas component give the most direct evidence for the presence of dark matter yet available. Using the same observations, \\citet{markevitch04} placed upper limits on the dark matter self-interaction cross section.   The goal of this work is to obtain a high-resolution, absolutely calibrated mass map with no assumptions on the physical properties of the underlying cluster potential. For this purpose we use HST ACS data, incorporating the gravitational lensing information from both multiple image systems (strong lensing) and from distortions of background sources (weak lensing). The superb spatial resolution delivered by ACS both increases the number density of background sources that can be used for weak lensing and reveals new strong lensing candidates, especially around the subcluster. The joint strong and weak lensing analysis thus benefits greatly from the ACS data, enabling us to increase the spatial resolution and signal strength of the mass map in the core of the cluster.   Encouraged by the success of the combined strong and weak lensing reconstruction method developed in \\citet{bradac04a} (hereafter \\citetalias{bradac04a}) and the results from applying the method to cluster RX~J1347.5$-$1145 \\citep{bradac04b}, we proceed to apply this method to {\\bullet}. Because the strong lensing data are richer in this case (RX~J1347.5$-$1145 has not been observed with either the ACS or WFPC2 cameras prior to this study), we improve the method as described in the text.    We perform the reconstruction in the following sequence. Using ACS (multi-color where available) HST images, we identify the multiply imaged systems. Having a spectroscopic redshift for one lensed system (from \\citealp{mehlert01}), we then use the predictive power of simple strong lens modeling to estimate the redshifts of the other systems. Using positions and redshifts of the strongly lensed images and the shape measurements of the ``weakly'' lensing sources, we perform a combined strong and weak lensing mass reconstruction, thereby significantly improving the constraints on the mass and positions of the main cluster and the colliding subcluster of {\\bullet}. In terms of the structure of this paper, in section~\\ref{sec:data} we describe the optical images used in this analysis, the basic image processing and the extraction of the strong~(\\S\\ref{sec:datasl}) and weak~(\\S\\ref{sec:datawl}) gravitational lensing data. We then infer the mass distribution of cluster {\\bullet} from these data in section~\\ref{sec:swunited_bullet}, following a demonstration of our methodology on suitable simulated data in section~\\ref{sec:swunited}. We discuss the possible sources of error in section~\\ref{sec:systematics} and summarize our conclusions in section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  Massive and interacting clusters, while quite rare, are remarkably well-suited to addressing outstanding issues in both galaxy evolution and fundamental physics. However, in order to study the mass distribution, methods relying on hydrostatic (X-rays) or dynamical equilibrium are ill-suited for such systems.   We have applied a mass reconstruction method based on strong and weak gravitational lensing to the cluster {\\bullet}. We use deep, high resolution optical data to identify objects belonging to the same multiple-image systems.  The same data are used to obtain weak lensing catalogs allowing us to obtain a strong+weak lensing mass map of the cluster core.  Our main conclusions are the following: \\begin{enumerate} \\item Using the combined strong and weak lensing mass reconstruction we derive a high-resolution, absolutely calibrated mass map; we get projected, enclosed mass $M_{\\rm main}(<250\\:\\mbox{kpc})= (2.8 \\pm 0.2) \\times 10^{14} M_{\\odot}$ around the main and $M_{\\rm sub}(<250\\:\\mbox{kpc})= (2.3 \\pm 0.2) \\times 10^{14} M_{\\odot}$ around the subcluster. \\item We detect the main cluster peak and a distinct mass concentration at the subcluster position, both clearly offset from the location of the X-ray gas in the system (at 10-$\\sigma$ and 6-$\\sigma$ significance for the main and the subcluster respectively - see Fig.~\\ref{fig:swunited} and Tab.~\\ref{tab:offset}). \\item The majority of the mass is spatially coincident with the galaxies, which implies that the cluster mass must be dominated by a relatively collisionless form of dark matter. The high resolution data allow us to significantly detect the shapes of both the main mass component and the subcluster with no prior assumptions on their positions or profiles. \\end{enumerate} We show that the cluster {\\bullet} is a very efficient strong lens; the area enclosed within the critical curve for highest redshift sources is comparable to that of A1689 \\cite{broadhurst05}. Consequently, we plan to conduct a pencil beam survey for high-redshift ($z\\sim 7$) objects. At present this is not possible, because we lack both reliable redshift information for the {\\bullet} multiple image systems, and multi-color high-resolution imaging data in the main cluster region. The latter is needed to identify counter images for the strong lens systems, and so improve the reliability of the mass map to the south-east of the main cluster.  Once these data are obtained (HST Cycle 15), we will improve our methodology to include an adaptive grid approach, which will greatly improve the constraints on the location of critical curves for these high-redshift sources.   The prospects for such a search are however very promising, as we already have two strong candidate high redshift objects. One is the red arc B, which is constrained to lie at $3.24 < z < 6$ based upon the redshift of arc A and detection in F814W. The other (see Fig.~\\ref{fig:ccurve}) is Spitzer IRAC-selected with a likely redshift of $z \\gtrsim 6$. Using this cluster as a gravitational telescope will therefore allow us to study galaxy properties near the epoch of reionization.   The high resolution mass map of this cluster, in addition to enabling the high-z galaxy search, will be valuable for several other ongoing science programs.  Most significantly, we are currently running simulations to reconstruct the dynamical history of the cluster merger, and provide an independent constraint on the dark matter self-interaction cross-section. The relative masses of the main and the subcluster, and the central mass distribution, two critical parameters for the simulations, were the limiting source of uncertainty prior to this work."
    },
    "0608/astro-ph0608122_arXiv.txt": {
        "abstract": "Minor mergers of galaxies are expected to be common in a hierarchical cosmology such as $\\Lambda$CDM. Though less disruptive than major mergers, minor mergers are more frequent and thus have the potential to affect galactic structure significantly. In this paper we dissect the case-by-case outcome from a set of numerical simulations of a single satellite elliptical galaxy accreting onto a massive elliptical galaxy. We take care to explore cosmologically relevant orbital parameters and to set up realistic initial galaxy models that include all three relevant dynamical components: dark matter halos, stellar bulges, and central massive black holes.  The effects of several different parameters are considered, including orbital energy and angular momentum, satellite density and inner density profile, satellite-to-host mass ratio, and presence of a black hole at the center of the host.  Black holes play a crucial role in protecting the shallow stellar cores of the hosts, as satellites merging onto a host with a central black hole are more strongly disrupted than those merging onto hosts without black holes. Orbital parameters play an important role in determining the degree of disruption: satellites on less bound or more eccentric orbits are more easily destroyed than those on more bound or more circular orbits as a result of an increased number of pericentric passages and greater cumulative effects of gravitational shocking and tidal stripping.  In addition, satellites with densities typical of faint elliptical galaxies are disrupted relatively easily, while denser satellites can survive much better in the tidal field of the host.  Over the range of parameters explored, we find that the accretion of a single satellite elliptical galaxy can result in a broad variety of changes, in both signs, in the surface brightness profile and color of the central part of an elliptical galaxy.  Our results show that detailed properties of the stellar components of merging satellites can strongly affect the properties of the remnants. ",
        "introduction": "There now exists a strong body of evidence that luminous elliptical galaxies are, in many ways, a distinct group of objects when compared to lower-luminosity elliptical galaxies.  This dichotomy is seen not only in their global structure -- for example, isophotal shapes \\citep{kormendy1996} and degree of rotational support \\citep{davies1983, bender1992} -- but also in their central properties: giant ellipticals ($M_v<-22$) have shallow central surface brightness profiles ($-d\\ln \\Sigma/d\\ln R < 0.26$), little rotation but velocity anisotropy, and boxy isophotal shapes, whereas spiral bulges and fainter ellipticals ($M_v>-20.5$) exhibit central power-law surface brightness profiles ($-d\\ln \\Sigma/d\\ln R > 0.5$), significant rotational support, and disky isophotal shapes \\citep{lauer1995, faber1997, ravindranath2001, lauer2005, ferrarese2006}.  Different environments and formation histories are plausibly responsible for these two classes of galaxies.  The fainter ellipticals are found in all environments including the field; their dense power-law centers may be a signature of dissipation in gas-rich mergers of disk galaxies (e.g., \\citealt{faber1997, genzel2001}).  The luminous ellipticals, on the other hand, are predominately found in galaxy clusters, some as brightest cluster galaxies and cD galaxies at or near the centers of clusters.  Their shallow central luminosity profiles and low-density cores appear particularly difficult to form and preserve in theoretical models.  Black holes provide an attractive model for the initial formation of the cores: mergers of two cuspy ($\\rho \\propto r^{-2}$) stellar bulges with central black holes result in the sinking of the black holes to the center of the merger remnant by dynamical friction and the formation of a black hole binary \\citep{begelman1980}.  This process, and the subsequent orbital decay (``hardening'') of the binary, result in energy transfer from the black holes to the stellar system, and the combined effect is to reduce the initial cusp to a profile that is nearly cored in projection \\citep{ebisuzaki1991, quinlan1997,milosavljevic2001}.  How the shallow surface brightness profiles of luminous ellipticals evolve and are preserved (for the most part) in subsequent mergers, however, is an open question and is the focus of this paper.  The stellar mass within the core radius (defined to be the radius where the local logarithmic slope of the surface brightness is $-0.5$) is of order 1\\% or less of the total stellar mass; a small amount of mass deposit or removal can therefore easily alter the inner profiles.  To maintain such shallow stellar profiles, massive elliptical galaxies, once formed, must somehow prevent additional stellar mass from being brought into their central regions through the frequent satellite accretion that is a hallmark of cold dark matter-based cosmological models.  In particular, cosmological $N$-body simulations that model only dark matter have revealed a rich spectrum of dense dark matter subhalos, some of which can survive complete tidal disruption and orbit within their host halos.  A logical expectation is that some of the massive core ellipticals have accreted at least one lower-mass galaxy with cuspier central density over their lifetimes.  It is therefore intriguing that surface brightness cores in luminous ellipticals are ubiquitous.  In this paper, we report the results from a set of merger simulations that explore the effects of a single satellite accreting onto a massive elliptical galaxy.  Our simulations use realistic galaxy models composed of all three components that are dynamically relevant for elliptical galaxies: dark matter halos, stellar spheroids, and central supermassive black holes (section~\\ref{sec:sims}).  As we show, each component plays an important role in the outcome of the accretion events and therefore should be included.  The dark matter halos set the gravitational potential well on large scales and determine the orbital evolution in the early stages of the merger (before the stellar bulges interact strongly).  The bulges dominate the gravitational potential on smaller ($\\sim$ kpc) scales and thus affect the orbital evolution at late times in the merger.  On sub-kpc scales, supermassive black holes can dominate the potential well and thus contribute to the structural evolution of satellites, sometimes quite strongly.  Even though the simulations are dissipationless, the three components exchange energy gravitationally with one another during the merger, leading to observable effects on the stellar component that would not be fully captured if any component were neglected.  An additional feature of our simulations is that we choose the orbital parameters based on distributions of orbits determined from large cosmological dark-matter-only simulations.  By applying a semi-analytic prescription of dynamical friction and tidal stripping (section~2.3), we are able to start our simulations not at the virial radius of the host, where the cosmological distributions of halo merger orbits are determined, but at the smaller scale of four effective radii of the host's bulge.  Our analysis focuses broadly on two aspects of satellite accretion: orbital decay and structural evolution of the satellite during the merger process (section~\\ref{sec:fate}) and observable signatures of infalling satellites on the host galaxies (section~\\ref{sec:observables}).  For the former, we investigate the dependence of an accreting satellite's fate on the initial orbital energy and angular momentum (section~3.1), the satellite's density (section~3.2), the presence of central black holes (section~3.3) and dark matter halo (section~3.4), and other factors such as the satellite-to-host mass ratio and the inner density profile of the satellite (section~3.5). For the observable consequences of satellite accretion, we investigate a number of different properties including surface brightness profiles (section~4.1), colors and color gradients (section~4.2), and mass deposit at large radii and kinematic structures (section~4.3). ",
        "conclusions": "In this paper we have used dissipationless simulations to examine in detail and over a wide range of parameter space the case-by-case outcome of a single satellite elliptical galaxy accreting onto a core elliptical galaxy. We have taken special care to explore cosmologically relevant orbital parameters and to set up realistic galaxy models that include all three relevant dynamical components: dark matter halos, stellar bulges, and central massive black holes.  The main results in this paper are as follows:  (i) Satellites on more energetic or more eccentric (lower angular momentum) orbits are destroyed more efficiently due to the stronger tidal effects from the host spheroid and the central black hole (see Figs.~2, 3, 9-12).  (ii) Denser satellite galaxies -- those with either a smaller $\\re$ or a steeper inner slope -- survive significantly better (see Figs.~\\ref{fig:orbitDense}, 5, 9-12).  For instance, the standard density satellite in Fig.~\\ref{fig:orbitDense} loses more mass by its 3rd pericentric passage than the total mass loss in the $2\\sigma$ dense satellite during its entire orbital evolution.  (iii) A supermassive black hole at the center of the host galaxy significantly increases the efficiency with which an accreting satellite is destroyed (see Figs.~6, 7, 10).  Without a central black hole, luminous elliptical galaxies would have more difficulty in protecting their shallow inner profile (however formed initially) from steepening due to mass deposit through satellite accretion, in agreement with previous 2-component numerical work \\citep{holley-bockelmann2000, merritt2001}.  (iv) Ignoring dark matter halos in the galaxy models causes significant differences primarily because dark matter halos lead to higher orbital velocities for a fixed ratio of $\\vorb/\\vcn$ (see Fig.~8).  Satellites with increased orbital velocity complete more orbits before merging with the host, resulting in greater destruction of the satellite galaxy.  (v) Our overall conclusion is that the accretion of a single satellite elliptical galaxy (1/30 to 1/10 mass of the host) can result in a broad range of changes, in both signs, in the surface brightness profile $\\Sigma(R)$ and color of the central part of an elliptical galaxy, as summarized in Figs.~\\ref{fig:sb_all}-12.  More specifically, Figs.~\\ref{fig:sb_all}-12 show that when a black hole is included in the host galaxy and when the satellite and host are both on the SDSS stellar size-mass relation, the fractional change in $\\Sigma(R)$ in the inner $\\sim 500$ pc of the host galaxy upon satellite accretion ranges from $\\sim -20\\%$ to $\\sim +40\\%$ When more compact satellites are considered in accordance to the large scatter about the SDSS size-mass relation, the inner $\\Sigma(R)$ of the host galaxy can increase by more than 100\\%.  Since accretion of a single satellite adds little mass and luminosity to the host, we suggest that a ``core'' galaxy at $M_v\\sim -22$ can potentially transition into a galaxy with power-law characteristics at a comparable luminosity by accretion of a relatively compact and massive ($\\sim$ 1:10) satellite.  A core-in-core structure would provide particularly strong evidence for this process (see also \\citealt{kormendy1984}).  The strong dependence of satellite mass loss on the satellite's density relative to the host and the inclusion of black holes has broader implications, for example, in interpreting the connection between dark matter subhalos and galaxies in cosmological simulations that model only dark matter.  Dark matter subhalos in these simulations are often stripped to the point of not being distinguishable from the host dark matter halo. In semi-analytic models of galaxy formation (e.g. \\citealt{gao2004a, wang2006}), these subhalos are tracked as ``galaxies'' by either following the most bound particle of the subhalo or extrapolating their positions and velocities using information from previous timesteps.  The results from our three-component (dark matter+stars+black hole) non-cosmological simulations suggest that the properties of the observable stellar component in the simulations are very sensitive to their assumed initial properties, and in particular, to the initial stellar density assigned to the galaxies and the inclusion of black holes in the host galaxies.  Future work that incorporates the results presented here into a full cosmological context should be valuable."
    },
    "0608/astro-ph0608314_arXiv.txt": {
        "abstract": "Two sequences of solar images obtained by the {\\emph{Transition Region and Coronal Explorer}} in three UV passbands are studied using wavelet and Fourier analysis and compared to the photospheric magnetic flux measured by the Michelson Doppler Interferometer on the {\\emph{Solar Heliospheric Observatory}} to study wave behaviour in differing magnetic environments. Wavelet periods show deviations from the theoretical cutoff value and are interpreted in terms of inclined fields. The variation of wave speeds indicates that a transition from dominant fast--magnetoacoustic waves to slow modes is observed when moving from network into plage and umbrae. This implies preferential transmission of slow modes into the upper atmosphere, where they may lead to heating or be detected in coronal loops and plumes. ",
        "introduction": "\\label{sec:intro} It has been believed for over forty years that oscillatory motions play an important role in the outer solar atmosphere \\citep{lei62}. This is particularly true of the wide variety of perturbations that are possible within the chromosphere, as this is the atmospheric layer which spans the transition from the domination of gas pressure to that of the magnetic field. These oscillations may be the signature of wave energy carried outward from the photosphere and deposited in the chromosphere, hence resulting in heating of the outer atmosphere. As such, spatial variations in oscillatory behaviour could yield information on the local topology of the magnetic environment \\citep[\\eg,][]{fin04}. \\citet{fos05a} showed that hydrodynamic, high--frequency waves in the range $5-50$~mHz do not contribute sufficient energy to heat the atmosphere in a region of weak magnetic field. Also, \\citet{soc05} found that electric currents in the chromosphere above a sunspot do not show a substantial spatial correlation with regions of either high temperature or large temperature gradient. Hence, it appears entirely plausible that differing forms of magnetoacoustic waves could supply some portion of the excess energy input through their existence in both extremes of the chromospheric environment.  Previous work on chromospheric oscillatory behaviour has reported on differences between spatial regions such as ``cell boundary'' and ``cell interior'' \\citep[\\eg,][]{dam84, deu90}, or the (equally as ambiguously defined) magnetic ``network'' and non--magnetic ``internetwork'' \\citep[\\eg,][]{lit93, mca04}. Although the findings of these authors are generally discussed in relation to the magnetic environment, spatial distinctions were often not based upon any magnetic field measurements, but instead on the magnitude of time--averaged emission within some spectral line or passband. This arbitrary form of classification may give the false impression that some sharp transition occurs when moving from the weak--field regime to that of higher field strengths. In reality, the observed distribution of solar magnetism shows a continuous variation of field strengths \\citep{dom06}, suggesting that forms of oscillatory behaviour which are (in any manner) magnetically related should vary continuously as well \\citep{law03}.  The comparison of chromospheric oscillations to surfaces of constant plasma--$\\beta$, as determined by potential--field extrapolations from photospheric longitudinal magnetic fields \\citep{mci03, mci04, fin04} has shown the spatial distribution of oscillatory power and propagation characteristics. We note that $\\beta$, the ratio of gas to magnetic pressure, is given by $(2 / \\gamma ) ( c_S / v_A )^2$, where $c_S$ is the sound speed (=$\\gamma P / \\rho$), $\\gamma$ is the ratio of specific heats (=$5/3$), $P$ is the gas pressure, $\\rho$ is the mass density, $v_A$ is the Alfv\\'{e}n speed (=$B / \\sqrt{4 \\pi \\rho}$), and $B$ is the magnetic field strength. However, the above studies mainly report upon the co--spatiality of oscillation behaviour with a chosen plasma--$\\beta$ level (usually the unity surface; $c_S \\approx v_A$) without yielding qualitative diagnostic feedback on the local atmospheric environment. In addition, the complicated two--dimensional characteristics of chromospheric wave propagation, both hydrodynamic \\citep{ros02} and magnetoacoustic \\citep{bog03} in nature, are only now beginning to be successfully explored through careful modeling. The results of such work are decidedly sobering, indicating a veritable plethora of interacting wave modes propagating at different speeds, sometimes in differing directions. Another consideration apparent from the simulations is the importance of knowing the plasma environment (\\ie, gas or magnetic domination) in which the waves exist.   \\begin{deluxetable*}{lccccccccc} \\tablecolumns{10} \\tablewidth{0pc} \\tablecaption{\\trace\\ UV and \\mdi\\ Data Set Summaries\\label{tab:datasets}} \\tablehead{ \\colhead{Data Set}  &  \\colhead{Date}  &  \\multicolumn{4}{c}{\\trace}\t\t\t\t\t\t\t\t\t\t\t      &  \\multicolumn{4}{c}{\\mdi}\\\\ &  \t\t       &  \\colhead{Start Time}\t &  \\colhead{Frames}  &  \\colhead{Cadence}  &  \\colhead{Initial Pointing}\t\t      &  \\colhead{Start Time}\t&  \\colhead{Frames}  &  \\colhead{Cadence}  &  \\colhead{Initial Pointing}\\\\ &  \t\t       &  \\colhead{({\\sc{UT}})}  &  \t\t      &  \\colhead{(s)}\t    &  \t\t\t\t\t\t      &  \\colhead{({\\sc{UT}})}  &\t\t     &  \\colhead{(s)}\t   &  } \\startdata 1\t\t    &  2004 Jul 03     &  15:50\t\t\t &  128\t\t      &  16.0\t\t    &  ($6^{\\prime\\prime}$, $-16^{\\prime\\prime}$)     &  15:37  \t\t&  1\t  \t     &  \\nodata\t\t   &  ($-10^{\\prime\\prime}$, $-122^{\\prime\\prime}$)\\\\ 2\t\t    &  2000 Sep 22     &  08:00\t\t\t &  312\t\t      &  22.7\t\t    &  ($-174^{\\prime\\prime}$, $124^{\\prime\\prime}$)  &  08:01  \t\t&  118  \t     &  60.0\t\t   &  ($-88^{\\prime\\prime}$, $44^{\\prime\\prime}$)% \\enddata \\end{deluxetable*}   Our paper seeks to address the variable distribution of chromospheric UV intensity oscillations by studying their occurrence as a function of both detection height and the underlying photospheric magnetic flux. Observational data and their sources are discussed in \\S~\\ref{sec:data} with descriptions of the reduction and alignment processes, while \\S~\\ref{sec:anamet} outlines the wavelet and Fourier analysis. The time--localized nature of the wavelet transform is used to record the available distributions of wavepacket periods and durations. These parameters are of interest as oscillation periods can indicate which form of wave may exist (\\ie, acoustic or magnetic) while the duration in terms of oscillatory cycles indicates the degree of energy carried if an oscillation is indeed the signature of a propagating wave. A Fourier phase difference technique is then used to directly study wave propagation characteristics. The results are presented and discussed in \\S~\\ref{sec:res_dis} in terms of wave behaviours and the plasma environment, while in \\S~\\ref{sec:conc} the implications of our work are summarized. ",
        "conclusions": "\\label{sec:conc} Chromospheric intensity variations have been studied using both time--localized wavelet and Fourier analysis. Investigation of the oscillation behaviours has identified three regimes with distinctly differing properties.  The increase of wavelet period with \\mdi\\ field over $\\vert 25-100 \\vert$~G marks the modification of the magnetoacoustic cutoff period by field inclination effects. This arises from the ``magnetic canopy'' being pulled progressively closer to the sampled HOFs so that they experience more inclined fields. Wavelet durations show a decrease in the 1550~\\AA\\ data over those in the 1600~\\AA\\ and 1700~\\AA\\ data, indicating a loss of wave energy. This can be explained by waves undergoing mode conversion or reflection at the $\\beta$--unity surface \\citep{bog03}. Fourier analysis indicates that wave phase speeds increase with magnetic field between the 1700~\\AA\\ and 1600~\\AA\\ formation heights. It is proposed here that this is due to an increase in magnetic filling factor within the pixel sampling, yielding a greater contribution to the observed wave speed from the low--$\\beta$ ``flux--associated'' fast--mode speed over the slower sound speed of the surrounding high--$\\beta$ plasma.  The similarity of wavelet periods and durations over field strengths in the range $\\vert 100-450 \\vert$~G implies that waves in this regime do not experience a drastic change in their environmental topology as evidenced by the weaker--field case. The $\\beta$--unity surface is believed to lie below the formation height of 1700~\\AA\\ emission, no transitional interface exists between the passbands, and wave energy (\\ie, wavelet duration) is conserved with altitude. In contrast to the weak--field regime wave speeds decrease, which may be a signature of a change from dominant fast--mode magnetoacoustic waves to slow--mode waves when moving from network into plage.  Wavelet period distributions broaden and display a doubly--distributed nature over $\\vert 550-950 \\vert$~G, marking the sampling of sunspot regions. Interpretation of the 5--min component as being of penumbral origin is given strong support by the virtual disappearance of these higher--period components when sampling field strengths $\\ge$$\\vert 950 \\vert$~G (\\ie, umbrae). The decrease of phase difference gradient values over the mixed penumbral/umbral field strength range ($\\vert 550-950 \\vert$~G) suggests that either penumbral waves travel at greater speeds than those in plage and umbral regions or the formation heights of the 1700~\\AA\\ and 1600~\\AA\\ passbands are closer in altitude (\\eg, because of inclined fields) in the penumbra.  In the future, we will be able to better determine the intricate relationship between magnetic topology and propagating waves using high--resolution, vector magnetic field information. Such data will soon be obtained by the full--Stokes spectropolarimeter in the instrument suite of the Solar Optical Telescope ({\\sc{sot}}) Focal Plane Package ({\\sc{fpp}}) on board {\\sl{Solar--B}}. Coordinated observations between this instrument, \\trace, and the multi--wavelength tunable filters also contained in the {\\sc{sot}}/{\\sc{fpp}} will yield diffraction--limited, high--cadence imaging at multiple heights through the outer Solar atmosphere unparalleled by previous observations. The combination of these forms of data with a range of atmospheric models taking into account the differences in atmospheric stratification for differing magnetic fields will yield a powerful diagnostic suite for studies of solar wave phenomena."
    }
}